source,target
 star., star.
 Lower initial masses are preferred by the initial mass function thus initially less massive. and therefore older. stars are the preferred. when we estimate the age of WRAL.," Lower initial masses are preferred by the initial mass function thus initially less massive, and therefore older, stars are the preferred when we estimate the age of WR11."
 We have searched through our binary models to find those with a reasonable match for the 57. Velorum system., We have searched through our binary models to find those with a reasonable match for the $\gamma^2$ Velorum system.
 For a certain binary model to match. we require that the following are all true at some point during the lifetime of the binary.," For a certain binary model to match, we require that the following are all true at some point during the lifetime of the binary."
 As already cliscussecl the binary is eccentric and all our models have circular orbits., As already discussed the binary is eccentric and all our models have circular orbits.
 We therefore assume that. for our binary svstems to fit. it must have a radius somewhere in the range of separations of the observed. binary.," We therefore assume that, for our binary systems to fit, it must have a radius somewhere in the range of separations of the observed binary."
 Ehe elfect of eccentricity would be for mass-transfer to begin. earlier (Churchetal.2009). and thus the initial separation would be greater for more eccentric systems to obtain the same evolution as we discuss here., The effect of eccentricity would be for mass-transfer to begin earlier \citep{church} and thus the initial separation would be greater for more eccentric systems to obtain the same evolution as we discuss here.
 We choose to only consider svstems with mass-transfer after the main sequence (Case D)., We choose to only consider systems with mass-transfer after the main sequence (Case B).
 This is because when niss-transfer events occur on the main sequence (Case A) they happen on a nuclear timescale ancl the tidal forces would have more time to circularise the orbit., This is because when mass-transfer events occur on the main sequence (Case A) they happen on a nuclear timescale and the tidal forces would have more time to circularise the orbit.
 Post main-sequence mass-transfer occurs on a thermal timescale anc therefore tidal forces have less time to circularise the orbit and the resulting binary remains eccentric., Post main-sequence mass-transfer occurs on a thermal timescale and therefore tidal forces have less time to circularise the orbit and the resulting binary remains eccentric.
 We do not attempt to model the secondary here bevonc matching its mass., We do not attempt to model the secondary here beyond matching its mass.
 The secondary does lose mass by stellar winds and mass is transferred from the primary to. the secondary., The secondary does lose mass by stellar winds and mass is transferred from the primary to the secondary.
 The accretion rate onto the secondary is limitec o AlesΕπο Where Thema is the secondaries therma imescale.," The accretion rate onto the secondary is limited to $M_2/ \tau_{\rm thermal}$ where $\tau_{\rm
 thermal}$ is the secondaries thermal timescale."
 Lf the mass-transfer rate is greater than this value he excess mass is lost. [rom the svstem., If the mass-transfer rate is greater than this value the excess mass is lost from the system.
 We do not include angular momentum transfer and thermohaline mixing in his model., We do not include angular momentum transfer and thermohaline mixing in this model.
 This is because they produce complex allects on he secondary. especially angular momentum transfer which will alter the amount of rotationally induced. mixing.," This is because they produce complex affects on the secondary, especially angular momentum transfer which will alter the amount of rotationally induced mixing."
 This makes the evolutionary outcome of the secondary uncertain (Cantielloetal.2007:Stanclille&ιάνιάσο2009).," This makes the evolutionary outcome of the secondary uncertain \citep{cantiello,stancliffe}."
. We consider the secondary in more detail in Section 2.3., We consider the secondary in more detail in Section 2.3.
 We list our models that match 5? Velorum in Table I.., We list our models that match $\gamma^{2}$ Velorum in Table \ref{initialp}.
 The range of initial primary masses and initial separations are narrow., The range of initial primary masses and initial separations are narrow.
 “Phe initial separations are similar to the observed: separation., The initial separations are similar to the observed separation.
 This indicates that the wicening of he system due to stellar-wind mass-loss is countered by the nmiass-transfer event tightening the binary., This indicates that the widening of the system due to stellar-wind mass-loss is countered by the mass-transfer event tightening the binary.
 For each of the svstems a binary interaction mus jwe occurred., For each of the systems a binary interaction must have occurred.
 Therefore we might simply assume that 57 Velorum should be a circular orbit rather than eccentric., Therefore we might simply assume that $\gamma^2$ Velorum should be a circular orbit rather than eccentric.
 1 is worth asking. how unusual it is for à \Woll-Rayvet binary o be eccentric.," It is worth asking, how unusual it is for a Wolf-Rayet binary to be eccentric."
 Figure 3 in vanderισα(2001). shows hat in general binaries with periods of 30 days or less tenc o be circular., Figure 3 in \citet{wrcat7} shows that in general binaries with periods of 30 days or less tend to be circular.
 However there are a number of svstems with »riods between 30 and LOO days that have eccentricities in he range of 0.3 to 0.6., However there are a number of systems with periods between 30 and 100 days that have eccentricities in the range of 0.3 to 0.6.
 Phis is because binarics with periods xelow 30 days experience Case A mass-transfer and therefore heir tidal forces have a long time to circularise the orbit., This is because binaries with periods below 30 days experience Case A mass-transfer and therefore their tidal forces have a long time to circularise the orbit.
 The binaries with longer periods experience Case D. mass- that proceeds at shorter thermal timescales giving," The binaries with longer periods experience Case B mass-transfer, that proceeds at shorter thermal timescales giving"
reveals a line profile very similar to that first discovered by T95.,reveals a line profile very similar to that first discovered by T95.
 We note. however. that since a simple power-law is used to derive the line profile seen in this figure. it is possible that some fraction of the broad excess at ~5 keV can have contributions from the reflection spectrum known to exist in this source (e.g. Lee et al.," We note, however, that since a simple power-law is used to derive the line profile seen in this figure, it is possible that some fraction of the broad excess at $\sim 5$ keV can have contributions from the reflection spectrum known to exist in this source (e.g. Lee et al."
 1998. 1999).," 1998, 1999)."
 To further investigate the degree to which the relativistic disk line can account for the line profile of Fig. 3..," To further investigate the degree to which the relativistic disk line can account for the line profile of Fig. \ref{fig-hegasca},"
 we fit the data with a power-law plus diskline model modified by absorption., we fit the data with a power-law plus diskline model modified by absorption.
" The power-law component is that described previously and the diskline parameters are fixed at the T95 values: accretion disk at inclination /230°. respectively inner (Ri= 3.4R4) and outer (Rou= LORs) radii. line energy = 6.35 keV (6.4 keV in the galaxy frame). and radial emissivity a=3 assuming a power-law-type emissivity function xA"" of the line."," The power-law component is that described previously and the diskline parameters are fixed at the T95 values: accretion disk at inclination $\,i = 30^\circ$ , respectively inner $R_{\rm in} = 3.4
R_{\rm S}$ ) and outer $R_{\rm out} = 10 R_{\rm S}$ ) radii, line energy = 6.35 keV (6.4 keV in the galaxy frame), and radial emissivity $\alpha = 3$ assuming a power-law-type emissivity function $\propto
R^{-\alpha}$ of the line."
" This model gives 4/d.o.f= 42/47. and Fe Ka flux (1.3720.29)«107phem~s7!. with equivalent width Wy,295+118 eV. This is in good agreement with the broad iron line measured using the data which has Wx,,=398+58 eV. comparable to previous (e.g. Lee et al."," This model gives $\chi^2 \, / d.o.f =
42/47$ , and Fe $\alpha$ flux $(1.37 \pm 0.29) \times 10^{-4} \rm \,
ph \, cm^{-2} \, s^{-1}$, with equivalent width $W_{\rm K\alpha} \sim
295 \pm 118$ eV. This is in good agreement with the broad iron line measured using the data which has $W_{\rm K\alpha} = 398 \pm
58$ eV, comparable to previous (e.g. Lee et al."
 1998. 1999) and measurements of this source (e.g. Iwasawa et al.," 1998, 1999) and measurements of this source (e.g. Iwasawa et al."
 1996. 1999).," 1996, 1999)."
" The apparent ""sharp drop’ at 6.7 keV (Fig. 3))", The apparent `sharp drop' at 6.7 keV (Fig. \ref{fig-hegasca}) )
 appears resolved (AE70 eV from peak to drop)., appears resolved $\Delta E > 70$ eV from peak to drop).
" A simple Galactic absorbed power-law (as discussed previously) plus Gaussian fit to the data gives Ey,as,=6.21£0.08 keV. Gaussian width &=0.660.11 (FWMH ~73.000 ». and iron line flux /k,,=(2.13+£0.31)«107phem«'."," A simple Galactic absorbed power-law (as discussed previously) plus Gaussian fit to the data gives $E_{\rm K\alpha(obs)} = 6.21 \pm
0.08$ keV, Gaussian width $\sigma = 0.66 \pm 0.11$ (FWMH $\sim
73,000$ ), and iron line flux $I_{\rm K\alpha} = (2.13 \pm 0.31)
\times 10^{-4} \, \rm \ph \, cm^{-2} \, s^{-1}$."
 We next address the formal detectability of the tron line narrow component and resolution of its width., We next address the formal detectability of the iron line narrow component and resolution of its width.
 We note that the count rate of the full 125 ks observation at the iron line energies is ~20—30ctsbin! (where bin =0.005A.. or an ACIS pixel) and therefore sufficient for statistically meaningful errors derived from spectral fitting.," We note that the count rate of the full 125 ks observation at the iron line energies is $\sim 20-30 \rm \
cts\, bin^{-1}$ (where bin =, or an ACIS pixel) and therefore sufficient for statistically meaningful errors derived from spectral fitting."
 Table 1 shows the narrow component parameters when fitting a Gaussian to the residuals of the best fit power law described in refsec-cont (see also Fig. 5))., Table 1 shows the narrow component parameters when fitting a Gaussian to the residuals of the best fit power law described in \\ref{sec-cont} (see also Fig. \ref{fig-hilo}) ).
 The 68. 90. and confidence contours (3 free parameters and 515 degrees of freedom) of Fig.," The 68, 90, and confidence contours (3 free parameters and 515 degrees of freedom) of Fig."
 4. show that the narrow component 1s resolved at confidence with a full width at half maximum (FWHM) ~11.000kms! based on the full observation. and ~3600kms! from the *high’ continuum flux period (see below). (," \ref{fig-confcontours} show that the narrow component is resolved at confidence with a full width at half maximum (FWHM) $\sim 11,000 \, \rm km\,s^{-1}$ based on the full observation, and $\sim 3600 \, \rm km\,s^{-1}$ from the `high' continuum flux period (see below). ("
The statistical significance of the excess counts in the narrow component over a power law plus to account for the broad component is >99.9%..),The statistical significance of the excess counts in the narrow component over a power law plus to account for the broad component is $> 99.9$ .)
 The line strengths shown in Table | are consistent with those found for the narrow core by Iwasawa (1996. 1999).," The line strengths shown in Table 1 are consistent with those found for the narrow core by Iwasawa (1996, 1999)."
 Wilms et al. (, Wilms et al. (
"2002) report an ""unresolved"" (at the EPIC resolution which is Z4« the HEG at the iron energies) narrow iron line with 2 38 eV. Fabian et al. (",2002) report an “unresolved” (at the EPIC resolution which is $\approxlt 4\times$ the HEG at the iron energies) narrow iron line with = 38 eV. Fabian et al. (
2002) attribute this to the blue wing of the disk line.,2002) attribute this to the blue wing of the disk line.
 If there is à constant narrow core due to fluorescence from material far from the black hole. such as a molecular torus. then 1t should not vary during our observation and be strongest during the period. of low continuum flux.," If there is a constant narrow core due to fluorescence from material far from the black hole, such as a molecular torus, then it should not vary during our observation and be strongest during the period of low continuum flux."
 To test for the presence of an intrinsically narrow core. we assess the variability nature of the narrow component. by separating the data into a “high’ (67 ks: 0.4-10 keV unabsorbed flux frien=5\107ergem™s! ) and ‘low’ (58 ks: flow=3s4107!ereem7 s!) flux states arbitrarily defined as above and below the mean count rate (Fig. 1))," To test for the presence of an intrinsically narrow core, we assess the variability nature of the narrow component, by separating the data into a `high' (67 ks; 0.4-10 keV unabsorbed flux $f_{\rm high} = 5 \times 10^{-11} \, \rm erg\,cm^{-2}\,s^{-1}$ ) and `low' (58 ks; $f_{\rm low} = 3 \times 10^{-11} \, \rm erg\,cm^{-2}\,s^{-1}$ ) flux states arbitrarily defined as above and below the mean count rate (Fig. \ref{fig-lc}) )"
 of the time averaged 125 ks CoacoivomdsIOergem?s y observation.," of the time averaged 125 ks $f_{\, \rm (0.4-10~keV)} = 
4 \times 10^{-11} \, \rm erg\,cm^{-2}\,s^{-1}$ ) observation."
 To ensure that there are sufficient counts. these data were binned to 0.01 to create the contour plots corresponding to the ‘low’ and ‘high’ state shown in Fig. 4..," To ensure that there are sufficient counts, these data were binned to 0.01 to create the contour plots corresponding to the `low' and `high' state shown in Fig. \ref{fig-confcontours}."
 It can be seen that the region of overlap between these states exists only between the confidence contours., It can be seen that the region of overlap between these states exists only between the confidence contours.
 This would indicate that the joint probability is ~1% that the iron line narrow component is a narrow core from distant material., This would indicate that the joint probability is $\sim 1$ that the iron line narrow component is a narrow core from distant material.
 This is depicted in Fig., This is depicted in Fig.
 5 which shows that a comparison of the high versus low flux data against the best fit power-law model reveals that the iron line narrow component is effectively absent from the low flux state: σ could not be well constrained during this statebecause a very broad line (FMHM =20.000 kms') is required by the fit.," \ref{fig-hilo} which shows that a comparison of the high versus low flux data against the best fit power-law model reveals that the iron line narrow component is effectively absent from the low flux state; $\sigma$ could not be well constrained during this statebecause a very broad line (FMHM $\approxgt 20,000 \, \rm km \, s^{-1}$ ) is required by the fit."
 We note the features blue-ward of 6.5 keV (e.g. at ~6.82 keV) seen during the high and lowstates are only marginally ος20) significant.," We note the features blue-ward of 6.5 keV (e.g. at $\sim 6.82$ keV) seen during the high and lowstates are only marginally $\approxlt 2
\sigma$ ) significant."
Active galactic nuclei CAGNS) are luminous in wide wavelength because of the vast amount of energy produced by accretion onto à central supermassive blackhole.,Active galactic nuclei (AGNs) are luminous in wide wavelength because of the vast amount of energy produced by accretion onto a central supermassive blackhole.
 Accordingly. they are observable at nearly entire wavelength despite locating at great distances. and have been searched in various wavelengths: e.g.. optical etal.2007:Véron-Cetty&Véron 2006).. infrared (Lowetal.1988. 1989).. or radio (Frayeretal.2004...," Accordingly, they are observable at nearly entire wavelength despite locating at great distances, and have been searched in various wavelengths: e.g., optical \citep{Schneider2007-AJ,Veron2006-AA}, , infrared \citep{Low1988-ApJ,Low1989-ApJ}, or radio \citep{Frayer2004-AJ}."
 X-ray emission is especially a characteristic property of AGNs., X-ray emission is especially a characteristic property of AGNs.
 The vast majority of X-ray sources are AGNS and all classes of AGNS appear in X-ray surveys., The vast majority of X-ray sources are AGNs and all classes of AGNs appear in X-ray surveys.
 There are many studies that collect AGN samples using X-ray data (e...Kim&Elvis1999:Watanabeetal.2004:Polletta 2007)...," There are many studies that collect AGN samples using X-ray data \citep[e.g., ][]{Kim1999-ApJ,Watanabe2004-ApJ,Polletta2007-ApJ}."
 Therefore. whether an object is a X-ray source can be a criterion to select AGNs.," Therefore, whether an object is a X-ray source can be a criterion to select AGNs."
 Colour selection is a powerful technique in extracting AGN candidates., Colour selection is a powerful technique in extracting AGN candidates.
 Aclassical method is known as the (V -excess (UVX:Sandage1965:Schmidt&Green1983:Boyleetal. 1990)... which extract bluer quasars.," Aclassical method is known as the $UV$ -excess \citep[UVX; ][]{Sandage1965-ApJ,Schmidt1983-ApJ,Boyle1990-MNRAS}, which extract bluer quasars."
 Richardsetal.(2002). selected quasars via their nonstellar colours using SDSS photometry., \citet{Richards2002-AJ} selected quasars via their nonstellar colours using SDSS photometry.
 Other selections are such as red quasar survey using optical and near-infrared combined colours (Glikmanetal.2007). or mid-infrared selected AGNSs using Spitzer data (Lacyetal.2004:Stern 2005). ," Other selections are such as red quasar survey using optical and near-infrared combined colours \citep{Glikman2007-ApJ}, or mid-infrared selected AGNs using Spitzer data \citep{Lacy2004-ApJS,Stern2005-ApJ}. ."
Colour selection using near-infrared photometry has also performed by some previous studies., Colour selection using near-infrared photometry has also performed by some previous studies.
 extracted obscured AGN candidates using the colour selection of (JAvy)2.0., extracted obscured AGN candidates using the colour selection of $(J-K_\textnormal{\tiny S})>2.0$.
 The A.X method using the excess in K-band. proposed by Warrenetal.(2000)... was used for extracting quasars (Jureketal.2008:MaddoxSmailNakosetal. 2009)..," The $KX$ method using the excess in K-band, proposed by \citet{Warren2000-MNRAS}, was used for extracting quasars \citep{Jurek2008-MNRAS,Maddox2008-MNRAS,Smail2008-MNRAS,Nakos2009-AA}."
 However. these extracted only peculiar AGNs (in the former case) or extracted AGNs using combined with optical photometry Cin the latter case).," However, these extracted only peculiar AGNs (in the former case) or extracted AGNs using combined with optical photometry (in the latter case)."
 Because the optical light suffers more extinctions than the infrared light and an AGN is surrounded by a dust torus. a selection using optical and near-infrared combined colours may miss AGNS (especially obscured AGNs). because of lack of optical detection.," Because the optical light suffers more extinctions than the infrared light and an AGN is surrounded by a dust torus, a selection using optical and near-infrared combined colours may miss AGNs (especially obscured AGNs), because of lack of optical detection."
 However. Kouzuma&Yamaoka(2010). proposed colour selection criteria to extract AGNs using only near-infrared colours.," However, \citet{Kouzuma2010-AA} proposed colour selection criteria to extract AGNs using only near-infrared colours."
 They demonstrated by both observed and simulated colours that AGNS are differentiated from several types of objects in a(// AS)ff) colour-colour diagram (CCD)., They demonstrated by both observed and simulated colours that AGNs are differentiated from several types of objects in a$(H-K_\textnormal{\tiny S})$ $(J-H)$ colour-colour diagram (CCD).
 This enables us to extract AGN candidates using only near-infrared photometry., This enables us to extract AGN candidates using only near-infrared photometry.
 In this paper. we first extract bright sources in both near- and X-ray by a cross-identification between the Two Micron All Sky Survey (2MASS) and ROSAT all-sky survey catalogues. and select AGN candidates on the basis of the infrared colour selection criteria proposed by Kouzuma&Ya-maoka (2010).," In this paper, we first extract bright sources in both near-infrared and X-ray by a cross-identification between the Two Micron All Sky Survey (2MASS) and ROSAT all-sky survey catalogues, and select AGN candidates on the basis of the near-infrared colour selection criteria proposed by \citet{Kouzuma2010-AA}."
 In addition. we investigate properties of candidates using not only near-infrared and X-ray data but also photometric data at other wavelengths derived by cross-identifications with some catalogues.," In addition, we investigate properties of candidates using not only near-infrared and X-ray data but also photometric data at other wavelengths derived by cross-identifications with some catalogues."
 In Section ??.. we introduce the 2MASS and ROSAT.In Section ??.. we describe the method to extract AGN," In Section \ref{DATA}, , we introduce the 2MASS and ROSAT.In Section \ref{EXTRACTION}, , we describe the method to extract AGN"
areal coverage (aud hence larger nuuber of chuups) in the simulations than iu the observatious.,areal coverage (and hence larger number of clumps) in the simulations than in the observations.
 ILaviug more clumps iu the sample iav increase the statistical weight of the high-iass end of the mass function. where the statistics withiu the observations are usually poor (because ligher mass clips are rarer).," Having more clumps in the sample may increase the statistical weight of the high-mass end of the mass function, where the statistics within the observations are usually poor (because higher mass clumps are rarer)."
 Unlike the simulated inages. which differ oulv in the distance to the simulated region. the observational data sets differ widely iu the telescope used to acquire the data. the wavelength at which the data were acquired. and the chuup-iudiugs algorithii used to extract the chuups.," 	Unlike the simulated images, which differ only in the distance to the simulated region, the observational data sets differ widely in the telescope used to acquire the data, the wavelength at which the data were acquired, and the clump-finding algorithm used to extract the clumps."
 As we mentioned in refsecidefine. this supports the argunuent that the properties of derived clump mass functions do not depend stronely on the choice of. cbunip-fiudiug algoritlin.," As we mentioned in \\ref{sec:define}, this supports the argument that the properties of derived clump mass functions do not depend strongly on the choice of clump-finding algorithm."
 That the break mass should scale with the distance to the region being observed is exactly what one would expect if the break mass were not a property of the eusenible of clumps themselves. but a function of the aneular resolution at which they are observed.," 	That the break mass should scale with the distance to the region being observed is exactly what one would expect if the break mass were not a property of the ensemble of clumps themselves, but a function of the angular resolution at which they are observed."
 If the break mass reflected. sav. the local Jeans mass or some," If the break mass reflected, say, the local Jeans mass or some"
Dust iuske a rich galaxy. clusters Intracluster Medium (ICM) would be subjected to harsh concditious.,Dust inside a rich galaxy cluster's Intracluster Medium (ICM) would be subjected to harsh conditions.
 Parameters of the bot. X-ray. component of the ICM in galaxy. clusters have been measurecl by 1tumerous studies.," Parameters of the hot, X-ray component of the ICM in galaxy clusters have been measured by numerous studies."
 Typical temperatures are generally observed in the range 10°. which is equivalent to thermal energies of AT~2—141 keV (Bahlcall1999).," Typical temperatures are generally observed in the range $ T_{gas} \sim 10^{7} - 10^{8} K $ , which is equivalent to thermal energies of $ kT \sim 2 - 14$ keV \citep{B99}."
.. Grains. which have tysical iuolecular bonding potentials on the order of a few tenths to a lew eV. are likely to dissoclate by sputtering due to collisions with thermal electrous.," Grains, which have typical molecular bonding potentials on the order of a few tenths to a few eV, are likely to dissociate by sputtering due to collisions with thermal electrons."
 The dust grain sputtering timescales are depeudent on the size of the graius. α aud the electrou deusity at a given location ii the ICM. ne(r): for graphite. silicate or iron graius (Draiue&Salpeter 1979).. Electron deusities ou the order of He09LOehen have been measured (Jones&Forman1992).. so typical sputtering timescales are on the order of 7.sp~109—109yr. The shortest sputtering timescales correspond to grains locate in the cletises regions of the ICM.," The dust grain sputtering timescales are dependent on the size of the grains, $ a$ and the electron density at a given location in the ICM, $ n_{e}(r) $: for graphite, silicate or iron grains \citep{DS79}, , Electron densities on the order of $ n_{e} \sim 10^{-3} \cdot h^{1/2} cm^{-3} $ have been measured \citep{JF92}, so typical sputtering timescales are on the order of $\tau_{sp} \sim 10^{6} - 10^{9} yr.$ The shortest sputtering timescales correspond to grains located in the densest regions of the ICM."
 To estimate whether there will be auy dust iu rich clusters. dust destruction timescales have to be compared to dust injection timescales.," To estimate whether there will be any dust in rich clusters, dust destruction timescales have to be compared to dust injection timescales."
 Dust could couceivably be introducec iuto the ICM. through several processes iucluding rau pressure strippiug of galaxies as they trave through the ICM. the accretion of primordial dust. galaxy or cluster mergers aix collisions. blowout from galaxies that experience multiple or intense starbursts. or cooling flows (Popescuetal. 2000)..," Dust could conceivably be introduced into the ICM through several processes including ram pressure stripping of galaxies as they travel through the ICM, the accretion of primordial dust, galaxy or cluster mergers and collisions, blowout from galaxies that experience multiple or intense starbursts, or cooling flows \citep{PTFV00}. ."
 Mauy of these processes have timescales on the order of To105—10?yrs. which is comparable to the longest sputtering timescales.," Many of these processes have timescales on the order of $ \tau \sim 10^{8} - 10^{9} yrs$, which is comparable to the longest sputtering timescales."
 Each of these modes o dus ---isertion allect different spatial scales aud locations within clusters., Each of these modes of dust insertion affect different spatial scales and locations within clusters.
 For example. cooling flows occur uear the central regious of a clusters gravitatioual poteutial while accretion processes sucl as mergers affect the outer portions ofthe ICM.," For example, cooling flows occur near the central regions of a cluster's gravitational potential while accretion processes such as mergers affect the outer portions of the ICM."
 Furthermore. bydrodyuamiueal processes may allow dust to reside iu chumps. which would be self-shielded (rom ionizing radiation aud electron collisions. resulting in a longer ellective sputtering timescale aud a smaller covering factor.," Furthermore, hydrodynamical processes may allow dust to reside in clumps, which would be self-shielded from ionizing radiation and electron collisions, resulting in a longer effective sputtering timescale and a smaller covering factor."
 Therefore. the amount aud clistribution of dust tn a rich galaxy. cluster may also be depenclent ou the history of the various depositiou processes in an individual cluster.," Therefore, the amount and distribution of dust in a rich galaxy cluster may also be dependent on the history of the various deposition processes in an individual cluster."
 The nature of the extinctiou aud recddeuing from dust iu clusters iuay. also depeud heavily on environmeut because the lifetimes of erains with some types of chemistries and structures may ve longer than the lifetimes of others uncer the same ambient conditious., The nature of the extinction and reddening from dust in clusters may also depend heavily on environment because the lifetimes of grains with some types of chemistries and structures may be longer than the lifetimes of others under the same ambient conditions.
 In addition. deposition JJOCCSSesS COILIC Lact as filters.," In addition, deposition processes could act as filters."
 Grain acceleration by radiation pressure from starlielt. for example. could. iuipose a bias of grain radius aud mass ou particles entering the ICM. from galaxies.," Grain acceleration by radiation pressure from starlight, for example, could impose a bias of grain radius and mass on particles entering the ICM, from galaxies."
 Iu orocess. this ¢ould result in differences between the initial aud resulting extinction curves of the allect grain populations.," In process, this could result in differences between the initial and resulting extinction curves of the affect grain populations."
 Certainly. the ratio of total-to-selective extinction. à=Ay /E(B—V)is SLIOWLL 0o vary iu the range 3Xοà:S6 within the Milky Way (Mathis1990) and more strongly in other galaxies. where 1.5<RyX7.2) (Falcoetal.1999).," Certainly, the ratio of total-to-selective extinction, $R_{V} = 
A_{V} / E(B - V)$ is known to vary in the range $3 \lesssim R_{V} \lesssim 6$ within the Milky Way \citep{M90} and more strongly in other galaxies, where $1.5 \lesssim R_{V} \lesssim 7.2$ \citep{f99}."
. These facts together with the uucertaiuties i erain moclels for dust iu the ICM require a search for dust im tle ICM to beflexible interms of detecting coilijuatious of varying degrees| of reddening aud obscuration., These facts together with the uncertainties in grain models for dust in the ICM require a search for dust in the ICM to beflexible interms of detecting combinations of varying degrees of reddening and obscuration.
"Electrons with high energies produce a substantial energy gain of photons, which scattered to the high-frequency tail.","Electrons with high energies produce a substantial energy gain of photons, which scattered to the high-frequency tail."
" For an optical depth t«1 and a sufficiently large Comptonization parameter y, a sequence of declining peaks in the high-frequency tail (x>> 10) of the scattered photon energy flux spectrum is present (e.g., see Loeb et al."," For an optical depth $\tau\ll1$ and a sufficiently large Comptonization parameter $y$, a sequence of declining peaks in the high-frequency tail $x\gg10$ ) of the scattered photon energy flux spectrum is present (e.g., see Loeb et al."
 1991)., 1991).
" These peaks correspond successively to singly, doubly (etc) scattered photons (see Figs."," These peaks correspond successively to singly, doubly (etc) scattered photons (see Figs."
 4a-4d in Loeb et al., 4a-4d in Loeb et al.
 1991)., 1991).
" Next, we study the high-frequency tail of the scattered CMB spectrum by high energy electrons at frequencies below the first peak frequency."," Next, we study the high-frequency tail of the scattered CMB spectrum by high energy electrons at frequencies below the first peak frequency."
 At high frequencies (i.e. x>> 10) the term xl(exp(x)—1) of Eq. (5)), At high frequencies (i.e. $x\gg10$ ) the term $x^3/(\exp(x)-1)$ of Eq. \ref{G}) )
" decreases strongly and, therefore, the Eq. (5))"," decreases strongly and, therefore, the Eq. \ref{G}) )"
" can be written as The treatment|Baa is πρconsiderably simplified if exp(xexp(—s))>>1 and in the case the generalized spectral function G(x,Τε) is given by The sub-exponential{ asfunction f(s)=-ᾱ-—xexp(-s) has a maximum at the frequency shift smax=In(x/3)."," can be written as The treatment is considerably simplified if $\exp(x \exp(-s))\gg1$ and in the case the generalized spectral function $G(x,
T_{\mathrm{e}})$ is given by The sub-exponential function $f(s)=-3 s-x \exp(-s)$ has a maximum at the frequency shift $s_{\mathrm{max}}=\ln(x/3)$."
 Since exp(xexp(—Smax))>>1 the approximate expression for the generalized spectral function is valid and we calculate the integral in Eq. (22))," Since $\exp(x
\exp(-s_{\mathrm{max}}))\gg1$ the approximate expression for the generalized spectral function is valid and we calculate the integral in Eq. \ref{Ghf}) )"
 by the Laplace’s method., by the Laplace's method.
 The approximate value of the sub-exponential function in a neighborhood of the point s=Smax is then the spectral function approximately Tmequals 2n As was shown OT.)by Kino et al. (, The approximate value of the sub-exponential function in a neighborhood of the point $s=s_{\mathrm{max}}$ is then and the spectral function approximately equals As was shown by Kino et al. (
"2007) and Antonuccio-Delogu Silk (2008), high gas temperatures (kyT.~1 MeV) are expected in AGN cocoons.","2007) and Antonuccio-Delogu Silk (2008), high gas temperatures $k_{\mathrm{b}} T_{\mathrm{e}}\sim1$ MeV) are expected in AGN cocoons."
" For example, we choose the temperature equaled to 500 keV to study the high-frequency tail of the generalized spectral function."," For example, we choose the temperature equaled to 500 keV to study the high-frequency tail of the generalized spectral function."
" For high temperatures the function P\(s,Τε) is wide and it is centered at high values of the frequency shift s."," For high temperatures the function $P_1 (s, T_{\mathrm{e}})$ is wide and it is centered at high values of the frequency shift $s$."
" The distribution of frequency shifts for single scattering P|(s,Τε) at the temperature of kyT.=500 keV is shown in Fig. [L2]."," The distribution of frequency shifts for single scattering $P_1 (s,
T_{\mathrm{e}})$ at the temperature of $k_{\mathrm{b}}
T_{\mathrm{e}}=500$ keV is shown in Fig. \ref{P1}."
" Figure shows that the values of the distribution of frequency [I2]shift lie in the narrow range (0.25, 0.28) if the frequency shift s is in the range (1.5, 3.0)."," Figure \ref{P1} shows that the values of the distribution of frequency shift lie in the narrow range (0.25, 0.28) if the frequency shift $s$ is in the range (1.5, 3.0)."
" The frequency shift Smax lies in this range when the dimensionless frequency is of 15«x90 and, therefore, the approximate value of the generalized spectral function in this frequency range is The quantitative dependence of the spectral function G(x,kpT.=500 keV) on the dimensionless frequency x is illustrated in Fig.[I3]."," The frequency shift $s_{\mathrm{max}}$ lies in this range when the dimensionless frequency is of $15<x<90$ and, therefore, the approximate value of the generalized spectral function in this frequency range is The quantitative dependence of the spectral function $G(x,
k_{\mathrm{b}} T_{\mathrm{e}}=500$ keV) on the dimensionless frequency $x$ is illustrated in Fig. \ref{500}."
 We conclude that the generalized spectral function should be flat in the broad frequency range when the temperature values are sufficiently high., We conclude that the generalized spectral function should be flat in the broad frequency range when the temperature values are sufficiently high.
" Since the non-relativistic spectral function g(x) is a rapidly decreasing function at high frequencies (x>10) in contrast with the spectral function G(x,Τε) at high electron temperatures, a measurement of the SZ effect at high frequencies provides an interesting test of the presence of high energy electrons."," Since the non-relativistic spectral function $g(x)$ is a rapidly decreasing function at high frequencies $(x>10)$ in contrast with the spectral function $G(x, T_{\mathrm{e}})$ at high electron temperatures, a measurement of the SZ effect at high frequencies provides an interesting test of the presence of high energy electrons."
 The direct detection of the SZ effect can in principle provide a unique diagnostic tool to study the physical conditions of the ICM at high redshift., The direct detection of the SZ effect can in principle provide a unique diagnostic tool to study the physical conditions of the ICM at high redshift.
" To reach this end, an accurate treatment of the spectral properties of the SZ signal is very important, and this was our main aim in this work."," To reach this end, an accurate treatment of the spectral properties of the SZ signal is very important, and this was our main aim in this work."
" Previous models were based on simplifying assumptions, like the assumption of pressure equilibrium and density homogeneity (Colafrancesco, 2005; Pfrommer et al.,"," Previous models were based on simplifying assumptions, like the assumption of pressure equilibrium and density homogeneity (Colafrancesco, 2005; Pfrommer et al.,"
 2005)., 2005).
" Here we have instead considered the SZ signal arising from a inhomogeneous exactly solvable configuration, i.e. a spherically symmetric Sedov-expanding region, and even more realistic models obtained from 2D fluid-dynamical simulations of jet/cocoon system propagating into the ISM/IGM. "," Here we have instead considered the SZ signal arising from a inhomogeneous exactly solvable configuration, i.e. a spherically symmetric Sedov-expanding region, and even more realistic models obtained from 2D fluid-dynamical simulations of jet/cocoon system propagating into the $/$ "
birthplace indicating a NS age ~5«LO? ves (Walter&Lattimer2002:Kaplanetal. 2002).,"birthplace indicating a NS age $\sim5\times10^5$ yrs \citep{walter02, kaplan02}."
. Deep aud observations with total exposure time ~ 500 ks were performed to search for spectral feature and pulsation., Deep and observations with total exposure time $\sim$ 500 ks were performed to search for spectral feature and pulsation.
 However. neither was found in tle X-ray data which had Ligh resolution aud seusitivitv (Ransomctal.2002:Drakeetal. 2002).," However, neither was found in the X-ray data which had high resolution and sensitivity \citep{ransom02, drake02}."
. This lack of observed pulsation is perplexing since modulation of an anisotropic temperature distribution suggested. bv the two-component blackbody model fit discussed below seems expected. aud pulsation modulation of N-rav emission has heen detected iu other isolated NSs with thermal emission (Παυσetal.1997:Zavlinct 2000).," This lack of observed pulsation is perplexing since modulation of an anisotropic temperature distribution suggested by the two-component blackbody model fit discussed below seems expected, and pulsation modulation of X-ray emission has been detected in other isolated NSs with thermal emission \citep{haberl97, zavlin00}."
. We propose that has been spun-down to à spin period longer than 10! sec by the propeller effect., We propose that has been spun-down to a spin period longer than $10^4$ sec by the propeller effect.
 The fraction of inaguetars among isolated NSs is 101 (Ikouveliotouetal.1998).. aud is mich larger among NSs froii which thermal cussion has been detected.," The fraction of magnetars among isolated NSs is $\sim 10^{-1}$ \citep{kouveliotou98}, and is much larger among NSs from which thermal emission has been detected."
 We &ud that the ouly plausible way to achieve the needed condition for carly transition iuto the propeller phase aud rapid enough spin-down thereafter is for this star to be a haguctar., We find that the only plausible way to achieve the needed condition for early transition into the propeller phase and rapid enough spin-down thereafter is for this star to be a magnetar.
 We first assune B>Lot! C. sufficiently lavee that ransition iuto the propeller phase occurs i much less han 5«10° vears.," We first assume $B\ge 10^{14}$ G, sufficiently large that transition into the propeller phase occurs in much less than $5\times10^5$ years."
" Then. we searched parameter space in the + à plane for the region in which could. hereafter. have been spun-down to a period longer than 104 see,"," Then, we searched parameter space in the $\gamma$ $\delta$ plane for the region in which could, thereafter, have been spun-down to a period longer than $10^4$ sec."
" luput paramicters are magnetic dipole moment ji. inconnue eas density ,,. and iuconung gas velocity ¢,,."," Input parameters are magnetic dipole moment $\mu$, incoming gas density $n_m$, and incoming gas velocity $v_m$."
 We fixed ο from the observed stellar proper motion at 200 kant., We fixed $v_m$ from the observed stellar proper motion at 200 $^{-1}$.
 ISM particles (αμαλα] hydrogen) at the magnetosphere radius will be charged aud be reflected by the stellar maguctosphere. [, ISM particles (mainly hydrogen) at the magnetosphere radius will be charged and be reflected by the stellar magnetosphere. [
Thermal photons fron: the NS. surface ionize all livdvogen iu the vicinity of the star.,Thermal photons from the NS surface ionize all hydrogen in the vicinity of the star.
" The ionization time of avdrogen at R,,~10+? cur (~107 s) is jiuch shorter than A,/0,, (~10° 8).", The ionization time of hydrogen at $R_m\sim10^{12}$ cm $\sim10^2$ s) is much shorter than $R_m/v_m$ $\sim 10^5$ s).
 Therefore. hvdrogen is fully-ionized at the magnetosphere boundary before the star reaches that region].," Therefore, hydrogen is fully-ionized at the magnetosphere boundary before the star reaches that region]."
 Figure 1/ shows the + 6 parameter space for which can spin-down to P10! sec within 5«107 vears., Figure \ref{fig_diagram} shows the $\gamma$ $\delta$ parameter space for which can spin-down to $P > 10^4$ sec within $5\times10^5$ years.
" Our rough interface model (4=61l.» 3/7) can achieve the required spiu-down for Bo~5«Late C (pay= 5) andy,=1 and B~5«1011 C (μοι= 0.5) aud nv,=LO? P (estimated? deusity of the nearby molecular cloud R CrA through which nav have passed (Cadamminietal. 1998)))."," Our rough interface model $\gamma=\delta=1, n=3/7$ ) can achieve the required spin-down for $B\sim5\times10^{15}$ G $\mu_{33}=5$ ) and $n_m=1$ $^{-3}$ and $B\sim5\times10^{14}$ G $\mu_{33}=0.5$ ) and $n_m=10^5$ $^{-3}$ (estimated density of the nearby molecular cloud R CrA through which may have passed \citep{giannini98}) )."
 The surface temperature of is roughly consistent with predictions of standard cooling curves at its age ~SS107 vears (Tsurutaetal.2002)., The surface temperature of is roughly consistent with predictions of standard cooling curves at its age $\sim5\times10^5$ years \citep{tsuruta02}.
". may radiate because of accretion at an icoming mass rate M=zhup51ü*Polhu, o. I"," may radiate because of accretion at an incoming mass rate $\dot{M}=\pi
R_m^2 v_0 \rho_m = 5\times10^7R_{m,12}^2 v_{m,7} n_m$ $^{-1}$."
Llowever. the ταν Iunimositv from such au accretion rate is below the detection limit of and observations (Rutledge2001)...," However, the X-ray luminosity from such an accretion rate is below the detection limit of and observations \citep{rutledge01_2}."
 A third heat source may be the coutiuuous dissipation of the large magnetic field energv in the stellar crust (ev)&I&ullrni1998)., A third heat source may be the continuous dissipation of the large magnetic field energy in the stellar crust \citep{heyl98}.
. Given featureless X-ray spectra. several approaches using spectral energy distribution (SED) in the optical and N-rav baud have been undertaken to reveal the surface composition (Ponsetal.2002).," Given featureless X-ray spectra, several approaches using spectral energy distribution (SED) in the optical and X-ray band have been undertaken to reveal the surface composition \citep{pons02}."
. Light clement atmospheres (ID and We) were ruled out since they predict about two orders of magnitudes larger optical fiux conrpared to the observed values (Pousctal.2002)., Light element atmospheres (H and He) were ruled out since they predict about two orders of magnitudes larger optical flux compared to the observed values \citep{pons02}.
. Ou the other haud. a blackbody model uuderpredicts optical flux bv a factor of ~7 (Walter&Lattimer2002).," On the other hand, a blackbody model underpredicts optical flux by a factor of $\sim 7$ \citep{walter02}."
.. Non-magnetized heavy clement atmosphere models (e.g. Si-ash or hon) predict correct SED over the optical aud Nav baud with a single temperature (ALS=10 eV) (Walter&Lattimer2002)., Non-magnetized heavy element atmosphere models (e.g. Si-ash or Iron) predict correct SED over the optical and X-ray band with a single temperature $kT^{\infty} = 40$ eV) \citep{walter02}.
. However. they have absorption features which deviate from the featureless N-ravw data (Burwitzetal.2003).," However, they have absorption features which deviate from the featureless X-ray data \citep{burwitz02}."
. A similar situation exists for magnetized heavy clement atmospheres (Rajagopaletal.1997)., A similar situation exists for magnetized heavy element atmospheres \citep{rajagopal97}.
. A two-component blackbody model was proposed to account for the multinvaveleneth SED aud the featureless N-rav spectra (Ponsetal.2002)., A two-component blackbody model was proposed to account for the multi-wavelength SED and the featureless X-ray spectra \citep{pons02}.
. Tot aud cold blackbody conrponeuts for N-rav and optical spectra respectively are cousisteut with both SED aud featureless X-ray data (Pousetal.2002:Draje&Romani 2002).," Hot and cold blackbody components for X-ray and optical spectra respectively are consistent with both SED and featureless X-ray data \citep{pons02, braje02}."
. IHTowever. it is lard to obtain blackbods-like spectra when siguificant atmosphere is present on the surface.," However, it is hard to obtain blackbody-like spectra when significant atmosphere is present on the surface."
 At sufficiently high magnetic field streneth and low temperature. a NS surface becomes very deuse liquid or solid with almost no atinosphere above it (Ruderman1971:Lai&Salpeter 1996).," At sufficiently high magnetic field strength and low temperature, a NS surface becomes very dense liquid or solid with almost no atmosphere above it \citep{ruderman71, lai96}."
. That is a maguetar is consistent with the two temperature model (anisotropy of surface temperature distribution caused by strong magnetic field} aud the model for a condeused irou surface because they both require strong magnetic field strength on the surface., That is a magnetar is consistent with the two temperature model (anisotropy of surface temperature distribution caused by strong magnetic field) and the model for a condensed iron surface because they both require strong magnetic field strength on the surface.
 The cluissivity of such a coudeused matter surface in 105 C is about 1/2 that of a blackbody because the stronely ⋅ ⋅Beaad-6/25 ⊳↘⋅ ⊈⋚∶↕∩⊔⊈⋚⊔≼∶⋝∐⋜↧↴∖↴↸∖↴∖↴↴∖↴↸∖∐↑↕⋜↧∐⋅↖↽∐∪↸∖∐∐↴∖∷∖↴↕↖↽↕↑⋅↖↽↕≯∪↥⋅↻≓↕⊔∪≼∐∖ ⊸∖⊽≓↥⋅⋜↧∙↖↽↴∖↴↴⋝∏↑↕↴∖↴∐↸∖⋜∐⋅↴⋝↕⋜↕↸⊳↨∖↽↴⋯≺↧∙↖↽↕⋡∪↥⋅⊏≓⋯∪≺∐∖∪↕∐∖↴∖↴∙↽∕∏∐∖∐↑∐∖ ⋜↧↖↽↸∖↥⋅⋜↧∶↴⊾⊾↸∖↸∖∐∏↴∖∷∖↴↕↖⇁↕↑⋅↖⇁↕," The emissivity of such a condensed matter surface in $10^{15}$ G is about 1/2 that of a blackbody because the strongly magnetized very dense $\rho\sim
560Z^{-3/5}B_{12}^{6/5}$ $^{-3}$, with $B=10^{12}B_{12}$ G) has essentially no emissivity for O-mode X-rays but is near blackbody for E-mode ones."
⋟∪↥⋅↑↕∐∖↥⋅⋜∥∐⋜↧↑↕∐∶↴⋁↴∖↴↿∐⋅↕⋟⋜↧↸⊳↸∖↕↴∖↴∿∶≩∩⋅↱↗∩⋰⋰⊒∶∕ ≼∐∖↻↸∖∐≼∐∐∶↴⋁∪∐↻∐∪↑∪∐↸∖∐↸∖↥⋅∶↴∙⊾⋅↖↽⋜⋯≼↧↕⊔⋜↧∶↴∙⊾∐↸∖↑↕↸⊳∐↸∖↕≼↧∶↴∙⊾↸∖∪∐∐∖⊓⋅⋅↖↽ (Zaueetal.2003)., Then the average emissivity for the radiating surface is $\sim$ depending on photon energy and magnetic field geometry \citep{zane03}.
. The actual NS area should then be inore than twice that of the apparent blackbody area so that the iferred NS radius becomes Z2Rpp (Zane 2005)..., The actual NS area should then be more than twice that of the apparent blackbody area so that the inferred NS radius becomes $\uax \sqrt{2} R_{BB}$ \citep{zane03}.
 Will we find other isolated NSs in the propeller phase?, Will we find other isolated NSs in the propeller phase?
 There are several NSs with discrepant supernova renuit and “canonical” spiu-dowun ages (Tes)., There are several NSs with discrepant supernova remnant and “canonical” spin-down ages $\tau_{csd}$ ).
 Iu some cases. the discrepancy may be due to the fact that NSs are in a propeller phase. (," In some cases, the discrepancy may be due to the fact that NSs are in a propeller phase. ("
Alternatively. they could have been born with a lone spin period close to preseut value.),"Alternatively, they could have been born with a long spin period close to present value.)"
" However. some of them have spin periods shorter than «1 sec. so it is difficult for them to euter a propeller phase unless the vuubicut gas deusity is very large (9,25l en. 5]."," However, some of them have spin periods shorter than $< 1$ sec, so it is difficult for them to enter a propeller phase unless the ambient gas density is very large $n_m \gg 1$ $^{-3}$ )."
 Detection of long pulsation periods (2LO sec) from isolated NSs mav be another indication of NSs iu a xopeller phase., Detection of long pulsation periods $> 10$ sec) from isolated NSs may be another indication of NSs in a propeller phase.
 Potential candidates for such isolated NSs would have ages between 10° years (enoush time ‘or spinning down to cuter the propeller phase} aud 10° vears (still detectable thermal cmussion)., Potential candidates for such isolated NSs would have ages between $10^3$ years (enough time for spinning down to enter the propeller phase) and $10^6$ years (still detectable thermal emission).
 Ta à maguetar. naenetic field decay processes cau keep the maguctar nore N-rav Dhuuimous than would be the case for canonical oss (Tevl&dulkuni 1998)...," In a magnetar, magnetic field decay processes can keep the magnetar more X-ray luminous than would be the case for canonical pulsars \citep{heyl98}. ."
 Then older maguctars wielt be still observable bv their N-rav cuission after, Then older magnetars might be still observable by their X-ray emission after
and are likely also in lower mass svstems.,and are likely also in lower mass systems.
 This will further complicate not only the absolute detection of a cluster. but could bias the estimation of flux in a single band.," This will further complicate not only the absolute detection of a cluster, but could bias the estimation of flux in a single band."
 In à future work (Scharf. in preparation) I will discuss the more complex issues involved with X-ray and $-Z detection biases lor clusters in the context of their use as cosmological probes.," In a future work (Scharf, in preparation) I will discuss the more complex issues involved with X-ray and S-Z detection biases for clusters in the context of their use as cosmological probes."
 CAS gratefully acknowledges helpful diseussions with D. IHelfand and F. Paerels and {he generous support of the Columbia Astrophysics Laboratory lor this work., CAS gratefully acknowledges helpful discussions with D. Helfand and F. Paerels and the generous support of the Columbia Astrophysics Laboratory for this work.
 ee In order toobtain (4.18) over cutolE A7.It reflectsthe factthat,"The partial wave reduction of $\calm^2+\nu^2$ is an ordinary differential operator of second order, and for computing the determinant of an ordinary differential operator there is the Gel'fand-Yaglom theorem stating that from which we then obtain Here the functions $\tilde f^-_n$ are identical to the functions $f^-_n$ of the previous subsection for the normalization."
lima divergent.in contrast to t," The normalization is fixed by writing, and one which establishes a direct contact with method I. The first version of the proof is based on the condition for a bound state $\lim_{r\to \infty}f^-_n(r,\nu^2)=0$."
"he sum$7,(J, (0. A?)). Sothe opera"," Furthermore a basic assumption is that $\bfJ_n(\nu^2)\to 1$ as $\nu^2\to \infty$ , i.e., that the determinant of $\bfM_n$ tends towards the one of $\bfM_{n,0}$ in this limit, within each partial wave subspace."
tions of , This is the case for potentials of finite range.
summation ancltaking, But then again we have to sum over $n$ and this sum will be logarithmically divergent.
 limit do notcommute., The renormalization for this case as forthe Green' s function method is discussed insubsection \ref{renorm}. .
or TESS.,or HESS.
 The svuchrotron extension of the flare emission in the hard N-rav/soft x-ray domain explains the x 101211 episode of a profound increase of fiux detected by INTEGRAL from the direction of (Bélangeretal.200 1).," The synchrotron extension of the flare emission in the hard X-ray/soft $\gamma$ -ray domain explains the $\simeq 40\,$ min episode of a profound increase of flux detected by INTEGRAL from the direction of \citep{bel04}."
. The abseuce of appareut TeV flux variations lav sugeest a proton origin for the TeV radiation (Aharonian&Neronov 2001)., The absence of apparent TeV flux variations may suggest a proton origin for the TeV radiation \citep{an04}.
. Indeed. proton and ion acceleration through first- and second-processes in the ADAF and through first-order acceleration at the wind termination shock could make cosimic ravs to produce TeV oenüssion tough unclear pputeractious with à~LO?cur? deux eas on pe scales; aud. could. form extended TeV eumissio- possibly already. detected with TESS (Aharonianetal. 2001).," Indeed, proton and ion acceleration through first- and second-processes in the ADAF and through first-order acceleration at the wind termination shock could make cosmic rays to produce TeV emission through nuclear $pp$ -interactions with $n\sim 10^{3}\,\rm cm^{-3}$ dense gas on $pc$ scales, and could form extended TeV emission possibly already detected with HESS \citep{HESS}."
 The hadronic origin of the TeV radiation in tlic ADAF itself in the BID vicinity requires. however. a- extremely dense gas target or extremely large proto- powers =10°?ores1," The hadronic origin of the TeV radiation in the ADAF itself in the BH vicinity requires, however, an extremely dense gas target or extremely large proton powers $\gtrsim 10^{39}\,\rm erg\, s^{-1}$."
 Finally. we note that the unideuti&ed ECGRET source BEC J1716-2851 towards the CC (Alaver-Tasselwauderetal.1998). is significantly displaced (Diugus&Looper2002) from the direction of the GCDIL. aud is uulikelv to be related to Ser A.," Finally, we note that the unidentified EGRET source 3EG J1746-2851 towards the GC \citep{EGRET1} is significantly displaced \citep{dh02}
 from the direction of the GCBH, and is unlikely to be related to Sgr $^\ast$."
 It is probably cinission from a voung pulsar. though not with the “mouse” PSR JL?l7-2058 (MeLauehliu&Cordes2003).," It is probably emission from a young pulsar, though not with the “mouse"" PSR J1747-2958 \citep{mc03}."
. A vouug pulsar with s-rav properties like Vela but with apparent οταν power zLOS larger could have been nüssed in pulsar surveys due to the large dispersion measure towards the GC., A young pulsar with $\gamma$ -ray properties like Vela but with apparent $\gamma$ -ray power $\approx 10 \times$ larger could have been missed in pulsar surveys due to the large dispersion measure towards the GC.
 We sugecst a deeper. higher radio-frequeucy aud X-ray search at the refined location of 3EG J1716-2851.," We suggest a deeper, higher radio-frequency and X-ray search at the refined location of 3EG J1746-2851."
 (Larsonctal.POLO) Yol1«X20 μι. hieh-- ESI!~107 (Fanctal.2006).. zz6.5.Γ Ίσα Lya va (LF) of Lvo enütters (LAEs) selected. via narrowbaud filters have revealed a possible «ecline in abuudauce between 2=5.7 and 2=7.0 (Nashikwvaetal.2006:Iveetal.2006:Otact2008:Ouchi 2010)). offering tautalizine evideice that this shor time interval (2200 Απ) ay corres»oud to one during which there is some evolution iu the reutral fraction.," \citep{Larson10} $7<z<20$ $x_{HI}$ $z$ $x_{HI}\simeq 10^{-3}$ \citep{Fan06}, $z\simeq 6.5$ $\alpha$ $\alpha$ $\alpha$ (LF) of $\alpha$ emitters (LAEs) selected via narrowband filters have revealed a possible decline in abundance between $z=5.7$ and $z=7.0$ \citealt{Kashikawa06,Iye06,Ota08,Ouchi10}) ), offering tantalizing evidence that this short time interval $\simeq$ 200 Myr) may correspond to one during which there is some evolution in the neutral fraction."
 But since a umber of astroplivsical facors can also affec the presence of Lyra endssion. if lay be dangerous to civectly link evolution in the Ίσα LF to reionizatiou (c.g.. Davaletal. 2010)).," But since a number of astrophysical factors can also affect the presence of $\alpha$ emission, it may be dangerous to directly link evolution in the $\alpha$ LF to reionization (e.g., \citealt{Dayal10}) )."
 These factors incude time-dependeit changes in the host ealaxy nuniber dcsusity. dust obseur‘ation aud interstellar eas content and kineniaic proportics.," These factors include time-dependent changes in the host galaxy number density, dust obscuration and interstellar gas content and kinematic properties."
 By culareing the LAE samples. it nav be possible to bypass sole of these complications by testing for the expected change iu their spatial clustering and lije profiles as the neutral era is entered (Ouchieal.20]0).," By enlarging the LAE samples, it may be possible to bypass some of these complications by testing for the expected change in their spatial clustering and line profiles as the neutral era is entered \citep{Ouchi10}."
. A complement:wy approach introduced ii Starsetal.(2010) (hereafer Paper D Is to. spectroscopically measure the fracion of strong Ly| cluitters within thepopulation., A complementary approach introduced in \citet{Stark10a} (hereafter Paper I) is to spectroscopically measure the fraction of strong $\alpha$ emitters within the.
 By tracing he redshift-depeudeut fraction. the lost ealaxv umber densitv is not a factor.," By tracing the redshift-dependent fraction, the host galaxy number density is not a factor."
 Evoluion lu dust obscuration can be independently racked sine the contiuuuu colors aud ISM kinematics through deep spectroscopy (Steideletal.2010:Bowens2009.20102:Vauzellactal. 2009).," Evolution in dust obscuration can be independently tracked using the continuum colors and ISM kinematics through deep spectroscopy \citep{Steidel10,Bouwens09a,Bouwens10a,Vanzella09}."
 Althougi demanding observationaIv. liehOo throughputC» spectroeraplsC» such as FORS2 on the ESO Very Large Telescope aud DEIMOS on the Keck II telescope have enable progress du recent vears (Paper I. Vauzellaetal.20097).," Although demanding observationally, high throughput spectrographs such as FORS2 on the ESO Very Large Telescope and DEIMOS on the Keck II telescope have enabled progress in recent years (Paper I, \citealt{Vanzella09}))."
 With theadditional information on the host galaxies possible for the LBC population. we cau hope to nore reliablv," With theadditional information on the host galaxies possible for the LBG population, we can hope to more reliably"
refer to the following SExtractor output parameters: We exclude. objects with FLAGS (extraction flags) >1 from our analysis.,refer to the following SExtractor output parameters: We exclude objects with FLAGS (extraction flags) $\ge 1$ from our analysis.
 We aake star/galaxv separation using the SSTAR paraicter (stcllavitv iudex) and keep the objects with CLASS.STARx0.2 as faint galaxy candidates., We make star/galaxy separation using the STAR parameter (stellarity index) and keep the objects with ${\rm CLASS\_STAR}\le 0.2$ as faint galaxy candidates.
 Objects with FWIALTAIAGCE (FWIIM. profile from a Gaussian fit to the core) <1 pixels are excluded since image shapes of extremely simall objects relative to the pixel scale may be affected by the anisotropic PSF., Objects with IMAGE (FWHM profile from a Gaussian fit to the core) $<4$ pixels are excluded since image shapes of extremely small objects relative to the pixel scale may be affected by the anisotropic PSF.
 Finally. all objects with faisC(25.5.27.5) are selected as background galaxy caucidates.," Finally, all objects with $I_{814{\rm W}}\in(25.5,27.5)$ are selected as background galaxy candidates."
" This detection and sclection procedure leads to the fal catalogs with total galaxy numbers of Ny= 200 and 211. correspouding nuniboer densities of p,= 15.8 and 56.5 arcu?. for the 1991 aud 1995 data. respectively."," This detection and selection procedure leads to the final catalogs with total galaxy numbers of $N_{\rm g}=$ 200 and 241, corresponding number densities of $n_{\rm g}=$ $45.8$ and $56.8$ $^{-2}$, for the 1994 and 1995 data, respectively."
 Our first aim is to obtain the distribution of woeak-lensing S/N in the data field., Our first aim is to obtain the distribution of weak-lensing S/N in the data field.
 Then. the ligh peaks in the S/N imaps cau be ideutified as clusters or mass overdeusities.," Then, the high peaks in the S/N maps can be identified as clusters or mass overdensities."
 To do this. we make use of a variant of aperture mass statistics.," To do this, we make use of a variant of aperture mass statistics."
" The statistics rely outhe fact that the shear 5=(541.52) and the convergence &:;—N/M are related to cach other through Ve=Dἄν, where Mods the surface mass density of the deflector. M4=(οπόλΩ.(DgDg. is the evitical surface mass ceusity. aud D={Diy} Uj= 1.2) is a differential operator defined by (kaiser 1995)."," The statistics rely on the fact that the shear $\gamma=(\gamma_1,\gamma_2)$ and the convergence $\kappa:=\Sigma/\Sigma_{\rm cr}$ are related to each other through $\vec{\nabla} \kappa
=\hat{D}\gamma
\equiv \vec{u}_{\gamma}$, where $\Sigma$ is the surface mass density of the deflector, $\Sigma_{\rm cr}=(c^2/4\pi G)D_{\rm s}/D_{\rm d}D_{\rm ds}$ is the critical surface mass density, and $\hat{D}=\{\hat{D}_{ij}\}$ $i,j=1,2$ ) is a differential operator defined by (Kaiser 1995)."
 Since (6. is the eracient of &. operating D further ou 4. vields Iu weak lensing liuüt (&1l aud [|| 1). the expectation value of the image ellipticity. Ele(0)). is the shear (0).," Since $\vec{u}_{\gamma}$ is the gradient of $\kappa$, operating $\hat{D}$ further on $\vec{u}_{\gamma}$ yields In weak lensing limit $\kappa\ll1$ and $|\gamma|\ll 1$ ), the expectation value of the image ellipticity, ${\rm E}[\epsilon(\vec{\theta})]$, is the shear $\gamma(\vec{\theta})$."
 Iu practice. however. background sources have intrinsic cllipticities ej. so that weak lensing analysis involves the smoothing procedure to reduce the noise.," In practice, however, background sources have intrinsic ellipticities $\epsilon_{({\rm s})}$, so that weak lensing analysis involves the smoothing procedure to reduce the noise."
 We denote the simoothed fields by augular brackets Co: eg. Ge)=[f40WIοUa)(y. where TV(0:d) is a inooth. continuous window function with a characteristic scale of 0.," We denote the smoothed fields by angular brackets $\left< \ \right>$: e.g., $\left<\kappa\right>=\int\,d^2\theta'\,
W(|\vec{\theta}-\vec{\theta}'|;\vartheta)\,\kappa(\vec{\theta}')$, where $W(\theta;\vartheta)$ is a smooth, continuous window function with a characteristic scale of $\vartheta$ ."
 Because of the commutativity between smoothing and the mass reconstruction (Van Waerbeke 2000). the simoothed quautities (4) and (5j satisfies the same relations as those between # ux 5: Vi=DOS=H (A0)=D2)ΟΠ... (divié;..rotáé;..).," Because of the commutativity between smoothing and the mass reconstruction (Van Waerbeke 2000), the smoothed quantities $\left<\kappa\right>$ and $\left<\gamma\right>$ satisfies the same relations as those between $\kappa$ and $\gamma$ : $\vec{\nabla}\left<\kappa\right>=\hat{D}\left<\gamma\right>
\equiv \vec{u}_{\left<\gamma\right>}$; $(\triangle\left<\kappa\right>,0)=
\hat{D}^2\left<\gamma\right>=\hat{D}\vec{u}_{\left<\gamma\right>}
=({\rm div}\vec{u}_{\left<\gamma\right>},
{\rm rot}\vec{u}_{\left<\gamma\right>})$ ."
 The Laplacian of (8). be. divis. is just the convergence convolved with the compcusated filter function ATT. aud hence equivalent to the aperture mass.," The Laplacian of $\left<\kappa\right>$, i.e. ${\rm div}\vec{u}_{\left<\gamma\right>}$, is just the convergence convolved with the compensated filter function $\triangle W$ , and hence equivalent to the aperture mass."
 Defining £4;:—Dej. we sec that the observable divi; traces the distribution of Xiu the luit of weal lIeusiug. as pointed out bv Lupping Kaiser (1997).," Defining $\vec{u}_{\left<\epsilon\right>}
:=\hat{D}\left<\epsilon\right>$, we see that the observable ${\rm div}\vec{u}_{\left<\epsilon\right>}$ traces the distribution of $\Sigma$in the limit of weak lensing, as pointed out by Luppino Kaiser (1997)."
 On the other hand. τονc measures the ‘pure’ noises," On the other hand, ${\rm rot}\vec{u}_{\left<\epsilon\right>}$ measures the `pure' noise."
" À discretized estimator for Dif;; is eiven by Dit.=(divéi.rotácο)Ξ (j. where εν and e, are the tangential aud the radial componeuts of the image ellipticity e=ει|fey defined by and (FM)=Slee27] with o= and p(0:0)=W""(0:d)Wd:3/0."," A discretized estimator for $\hat{D}\vec{u}_{\left<\gamma\right>}$ is given by $\hat{D}\vec{u}_{\left<\epsilon\right>}=
({\rm div}\vec{u}_{\left<\epsilon\right>},
 {\rm rot}\vec{u}_{\left<\epsilon\right>})=
-n_{\rm g}^{-1}\sum_{m=1}^{N_{\rm
g}}p(|\vec{\theta}-\vec{\theta}_m|;\vartheta)\,
[\epsilon_{\rm t}(\vec{\theta}_m;\vec{\theta}),
 \epsilon_{\rm r}(\vec{\theta}_m;\vec{\theta})] 
$ , where $\epsilon_{\rm t}$ and $\epsilon_{\rm r}$ are the tangential and the radial components of the image ellipticity $\epsilon=\epsilon_1+i\epsilon_2$ defined by and $\epsilon_{\rm r}(\vec{\theta};\vec{\theta}_0):=-
\Im[\epsilon(\vec{\theta})e^{-2i\phi}]$ with $\phi= {\rm Arg}(\vec{\theta}-\vec{\theta}_0)$ and $p(\theta;\vartheta)=W''(\theta;\vartheta)-W'(\theta;\vartheta)/\theta$."
" The noise propertics of D?fe) due to the intriusic source ellipticities are contaiued in the covariaut matrix στ,=|p?fel)Dp?Cel)|ii: where σε is the disperson of the intrinsic source ellipticities and we have assed. Ele,je(9,lijο70minOe2."," The noise properties of $\hat{D}^2\left<\epsilon\right>$ due to the intrinsic source ellipticities are contained in the covariant matrix $\sigma^2_{ij}:={\rm E}[
\hat{D}^2\left<\epsilon^{\rm (s)}\right>\,
\hat{D}^2\left<\epsilon^{\rm (s)}\right>]_{ij}$ : where $\sigma_{\epsilon}$ is the dispersion of the intrinsic source ellipticities and we have assumed ${\rm E}[\epsilon^{\rm (s)}(\vec{\theta}_m)
\epsilon^{\rm (s)}(\vec{\theta}_n)]_{ij}=\delta_{ij}\delta_{mn}
\sigma_{\epsilon}^2/2$ ."
 The Wrouceker delta ó;; iu equation (3)) ensures that the dispersions of divé;; aud roté44c are the sale and that the two fields are statistically uncorrelated., The Kronecker delta $\delta_{ij}$ in equation \ref{eq:sigma}) ) ensures that the dispersions of ${\rm div}\vec{u}_{\left<\epsilon\right>}$ and ${\rm rot}\vec{u}_{\left<\epsilon\right>}$ are the same and that the two fields are statistically uncorrelated.
 The local weal-leusine S/N v0) at position 0 is then defined by The resulting formula (eq. [1]]], The local weak-lensing S/N $\nu(\vec{\theta})$ at position $\vec{\theta}$ is then defined by The resulting formula (eq. \ref{eq:S/N}] ])
 is equivalent to the oue derived by Schucider (1996) using aperture mass statistics., is equivalent to the one derived by Schneider (1996) using aperture mass statistics.
 Tu the preseut Letter. we use a Cassian window function of the form Wed:0)=exp02/03E. ia which case p(0:d)=MO/dyexpt0/02)x04patidy: pat:d) has its maxima at Ó=J and /falls off= rapidly at 0>d.," In the present Letter, we use a Gaussian window function of the form $W_{\rm
G}(\theta;\vartheta)=\exp(-\theta^2/\vartheta^2)/\pi\vartheta^{2}$, in which case $p(\theta;\vartheta)=4(\theta/\vartheta)^2
\exp(-\theta^2/\vartheta^2)/\pi\vartheta^4\equiv p_{\rm
G}(\theta;\vartheta)$; $p_{\rm G}(\theta;\vartheta)$ has its maximum at $\theta =\vartheta$ and falls off rapidly at $\theta>\vartheta$."
" These statistics can be applied to the stroue-leusiug recie (o.@.. cluster central region). because the coutribution of nuaee cllipticitics to the aperture mass comes manly from ealaxies within au anuulus at radius ( aud thus we can avoid the strong Ieusine τοσο,"," These statistics can be applied to the strong-lensing regime (e.g., cluster central region), because the contribution of image ellipticities to the aperture mass comes mainly from galaxies within an annulus at radius $\vartheta$ and thus we can avoid the strong lensing regime."
 For each iudepeudoeutHST observation. we perform a local S/N. analysis using equation (1)]).," For each independent observation, we perform a local S/N analysis using equation \ref{eq:S/N}) )."
 To obtain the projected mass distribution. we perform a dmass reconstruction to the HST/WEC field.," To obtain the projected mass distribution, we perform a mass reconstruction to the /WFC field."
 Taking into account the hieh redslift (2= (0.897) of C 1601]BOL and small field-of-view (2/5 on a side} of the HST/WEC field. we adopt a non-linear fiuite-fiek inversion method developed by Seitz Schneider (1997). which takes account of the source redshift distribution.," Taking into account the high redshift $z=0.897$ ) of Cl 1604+4304 and small field-of-view $2\farcm 5$ on a side) of the /WFC field, we adopt a non-linear finite-field inversion method developed by Seitz Schneider (1997), which takes account of the source redshift distribution."
 Since little is known about the redshitt distribution of field ealaxies. weassume a source redshift distributiou of the form pi(:)=ioxp|(2a)!|T(3nui (Brainerd. Blauctord. Simail 1996). in which case the mean redshift (C) is given by (2)τιοο.," Since little is known about the redshift distribution of field galaxies, weassume a source redshift distribution of the form $p_{z}(z)=\beta z^2 \exp[-(z/z_0)^{\beta}]/\Gamma(3/\beta)z_0^3$ (Brainerd, Blandford, Smaili 1996), in which case the mean redshift $\left<z\right>$ is given by $\left<z\right>=z_0\Gamma(4/\beta)/\Gamma(3/\beta)$."
 Tn the presen Letter. we consider ouly the case (625.7)=(1.0.1.0).," In the present Letter, we consider only the case $(\left<z\right>, \beta)=(1.0, 1.0)$."
 Once a smoothed ellipticity field (ej(0) is obtained from the observed inage ellipticities. a convergence map can be obtained through the iutegral equation where quantities with oc-ubscrpt represent the values for sources at infinite redshift. #\ is the unknown coustaut whichrepreseuts the average ofας withinthe data field U. If is the kernel which is the eracieut of the scalar field that satisfies the Neumnaun boundary problem (see Seitz Schneider 1996). aud ος= D5&4.," Once a smoothed ellipticity field $\left<\epsilon\right>(\vec{\theta})$ is obtained from the observed image ellipticities, a convergence map can be obtained through the integral equation where quantities with $\infty$ -subscript represent the values for sources at infinite redshift, $\bar{\kappa}_{\infty}$ is the unknown constant whichrepresents the average of$\kappa_{\infty}$ withinthe data field $\cal U$ , $\vec{H}$ is the kernel which is the gradient of the scalar field that satisfies the Neumann boundary problem (see Seitz Schneider 1996), and $\vec{u}_{\gamma_{\infty}}=\hat{D}\gamma_{\infty}$ ."
 Iu general. the shear is not direct observable. so that the iutegral equation (5)) is nou-linear aud solved iteratively: A mass recoustruction," In general, the shear is not direct observable, so that the integral equation \ref{eq:MR}) ) is non-linear and solved iteratively: A mass reconstruction"
result does not come from a circular argument.,result does not come from a circular argument.
" The 2 0,4 test emplovs multiple diagnostics: the 58 jm AGN/SB spectral shapes. the AGN/SB ratios between the 58 jn and the jan emission. the correlation between the 30 sam and the total ULIRG luminosity."," The $R$ $\alpha_\mathit{bol}$ test employs multiple diagnostics: the 5–8 $\mu$ m AGN/SB spectral shapes, the AGN/SB ratios between the 5–8 $\mu$ m and the 8--1000 $\mu$ m emission, the correlation between the 30 $\mu$ m and the total ULIRG luminosity."
 If anv of these elements were a strong function of redshift. evidence lor either an ensemble devialion or dramatic outliers would be found.," If any of these elements were a strong function of redshift, evidence for either an ensemble deviation or dramatic outliers would be found."
 Ht is also worth noting. as in Fig. 11((," It is also worth noting, as in Fig. \ref{hz}( ("
"b). that our estimates of Ly, (reconstructed [rom a single far-IR point) are well matched to the tabulated values. (hat are computed from a broader band photometry. even if still limited.","b), that our estimates of $L_\mathit{IR}$ (reconstructed from a single far-IR point) are well matched to the tabulated values, that are computed from a broader band photometry, even if still limited."
 In more detail. the distribution of the hieh-redshift entries in the A αι plot suggests a small change of shA27. the bolometric correction lor SD-dominated sources.," In more detail, the distribution of the high-redshift entries in the $R$ $\alpha_\mathit{bol}$ plot suggests a small change of $R^\mathit{sb}$, the bolometric correction for SB-dominated sources."
 There are (wo possible explanations for this effect: 1) a missed AGN detection. due either to the bad quality of the single spectra or to a modification of the AGN/SB templates: 2) an underestimate of Lg. since the 30 yan flux does not properly represent the tvpical dust temperature of a SD environment.," There are two possible explanations for this effect: 1) a missed AGN detection, due either to the bad quality of the single spectra or to a modification of the AGN/SB templates; 2) an underestimate of $L_\mathit{IR}$, since the 30 $\mu$ m flux does not properly represent the typical dust temperature of a SB environment."
 The latter argument is perhaps (he most likely., The latter argument is perhaps the most likely.
 Evidence against dramatic spectral variations has been found also around 2—2.3. by applving our AGN/SD decomposition to the stacked spectra of 24 jan-selected sources and subamillimetre galaxies (Watabe et al.," Evidence against dramatic spectral variations has been found also around $z \simeq 2.3$, by applying our AGN/SB decomposition to the stacked spectra of 24 $\mu$ m-selected sources and submillimetre galaxies (Watabe et al."
 2009)., 2009).
 The assumption of AGN/SD templates and bolometrie corrections similar to the local ones leads to an average AGN content which is fully consistent with the main properties of hot populations (Sajina et al., The assumption of AGN/SB templates and bolometric corrections similar to the local ones leads to an average AGN content which is fully consistent with the main properties of both populations (Sajina et al.
 2007: et al., 2007; Men{\'e}nndez-Delmestre et al.
 2009)., 2009).
 We conclude that. as a first approximation. the SED large-scale properties of ULIRG-like svstems are not subject to significant evolution with redshilt.," We conclude that, as a first approximation, the SED large-scale properties of ULIRG-like systems are not subject to significant evolution with redshift."
 This hints to a variant of our diagnostics. where the fitting of templates to the mid-IB. spectra is replaced. with (he measurements of mid-Ilt colours or spectral slopes.," This hints to a variant of our diagnostics, where the fitting of templates to the mid-IR spectra is replaced with the measurements of mid-IR colours or spectral slopes."
 With the advent ol andHerschel. (he combined spectral and. photometric coverage will enable the measure of both the 3.8 jm rest-frame slope (see also Risalità et al.," With the advent of and, the combined spectral and photometric coverage will enable the measure of both the 3–8 $\mu$ m rest-frame slope (see also Risaliti et al."
 2010) and the bolometric correction: a simple D A diagnostic diagram will (hen provide the quantitative analvsis of much fainter I. sources in the deep fields., 2010) and the bolometric correction: a simple $\Gamma$ $R$ diagnostic diagram will then provide the quantitative analysis of much fainter IR sources in the deep fields.
 The Spitzer—-IRS unprecedented sensitivity allowecl a deeper investigation of the role οἱ supermassive black hole accretion and intense star [formation as the eneine underlving extreme IR activity., The -IRS unprecedented sensitivity allowed a deeper investigation of the role of supermassive black hole accretion and intense star formation as the engine underlying extreme IR activity.
 In particular. the 58 jan resi-l[rame wavelength range has proven to be," In particular, the 5–8 $\mu$ m rest-frame wavelength range has proven to be"
 2=D ;z20/—230. o ~10°AL. 26. τοL. à Z32 1.Syan 210. ~1000 arcmin7. Lya , $z=5$ $z\approx 20-30$ $\sigma$ $\sim 10^5~{\rm M_\odot}$ $z\gsim 6$ $z\gsim 4$ $\alpha$ $\gsim 32$ $1-5\mu$ $z\gsim 10$ $\sim1000$ $^{-2}$ $\alpha$ 
Ry distribution at Ηχος Iuminosity is smaller (han that predicted by analytic theory.,$R_d$ distribution at fixed luminosity is smaller than that predicted by analytic theory.
 Our simulations show that. under the influence of a bar. A24; may. increase hy a factor of 2 or more at constant global angular momentum.," Our simulations show that, under the influence of a bar, $R_d$ may increase by a factor of 2 or more at constant global angular momentum."
 Since less extended clisks are likely to be more bar-unstable. (he secular evolution of (hese disks may be responsible for at least part of this discrepancy.," Since less extended disks are likely to be more bar-unstable, the secular evolution of these disks may be responsible for at least part of this discrepancy."
 We would like to thank the anonymous referee for useful suggestions., We would like to thank the anonymous referee for useful suggestions.
ancl a X7 estimator. we obtain 2=0.4340.05. 3—045+ for SCDM. and. B=0.42x 0.05. 3=0.56x0.10 for ACDAL,"and a $\chi^2$ estimator, we obtain $B=0.43\pm0.05$, $\beta=0.45\pm0.09$ for SCDM, and $B=0.42\pm0.05$ , $\beta=0.56\pm0.10$ for $\Lambda$ CDM."
 Using the BCESCY.Xa) estimator of Akritas Bershacly (1996). which accounts for errors in both axes and he presence of possible intrinsic scatter. we obtain ο=VASdE 0.07. 9=0.4630.08 for SCDM and 2B=0.51xz0.05. j—048d0.06 for ACDAL.," Using the $X_2|X_1$ ) estimator of Akritas Bershady (1996), which accounts for errors in both axes and the presence of possible intrinsic scatter, we obtain $B=0.48 \pm 0.07$ , $\beta=0.46 \pm0.08$ for SCDM and $B=0.51 \pm 0.05$, $\beta=0.48\pm0.06$ for $\Lambda$ CDM."
 We conclude that. the slopx of the. temperature-uminosity relation for the present. sample of hot. relaxed clusters is consistent with the predicted. value of 3=0.5 (Section 1).," We conclude that the slope of the temperature-luminosity relation for the present sample of hot, relaxed clusters is consistent with the predicted value of $\beta=0.5$ (Section 1)."
 Fixing 3?=0.33 results in a poor fit: V7=12.5 .or 5 degrees of ⋅⋅[reedom. as opposed to V72=6.7 withH —0.5 (ΑςΟΝΕ.," Fixing $\beta=0.33$ results in a poor fit: $\chi^2=12.5$ for 5 degrees of freedom, as opposed to $\chi^2=6.7$ with $\beta=0.5$ $\Lambda$ CDM)."
 We have shown that within ασ reso. Corresponding to a fixecl density contrast A=2500 with respect to the critical density at the redshifts of the clusters. the temperature profiles for the present sample of Luminous. relatively relaxed lensing clusters exhibit an approximately universal form which rises within roo~OBreso) anel then remains approximately constant oul to. reso.," We have shown that within radii $r_{2500}$, corresponding to a fixed density contrast $\Delta =2500$ with respect to the critical density at the redshifts of the clusters, the temperature profiles for the present sample of luminous, relatively relaxed lensing clusters exhibit an approximately universal form which rises within $r \sim 0.3\, r_{2500}$ and then remains approximately constant out to $r_{2500}$."
 “Phe enclosed. masses. bolometric luminosities and mean gas mass-weighted temperatures within these radii scale. in manner consistent with the predictions from the. simple virial relations outlined in Section 1.," The enclosed masses, bolometric luminosities and mean gas mass-weighted temperatures within these radii scale in manner consistent with the predictions from the simple virial relations outlined in Section 1."
 We have confirmed the presence of a svstematic ollset of ~40 per cent between the normalizations of the observed ancl predicted Mozou—οπου curves. in the sense that the predicted. temperatures are too low for a given mass. for both the SCDAL and [CDM cosmologies.," We have confirmed the presence of a systematic offset of $\sim 40$ per cent between the normalizations of the observed and predicted $M_{2500}-kT_{2500}$ curves, in the sense that the predicted temperatures are too low for a given mass, for both the SCDM and $\Lambda$ CDM cosmologies."
 An important aspect of the present study. is. that independent confirmation of the X-ray mass measurements is available from gravitational lensing studies., An important aspect of the present study is that independent confirmation of the X-ray mass measurements is available from gravitational lensing studies.
 For both bell 2390 and 15NJ1347-1145. the X-ray and. weak lensing mass profiles are consistent. within their GS por cent confidence limits.," For both Abell 2390 and RXJ1347-1145, the X-ray and weak lensing mass profiles are consistent within their 68 per cent confidence limits."
 For Abell 1835. 2390. M82137-2353 and. PINSQOT45-191. the observed. strong lensing configurations (on scales r~20SO tkpe) can be explained by mass models within the 68 per cent Chandra: confidence contours. although redshift measurements for the ares (which are required. to define the lensing masses precisely) are not available in all Thus. the presence of significant non-thermal pressure support arising from turbulent and/or bulk motions and/or magnetic Lelds) can be excluded.," For Abell 1835, 2390, MS2137-2353 and PKS0745-191, the observed strong lensing configurations (on scales $r \sim 20-80\,h^{-1}$ kpc) can be explained by mass models within the 68 per cent Chandra confidence contours, although redshift measurements for the arcs (which are required to define the lensing masses precisely) are not available in all Thus, the presence of significant non-thermal pressure support arising from turbulent and/or bulk motions and/or magnetic fields) can be excluded."
 We conclude that the systematic uncertainties associated. with the individual mass measurements are small (<20 per cent)., We conclude that the systematic uncertainties associated with the individual mass measurements are small $<20$ per cent).
 The offset. between the observed. and. simulated: mass-temperature curves cannot be explained. by invoking an earlier. formation redshift for the observed. clusters (ve assume that the clusters form at the redshifts they are observed) since. for the measured NEW mass distributions. AMozou(2) drops as fast or faster than ος). rises às the formation redshift is increased.," The offset between the observed and simulated mass-temperature curves cannot be explained by invoking an earlier formation redshift for the observed clusters (we assume that the clusters form at the redshifts they are observed) since, for the measured NFW mass distributions, $M_{2500}(z)$ drops as fast or faster than $E(z)$ rises as the formation redshift is increased."
 Our results suggest. that on the spatial scales studied: here. important physics may be missing from the reference simulations.," Our results suggest that on the spatial scales studied here, important physics may be missing from the reference simulations."
 One possible candidate is radiative cooling of the N-rav gas. which the Chandra cata show to be significant. within r~0.2rozoo Allen 2001a.b.c: David 2001: Schmidt 2001).," One possible candidate is radiative cooling of the X-ray gas, which the Chandra data show to be significant within $r
\sim 0.2 r_{2500}$ Allen 2001a,b,c; David 2001; Schmidt 2001)."
 Pearce (2000) show that the introduction of radiative cooling into their hvdrodynamical simulations can lead to central temperature drops similar to those in Fig. Ll.., Pearce (2000) show that the introduction of radiative cooling into their hydrodynamical simulations can lead to central temperature drops similar to those in Fig. \ref{fig:kt}.
 Phese authors also argue that cooling can. [ead to a significant increase in the mass-weighted temperature within rresus (às cooled. low-entropy gas is deposited ancl warmer. high-entropy material Lows inwards and is compressed). which may be sullicient to account for. the discrepancy between the observed ancl sipiulated: curves.," These authors also argue that cooling can lead to a significant increase in the mass-weighted temperature within $r \sim r_{2500}$ (as cooled, low-entropy gas is deposited and warmer, high-entropy material flows inwards and is compressed), which may be sufficient to account for the discrepancy between the observed and simulated curves."
 Detailec simulations of the Mozuo—A500 relation for large a sample of massive clusters. including the elfects of radiative cooling. are required to address this issue.," Detailed simulations of the $M_{2500}-kT_{2500}$ relation for large a sample of massive clusters, including the effects of radiative cooling, are required to address this issue."
 The results presented in this paper should. provide a useful calibrator for future studies of the X-ray. properties of galaxy clusters., The results presented in this paper should provide a useful calibrator for future studies of the X-ray properties of galaxy clusters.
 In future work we will examine the constraints that the present data place on radial variations in the X-ray gas mass fraction in theclusters andl. therefore; Ou.," In future work we will examine the constraints that the present data place on radial variations in the X-ray gas mass fraction in theclusters and, therefore, $\Omega_{\rm m}$ ."
 We will also explore the ability of dilferent parameterized mass mocels to explain the observed: X-ray gas temperature and. density proliles., We will also explore the ability of different parameterized mass models to explain the observed X-ray gas temperature and density profiles.
 SWA and ACE acknowledge the support. of the Itoval Society., SWA and ACF acknowledge the support of the Royal Society.
distribution modified by cluster disruption processes.,distribution modified by cluster disruption processes.
 We recall that based on our magnitude-Iimited sample analysis (Fig., We recall that based on our magnitude-limited sample analysis (Fig.
 6bb). we also concluded that the initial CALL slope appeared to be significantly shallower. at the 30 level. than the power-law slope. à=2. expected for voung star eluster πμο...," \ref{agemasshist.fig}b b), we also concluded that the initial CMF slope appeared to be significantly shallower, at the $3\sigma$ level, than the power-law slope, $\alpha =
-2$, expected for young star cluster systems."
 Since we observe this behaviour towards the low-mass end for all mass anc age subsets of the full mass-Iimited cluster sample that include masses logCA4/M.)Ἑ3. as well as in our magnittcle-limited sample. random stochastic ellects are unlikely to be the primary cause (see also Cirarcli Bica 1993: Santos Frogel 1997).," Since we observe this behaviour towards the low-mass end for all mass and age subsets of the full mass-limited cluster sample that include masses $\log(M_{\rm cl}/{\rm M}_\odot) \lesssim 3$, as well as in our magnitude-limited sample, random stochastic effects are unlikely to be the primary cause (see also Girardi Bica 1993; Santos Frogel 1997)."
 Lt is likely that this is a real elfect (see also Elmegreen Efremov 1997). and that the CME slopes flatten significantly for vounger ages and lower-mass clusters in the LAIC.," It is likely that this is a real effect (see also Elmegreen Efremov 1997), and that the CMF slopes flatten significantly for younger ages and lower-mass clusters in the LMC."
 Elmegreen Efremov (1907) argue that the vounger clusters are mostly. unbound OB associations (supporting Sica et al., Elmegreen Efremov (1997) argue that the younger clusters are mostly unbound OB associations (supporting Bica et al.
 1996). of which some 90 per cent will clisperse by the time their constituent stars will reach an age of ~107 vr.," 1996), of which some 90 per cent will disperse by the time their constituent stars will reach an age of $\sim 10^8$ yr."
" This is consistent with modern ideas on the formation and dissolution of star clusters within the first ~LO"" vr of their existence: most (7090 per cent) of these newly formed clusters will disperse on these time-scales. a process coined ""infant mortality (e.g... Boily Ixroupa 2003: Vesperini Zepf 2003: Whitmore 2004: Bastian et al."," This is consistent with modern ideas on the formation and dissolution of star clusters within the first $\sim 10^7$ yr of their existence; most $\sim 70-90$ per cent) of these newly formed clusters will disperse on these time-scales, a process coined “infant mortality” (e.g., Boily Kroupa 2003; Vesperini Zepf 2003; Whitmore 2004; Bastian et al."
 2005: Mengel et al., 2005; Mengel et al.
 2005: see also Tremonti ct al., 2005; see also Tremonti et al.
 2001)., 2001).
 LP the process of infant mortality is mass dependent. (but sce Whitmore 2004 for counterarguments). in the sense that the lowest-mass clusters will dissolve. preferentially. the result. will be a of the CALF slopes with increasing mean age. which is contrary to the apparent change of slope in Fig. SN.," If the process of infant mortality is mass dependent (but see Whitmore 2004 for counterarguments), in the sense that the lowest-mass clusters will dissolve preferentially, the result will be a of the CMF slopes with increasing mean age, which is contrary to the apparent change of slope in Fig. \ref{clfs2.fig}."
 Whether or not the voungest C510 Myr old) clusters will dissolve because of this infant. mortality scenario. the significant dilferences among the LMCS CALF slopes as à function of age are predominantly driven by the voungest subset of our cluster sample.," Whether or not the youngest $\lesssim 10$ Myr old) clusters will dissolve because of this infant mortality scenario, the significant differences among the LMC's CMF slopes as a function of age are predominantly driven by the youngest subset of our cluster sample."
 Thus. these results may imply. that the initial CME slope of thecombined. (Le. both bound. and unbound) LAIC cluster system is well represented. by à. power-law. although we cannot disentangle the unbound from the bound clusters at the voungest ages.," Thus, these results may imply that the initial CMF slope of the (i.e., both bound and unbound) LMC cluster system is well represented by a power-law, although we cannot disentangle the unbound from the bound clusters at the youngest ages."
 In addition. we recently presented observational evidence and theoretical arguments against an initial power-law CALF in MS2's intermecdiate-age starburst region. MS2. D (de Ciijs. Parmentier Lamers 2005). and in favour of an initial log-normal. €ME. in the Antennae interacting system. NGC 4038/39. (Ancers et al.," In addition, we recently presented observational evidence and theoretical arguments against an initial power-law CMF in M82's intermediate-age starburst region, M82 B (de Grijs, Parmentier Lamers 2005), and in favour of an initial log-normal CMF in the Antennae interacting system, NGC 4038/39 (Anders et al."
 2005)., 2005).
 Our detailed. analvsis of the LMC cluster mass distributions as a function of age and mass may therefore have uncovered supporting new evidence that star clusters at least in the low-density environment. of the LMC may not form following a power-law distribution that holds strength down to masses much below a few 10* M. where the power-law fits to the subsamples in Fig.," Our detailed analysis of the LMC cluster mass distributions as a function of age and mass may therefore have uncovered supporting new evidence that star clusters – at least in the low-density environment of the LMC – may not form following a power-law distribution that holds strength down to masses much below a few $\times 10^3$ $_\odot$, where the power-law fits to the subsamples in Fig."
 S. break down., \ref{clfs2.fig} break down.
 ὃν combining integrated properties with resolved: stellar population studies. the LMC cluster svstem olfers the unique chance to independently. check the accuracy of age (and corresponding mass) determinations based on broad-band SEDs.," By combining integrated properties with resolved stellar population studies, the LMC cluster system offers the unique chance to independently check the accuracy of age (and corresponding mass) determinations based on broad-band SEDs."
 In this paper. we have reanalyzed the broad-band LMC cluster SEDs based on the cata of Massey. (2002) and 1103. using a newly developed SED analysis approach.," In this paper, we have reanalyzed the broad-band LMC cluster SEDs based on the data of Massey (2002) and H03, using a newly developed SED analysis approach."
 We compare our new age determinations with (i) those of LI03 using the same data set but a dillerent approach. and. (ii) those of Pictrzvisski Udalski (2000) using CAID fitting. in order to set the tightest limits vet on the accuracy of (absolute) age determinations based. on broad-band SEDs. and therefore on the usefulness of such an approach.," We compare our new age determinations with (i) those of H03 using the same data set but a different approach, and (ii) those of Pietrzyńsski Udalski (2000) using CMD fitting, in order to set the tightest limits yet on the accuracy of (absolute) age determinations based on broad-band SEDs, and therefore on the usefulness of such an approach."
 We note a significant svstematic effect. between the age cdillerences. of H03 on the one hand. and those of both the OGLE-L team and our own redeterminations on the other.," We note a significant systematic effect between the age differences of H03 on the one hand, and those of both the OGLE-II team and our own redeterminations on the other."
 Lt appears that these systematic dillerences are caused by 1095 conversions of the photometry to a cillerent filter system., It appears that these systematic differences are caused by H03's conversions of the photometry to a different filter system.
 We emphasize and warn that theacfiial filter systems used. for the observations should. be used. for the most accurate parameter analvsis. instead of using filter conversion equations. in order to achieve more accurate derivations of the cluster ages and the corresponding masses.," We emphasize and warn that the filter systems used for the observations should be used for the most accurate parameter analysis, instead of using filter conversion equations, in order to achieve more accurate derivations of the cluster ages and the corresponding masses."
 Based on this comparison. and additionally on a cetailed assessment of the age-metallicity and age-extinction degeneracies. we conclude that our. broad-band SED fits viele reliable ages. with statistical uncertainties within Alog(Agesvr)20.4 overall.," Based on this comparison, and additionally on a detailed assessment of the age-metallicity and age-extinction degeneracies, we conclude that our broad-band SED fits yield reliable ages, with statistical uncertainties within $\Delta\log( \mbox{Age/yr}) \simeq 0.4$ overall."
" ""Thus. in addition to our conclusion in de Cirijs et al. ("," Thus, in addition to our conclusion in de Grijs et al. ("
2005) that we can retrieve prominent features in the cluster age distribution to within ελΙουλσον)ες0.35 using a variety of approaches based. on broad-band SEDs modelling. we have now also shown that the associated.κό statistical uncertainties involved in cluster age determinations are ofa very similar magnitucoe.,"2005) that we can retrieve prominent features in the cluster age distribution to within $\Delta \langle \log( {\rm Age / yr} ) \rangle
\le 0.35$ using a variety of approaches based on broad-band SEDs modelling, we have now also shown that the associated statistical uncertainties involved in cluster age determinations are ofa very similar magnitude."
 The LAIC’s CER has been roughly constant outside of the well-known age gap between 3 and 13 Car. when the CER was a factor of ~5 lower (assuming a roughly constant rate during the entire period).," The LMC's CFR has been roughly constant outside of the well-known age gap between $\sim 3$ and 13 Gyr, when the CFR was a factor of $\sim 5$ lower (assuming a roughly constant rate during the entire period)."
 There are no clear observational signatures of an enhanced. CER. associated with the [ast tidal encounter between the LMC. and the SAIC. while we argue that the combination of the relevant time-scales. Le. the LAIC’s rotation period and the time since the last LMC-SMC encouter. has been insullicicnt to wash out any such signatures. if they had been present.," There are no clear observational signatures of an enhanced CFR associated with the last tidal encounter between the LMC and the SMC, while we argue that the combination of the relevant time-scales, i.e., the LMC's rotation period and the time since the last LMC-SMC encouter, has been insufficient to wash out any such signatures, if they had been present."
 An alternative triggering mechanism for the voung(er) clusters may be neeceel., An alternative triggering mechanism for the young(er) clusters may be needed.
" Using a simple approach to derive the characteristic cluster disruption time-scale. we [find that log(P/vr)= 0.1. where fan=CIP(UL4U/104M. ""7. for the EMC cluster system."," Using a simple approach to derive the characteristic cluster disruption time-scale, we find that $\log(t_4^{\rm dis}/{\rm yr}) =
9.9 \pm 0.1$ , where $t_{\rm dis} = t_4^{\rm dis} (M_{\rm cl}/10^4 {\rm
M}_\odot)^{0.62}$ , for the LMC cluster system."
 This is consistent with earlier. preliminary," This is consistent with earlier, preliminary"
however it is double-valued at low abundances.,however it is double-valued at low abundances.
" Note that the 2005 observations could only be corrected for extinction by measuring the H8 to Hy ratio, since Ha was not available."," Note that the 2005 observations could only be corrected for extinction by measuring the $H\beta$ to $H\gamma$ ratio, since $H\alpha$ was not available."
 This provides a poorer correction than the Ha to Hf ratio., This provides a poorer correction than the $H\alpha$ to $H\beta$ ratio.
" In contrast to thep-method, the [O 1u]/[N 11 Stasifisska's method is single valued, with the advantage that it is relatively independent of the reddening (?).."," In contrast to the, the [O ]/[N ] Stasińsska's method is single valued, with the advantage that it is relatively independent of the reddening \citep{Stasinska06}."
 A constant gradient has been the most commonly adopted law in the study of metallicity of galaxies., A constant gradient has been the most commonly adopted law in the study of metallicity of galaxies.
" This is a first approximation, not a choice dictated by theory."," This is a first approximation, not a choice dictated by theory."
" For instance, most chemical abundance models adopt star formation rates proportional to some power of the gas density, and most often the gas density presents a peak at several kpc from the center."," For instance, most chemical abundance models adopt star formation rates proportional to some power of the gas density, and most often the gas density presents a peak at several kpc from the center."
 The stellar density in the disk of our Galaxy is better described by Kormendy's function than by an exponential law (see ?))., The stellar density in the disk of our Galaxy is better described by Kormendy's function than by an exponential law (see \citealt{Lepine00}) ).
 We make here the choice of fitting the data with 4th order polynomials., We make here the choice of fitting the data with 4th order polynomials.
" This does not correspond to a theoretical model of galactic structure; the polynomial fitting is only adopted as a smoothing method instead other methods like splines or Gaussian filtering, which could be used alternatively."," This does not correspond to a theoretical model of galactic structure; the polynomial fitting is only adopted as a smoothing method instead other methods like splines or Gaussian filtering, which could be used alternatively."
 One advantage of the polynomial method is that it gives the position of minima and inflections in an easy way., One advantage of the polynomial method is that it gives the position of minima and inflections in an easy way.
 The choice of the 4th order corresponds to the amount of details that we are willing to reveal., The choice of the 4th order corresponds to the amount of details that we are willing to reveal.
" For instance in our Galaxy, the observations of Cepheids (?)) show that there is a strong metallicity gradient in the inner regions, followed by a plateau, followed again by a another strong gradient."," For instance in our Galaxy, the observations of Cepheids \citealt{Andrievsky04}) ) show that there is a strong metallicity gradient in the inner regions, followed by a plateau, followed again by a another strong gradient."
" Such a behavior, that we expect to observe in other galaxies as well, can be approximated by a 4th order polynomial."," Such a behavior, that we expect to observe in other galaxies as well, can be approximated by a 4th order polynomial."
" It is known that the different statistical methods used to calculate the abundances can result in different values for a same observation, and they can produce different dispersion of the data with respect to an average abundance profile, and also small differences in the gradients, specially for those methods which do not depend directly on theO ionic abundances."," It is known that the different statistical methods used to calculate the abundances can result in different values for a same observation, and they can produce different dispersion of the data with respect to an average abundance profile, and also small differences in the gradients, specially for those methods which do not depend directly on the ionic abundances."
 In our observations the largest differences occurred between thep-method and all the other mentioned methods., In our observations the largest differences occurred between the and all the other mentioned methods.
" This seems to be a consequence of the strong dependence of thep-method on the [O 11] emission line, which is observed at the limit of the instrumental sensitivity and it is strongly affected by the reddening and the parallactic angle misalignments."," This seems to be a consequence of the strong dependence of the on the [O ] emission line, which is observed at the limit of the instrumental sensitivity and it is strongly affected by the reddening and the parallactic angle misalignments."
" For IC0167 most of the methods applied here point towards the existence of a very smooth minimum in theO abundance profile, which is shallower than the one of our Galaxy, according to ?).."," For IC0167 most of the methods applied here point towards the existence of a very smooth minimum in the abundance profile, which is shallower than the one of our Galaxy, according to \cite*{Mishurov02}."
" In particular, the distribution produced by thep-method is more dispersed than the results for the [O iJ/[N Π] method, however the minima in theO abundance obtained with the two methods almost coincide."," In particular, the distribution produced by the is more dispersed than the results for the [O ]/[N ] method, however the minima in the abundance obtained with the two methods almost coincide."
 'The median of the minima and inflections found using all the methods applied to this galaxy is 14.7+2.6 kpc., The median of the minima and inflections found using all the methods applied to this galaxy is $14.7 \pm 2.6$ kpc.
 Obviously it is necessary to be cautious in the interpretation of the fitted curves., Obviously it is necessary to be cautious in the interpretation of the fitted curves.
 The trend could be fitted by a straight line and the evidence for an inflection is not very strong in this case., The trend could be fitted by a straight line and the evidence for an inflection is not very strong in this case.
" On the other hand, for NGC1042, the abundance dispersion is less pronounced and a plateau is present in theO abundance profiles of this galaxy obtained by means of the [O 11]/[N 1]p and R23-based methods."," On the other hand, for NGC1042, the abundance dispersion is less pronounced and a plateau is present in the abundance profiles of this galaxy obtained by means of the [O ]/[N ] and R23-based methods."
" This plateau could possibly be interpreted as being a consequence of the corotation as well, but in this case the minimum is too shallow to contrast with the dispersion of the data."," This plateau could possibly be interpreted as being a consequence of the corotation as well, but in this case the minimum is too shallow to contrast with the dispersion of the data."
 The dispersion resulting from the two calibrations of [N u]/Ha method of ?) and the [Ar 111]/[O 1] is so high that no consistent information can be extracted., The dispersion resulting from the two calibrations of [N $H\alpha$ method of \cite{PP04} and the [Ar ]/[O ] is so high that no consistent information can be extracted.
" Since the point of inflection is a point which is close to being a minimum (see the lower curves in Figure 6 for NGC1042), we could speculate that it is also an indicator of thecorotation."," Since the point of inflection is a point which is close to being a minimum (see the lower curves in Figure \ref{fig6} for NGC1042), we could speculate that it is also an indicator of thecorotation."
 The median of the inflexions of the abundance curves for NGC1042 is 8.2+2.8 kpc., The median of the inflexions of the abundance curves for NGC1042 is $8.2 \pm 2.8$ kpc.
" Interestingly, at this same position a small bump also can be observed in the color distribution of this galaxy (Figure 4))."," Interestingly, at this same position a small bump also can be observed in the color distribution of this galaxy (Figure \ref{fig4}) )."
" This might be connected with corotation as well, since a gap in the density distribution of young stars can be expected at corotation, as it happens in our Galaxy (?).."," This might be connected with corotation as well, since a gap in the density distribution of young stars can be expected at corotation, as it happens in our Galaxy \citep*{Amores09}."
 The lack of young stars would make this region slightly redder than its neighborhood., The lack of young stars would make this region slightly redder than its neighborhood.
" Finally, a bimodal behavior in the radial distribution ofO is found for NGC6907 using the [O iu]/[N τῇ method."," Finally, a bimodal behavior in the radial distribution of is found for NGC6907 using the [O ]/[N ] method."
" Since the Argonium line is not detected along all the galaxy, the spatial distribution of metallicity derived from this element is inconclusive."," Since the Argonium line is not detected along all the galaxy, the spatial distribution of metallicity derived from this element is inconclusive."
" The same happens using the [N u]/Ha method, which results inO abundances with large dispersion at radii beyond 15 kpc."," The same happens using the [N $H\alpha$ method, which results in abundances with large dispersion at radii beyond 15 kpc."
" Unfortunately the observations of 2005 were not performed near the parallactic angle and for this reason the measurements of the [O iij lines are not reliable, causing a huge dispersion of theO abundances derived by the R23 andp-method."," Unfortunately the observations of 2005 were not performed near the parallactic angle and for this reason the measurements of the [O ] lines are not reliable, causing a huge dispersion of the abundances derived by the R23 and."
 Since the nucleosynthetic origin of the nitrogen may shift from secondary to primary in low abundances regions we need to be more cautious to interpret the inflexions in the NGC6907 abundance distribution., Since the nucleosynthetic origin of the nitrogen may shift from secondary to primary in low abundances regions we need to be more cautious to interpret the inflexions in the NGC6907 abundance distribution.
 The median of the radii of the minima found using the other statistical methods is 21.13.7 kpc., The median of the radii of the minima found using the other statistical methods is $21.1 \pm 3.7$ kpc.
" ?) have shown that for this galaxy there is a minimum in the gas distribution between 20 and 30 kpc, which can be relatedI with the same effect of the corotation found in our galaxy by ?).."," \cite{Sca2008a} have shown that for this galaxy there is a minimum in the gas distribution between 20 and 30 kpc, which can be related with the same effect of the corotation found in our galaxy by \cite{Amores09}."
" The interpretation of the gradient of metallicity for NGC6907 may seem dangerous, since this galaxy is interacting with NGC6908 as discussed by ?).."," The interpretation of the gradient of metallicity for NGC6907 may seem dangerous, since this galaxy is interacting with NGC6908 as discussed by \cite{Sca2008a}."
" However the same authors have shown that the influence of this galaxy on the NGC6907 gas distribution is restricted to about 20°around NGC6908, in azimuthal angles measured in NGC6907 galactic plane."," However the same authors have shown that the influence of this galaxy on the NGC6907 gas distribution is restricted to about around NGC6908, in azimuthal angles measured in NGC6907 galactic plane."
" Except for the spectra sampled by the slit 25 in the 2006 run, all the other slits associated to radii greater than 15 kpc are situated on the side opposite to the one where the interaction with NGC6908 takes place."," Except for the spectra sampled by the slit 25 in the 2006 run, all the other slits associated to radii greater than 15 kpc are situated on the side opposite to the one where the interaction with NGC6908 takes place."
" Similarly to what happens with the rotation curve, which is not perturbed on that side (?), we expect that the metallicities of the regions presented here are not affected by the interaction."," Similarly to what happens with the rotation curve, which is not perturbed on that side \citep{Sca2008a}, we expect that the metallicities of the regions presented here are not affected by the interaction."
 Are the polynomial fits really better than the traditional straight line fits?, Are the polynomial fits really better than the traditional straight line fits?
" It should be remembered that we do not consider that the 4th order polynomials are models, but only a technique used for locating changes in the slope."," It should be remembered that we do not consider that the 4th order polynomials are models, but only a technique used for locating changes in the slope."
" If the changes in slopes and plateaux are real then, in principle, the polynomial fit should be better."," If the changes in slopes and plateaux are real then, in principle, the polynomial fit should be better."
" But of course, a 4th order polynomial always produces a better fit than straight line, as long as the rms deviations of the data with respect to the fitted curve are taken as the measure of the quality of the fit."," But of course, a 4th order polynomial always produces a better fit than straight line, as long as the rms deviations of the data with respect to the fitted curve are taken as the measure of the quality of the fit."
" To determine if a polynomial fit is effectively a meaningful choice, one must compare the chi-square divided by the degree of freedom N—k 1, where N is the number of data points and k the order of the polynomial (see the"," To determine if a polynomial fit is effectively a meaningful choice, one must compare the chi-square divided by the degree of freedom $N-k-1$ , where $N$ is the number of data points and $k$ the order of the polynomial (see the"
"quantities in. the absorption. profiles could also be partly affected by relatively strong telluric lines at 6560.555 A and 6564.206 A,",quantities in the absorption profiles could also be partly affected by relatively strong telluric lines at 6560.555 $\AA$ and 6564.206 $\AA$.
 Our hypothesis and results discussed above are potentially good for determining the physical properties of the disk (Le. its density and velocity structure). which could help in understanding its nature.," Our hypothesis and results discussed above are potentially good for determining the physical properties of the disk (i.e., its density and velocity structure), which could help in understanding its nature."
 But for a precise analysis of this kind. one should model the light of the primary going through the various model disks and compare the results with the observed profiles.," But for a precise analysis of this kind, one should model the light of the primary going through the various model disks and compare the results with the observed profiles."
" This task is beyond the scope of this, rather qualitative, analysis."," This task is beyond the scope of this, rather qualitative, analysis."
" As already noted at the beginning of this section, the CI of the absorption core starts to decrease at least three years before the predicted beginning of primary eclipse."," As already noted at the beginning of this section, the CI of the absorption core starts to decrease at least three years before the predicted beginning of primary eclipse."
 This observational result can be used to constrain the structure of the system., This observational result can be used to constrain the structure of the system.
 We refer to the period when absorption cores deepen below their out-of-eclipse mean value as “spectroscopiceclipse’., We refer to the period when absorption cores deepen below their out-of-eclipse mean value as spectroscopic.
 From Fig. 7..," From Fig. \ref{Ha_orig},"
 we estimate that spectroscopic eclipse began around HID 2454000., we estimate that spectroscopic eclipse began around HJD 2454000.
 The ?. orbital solution implies that the angle between the sightline at the start of spectroscopic eclipse and the mideclipse sightline 1s — 857., The \citet{Chadima2010} orbital solution implies that the angle between the sightline at the start of spectroscopic eclipse and the mideclipse sightline is $\sim$ $^\circ$.
 JH determined that the 2009 photometric eclipse began on HJD 2455056 when the sightline angle was only ~20° relative to the mideclipse line of sight., JH determined that the 2009 photometric eclipse began on HJD 2455056 when the sightline angle was only $\sim$ $^\circ$ relative to the mideclipse line of sight.
 In Fig. 12..," In Fig. \ref{Roche},"
" we show the critical Roche lobes for the high- and low-mass models presented in Sect. ??.,"," we show the critical Roche lobes for the high- and low-mass models presented in Sect. \ref{intro},"
 and the sightlines at the onset of spectroscopic and photometric eclipse., and the sightlines at the onset of spectroscopic and photometric eclipse.
 It is evident that the additional absorption at the. start of spectroscopic eclipse cannot be caused by material near the secondary because the binary orientation. is. close to maximum separation at that time., It is evident that the additional absorption at the start of spectroscopic eclipse cannot be caused by material near the secondary because the binary orientation is close to maximum separation at that time.
 Therefore there must be some circumbinary material responsible for this additional absorption. suggestive of the ?| model.," Therefore there must be some circumbinary material responsible for this additional absorption, suggestive of the \citet{Struve1956} model."
" However. trom our observations. 1t Is seen that this envelope is not homogenous, as proposed by ?.."," However, from our observations, it is seen that this envelope is not homogenous, as proposed by \citet{Struve1956}."
 The line of sight at the onset of photometric eclipse intersects both critical. Roche lobes around the secondary., The line of sight at the onset of photometric eclipse intersects both critical Roche lobes around the secondary.
 However. in the case of the high-mass model. the disk around the secondary would almost extend to the critical Roche lobe.," However, in the case of the high-mass model, the disk around the secondary would almost extend to the critical Roche lobe."
 An unusual and prominent change in the line profile was observed during 2005-2006. long before the onset of spectroscopic eclipse.," An unusual and prominent change in the line profile was observed during 2005–2006, long before the onset of spectroscopic eclipse."
 The blue emission wing disappeared and was replaced by a deep. blueshifted absorption core (Fig. 13..," The blue emission wing disappeared and was replaced by a deep, blueshifted absorption core (Fig. \ref{outburst},"
 also Fig. 2))., also Fig. \ref{ha_fig1}) ).
" Initially, the first spectrum (in Apr 2005) has a profile. while the second and the third spectra observed later that year show the gradual appearance of the blueshitted absorption core."," Initially, the first spectrum (in Apr 2005) has a profile, while the second and the third spectra observed later that year show the gradual appearance of the blueshifted absorption core."
 By Mar 2007. the final spectrum of this series," By Mar 2007, the final spectrum of this series"
"Notice, though. the three separately tageed UCDs iu Figure 6..","Notice, though, the three separately tagged UCDs in Figure \ref{fig:amr}."
 These correspond to the bright. red UCDs hat we linked earlier to the stripped remnants of iore nassive ealaxies: they are again quite distinct from the eoncral run of dE nuclei aud UCDs.," These correspond to the bright, red UCDs that we linked earlier to the stripped remnants of more massive galaxies; they are again quite distinct from the general run of dE nuclei and UCDs."
 Another intriguing result from Figure ο ds that he AMIRs of the faint and bright dE πο] may be systematically differcut in the seuse that the faint nuclei jiwe lower iietallicity at a eiven age (equivalent to a inass-inetallicity correlation at each age) aud lower o at a given metallicitv., Another intriguing result from Figure \ref{fig:amr} is that the AMRs of the faint and bright dE nuclei may be systematically different in the sense that the faint nuclei have lower metallicity at a given age (equivalent to a mass-metallicity correlation at each age) and lower $\alpha$ /Fe] at a given metallicity.
 The trends for the UCDs are fFo]nof clear from the existing data., The trends for the UCDs are not clear from the existing data.
 Spectroscopic analyses are needed iu particulary for the new class of ow-luninosity object. as well as for ordinary compact GC's around. M8ST.," Spectroscopic analyses are needed in particular for the new class of low-luminosity object, as well as for ordinary compact GCs around M87."
 Futher discussion of the stellar populations Huplications in a wider context will be provided iu Section ?7.., Further discussion of the stellar populations implications in a wider context will be provided in Section \ref{sec:disc}.
 Understanding the origins of UCDs will ultimately require information in addition to size. liminosity. age. auc metallicity trends.," Understanding the origins of UCDs will ultimately require information in addition to size, luminosity, age, and metallicity trends."
 Two additional discriminators are their spatial auc velocity clistributions (which are both projections of an uncderling distribution)., Two additional discriminators are their spatial and velocity distributions (which are both projections of an underlying distribution).
 UCDs that are tidally stripped uuclei may be expected to reside on prefereutiallv radial orbits that result in a centrallyconcentrated number density distribution. a projected velocity dispersion profile that declines stronely with distance. aud a peaky. broacd-winged shape to their linc-ofsight velocity distribution (see Bassinoctal.1991:Thomasctal. 2008)).," UCDs that are tidally stripped nuclei may be expected to reside on preferentially radial orbits that result in a centrally-concentrated number density distribution, a projected velocity dispersion profile that declines strongly with distance, and a peaky, broad-winged shape to their line-of-sight velocity distribution (see \citealt{1994ApJ...431..634B,2003MNRAS.344..399B,2007MNRAS.380.1177B,2008MNRAS.385.2136G,2008MNRAS.389..102T}) )."
" Alternatively, UCDs that formed as extended star clusters uueht show an increasing dispersion profile aud a flat-topped velocitytion’."," Alternatively, UCDs that formed as extended star clusters might show an increasing dispersion profile and a flat-topped velocity."
 The density distribution of the λος UCDs will require further analysis that carefully considers selection effects. but their kinematics have heen analyzed in detail iu S|11 and add to the indications frou the color-magnitude diagram that the UCDs are distinct from the general GC population. aud not simply a tail of the GCs to larges," The density distribution of the M87 UCDs will require further analysis that carefully considers selection effects, but their kinematics have been analyzed in detail in S+11 and add to the indications from the color-magnitude diagram that the UCDs are distinct from the general GC population, and not simply a tail of the GCs to large."
"izes?! Briefly, over the distance range of RR.— 1035 kpce (Gwhere the data are available for both UCDs and compact GCs). the UCDs and. intermecdiate-size objects show a broader. flatter distribution of recession velocities than the GCs (considering only the blue GC subpopulation for a fair comparison)."," Briefly, over the distance range of $R \sim$ 10--35 kpc (where the data are available for both UCDs and compact GCs), the UCDs and intermediate-size objects show a broader, flatter distribution of recession velocities than the GCs (considering only the blue GC subpopulation for a fair comparison)."
 To illustrate this point further. we plot velocity vs. suze in Figure T. where the UCDs. intermecdiate-size objects. blue GCs are shown together.," To illustrate this point further, we plot velocity vs. size in Figure \ref{fig:kin}, where the UCDs, intermediate-size objects, blue GCs are shown together."
 The velocity distribution of compact objects appears to have the expected Caussian distribution. but the lager objects show a teudejicv fo. avoid he systemic velocity.," The velocity distribution of compact objects appears to have the expected Gaussian distribution, but the larger objects show a tendency to avoid the systemic velocity."
 This behavior ποσα» to set in for rjZ 5 pe. supporting our sugeestion from color cousideratious (Section ?7)) that many of the intermeciate-size objects should be identified as simall UCDs.," This behavior seems to set in for $r_{\rm h} \ga$ 5 pc, supporting our suggestion from color considerations (Section \ref{sec:cmd2}) ) that many of the intermediate-size objects should be identified as small UCDs."
" With 5 pe as the QC-UCD boundary. the velocity dispersions of the GCs aud the UCDs are 3102730 aand 500£90L. respectively,"," With 5 pc as the GC-UCD boundary, the velocity dispersions of the GCs and the UCDs are $340\pm30$ and $500\pm90$, respectively."
 A Ἱκομποροτον-Sunirnov test fuds that the velocity distributions are different at the confidence level., A Kolmogorov-Smirnov test finds that the velocity distributions are different at the confidence level.
 SJ11 discussed the Ms? UCD velocities iu more detail. includiug the treuds with cistance and huninosity.," S+11 discussed the M87 UCD velocities in more detail, including the trends with distance and luminosity."
 The UCD velocity dispersion profile remains constant. and the shape of the velocity distribution chauges im a complicated wax. ucither of which is uniquelv aud straightforwardly explained by cither of the formation scenarios under consideration.," The UCD velocity dispersion profile remains constant, and the shape of the velocity distribution changes in a complicated way, neither of which is uniquely and straightforwardly explained by either of the formation scenarios under consideration."
 It is possible that the blue UCDs comprise a iix of two different populations. with objects of dE unclei aud star-cluster origius beconune more dominant at the bright and faint cuds of the luminosity ranec. respectively;," It is possible that the blue UCDs comprise a mix of two different populations, with objects of dE nuclei and star-cluster origins becoming more dominant at the bright and faint ends of the luminosity range, respectively."
 Further theoretical work aud better statistics are needed to draw firmer conclusions about UCD origins from kinematics., Further theoretical work and better statistics are needed to draw firmer conclusions about UCD origins from kinematics.
 Up to this poimt. we lave focused on the AIST QGC/UCD system as a high-quality. well-characterizecl. rolmogencous dataset from a sinele environment.," Up to this point, we have focused on the M87 GC/UCD system as a high-quality, well-characterized, homogeneous dataset from a single environment."
 Now we seek to understand UCDs in a broader context. usine iterature data and results from other systems.," Now we seek to understand UCDs in a broader context, using literature data and results from other systems."
 We start * exanudnimeg basic trends m size and luuinositv. aud heu attempt to survey a broad range of their properties in order to converge on an iuteerated view of their ormational histories.," We start by examining basic trends in size and luminosity, and then attempt to survey a broad range of their properties in order to converge on an integrated view of their formational histories."
 To orient the discussion. we consider a basic though 1on-exliaustive set of four formation scenarios for UCDs.," To orient the discussion, we consider a basic though non-exhaustive set of four formation scenarios for UCDs."
" The first is that they are ""eijaut. CCS”. an extension of he normal CC population to very high masses. which wturally lead to large sizes owing to mass-depeudeucies of formation or internal evolution (e.¢.. Murray2009:Colesetal. 20103)."," The first is that they are “giant GCs”, an extension of the normal GC population to very high masses, which naturally lead to large sizes owing to mass-dependencies of formation or internal evolution (e.g., \citealt{2009ApJ...691..946M,2010MNRAS.408L..16G}) )."
 The second is that they are produced X normal star clusters that have collided (c.e.. Fellhaucr&Iroupa 2002)).," The second is that they are produced by normal star clusters that have collided (e.g., \citealt{2002MNRAS.330..642F}) )."
" We refer to these as merged GCs,", We refer to these as merged GCs.
 The third is that they pertain to an independent mode of diffuse star cluster formation that includes the lower huninosity ECs (ce... Drüusetal. 2011)).," The third is that they pertain to an independent mode of diffuse star cluster formation that includes the lower luminosity ECs (e.g., \citealt{2011A&A...529A.138B}) )."
 The fourth is that they are stripped galactic nuclei (e.g. Bekkietal.2001:Cocrdtetal. 2008)).," The fourth is that they are stripped galactic nuclei (e.g., \citealt{2001ApJ...552L.105B,2008MNRAS.385.2136G}) )."
" Two ""sinokiug guns” provide direct evidence that UCDs can form iu at least two cistinet wavs: W3 is likely a mereed GC (Marastouetal.20014:Felliauer&Ivoupa 2005).. and NGC 1516 UDI is likely a stripped uucleus (Norris&Roaunappan2011).."," Two “smoking guns"" provide direct evidence that UCDs can form in at least two distinct ways: W3 is likely a merged GC \citep{2004A&A...416..467M,2005MNRAS.359..223F}, and NGC 4546 UD1 is likely a stripped nucleus \citep{2011MNRAS.414..739N}."
 Below. as we review the sundry propertics of UCDs. we will conuneut at cach stage ou the compatibility of the data with these different formation scenarios. and then try to tie toecther the various lues of evidence into au integrated picture of CCD origins.," Below, as we review the sundry properties of UCDs, we will comment at each stage on the compatibility of the data with these different formation scenarios, and then try to tie together the various lines of evidence into an integrated picture of UCD origins."
 We asseiible from the literature a compilation of the sizes and Duuiuosities of hot stellar svstenis. from the largest galaxies to the smallest GCs.," We assemble from the literature a compilation of the sizes and luminosities of hot stellar systems, from the largest galaxies to the smallest GCs."
 We restrict the suuple to objects with distances confined either bv, We restrict the sample to objects with distances confirmed either by
simulations.,simulations.
 The low mass loss rate simulation has Alyw=5OAL.ve +., The low mass loss rate simulation has $\Mdot_{\rm W} = 5 \times 10^{-5} \Msol \pyr$ .
 In the high mass loss rate simulation Aly=d10*ALve eiving a bubble of half the temperature of the low mass loss rate simulation.," In the high mass loss rate simulation $\Mdot_{\rm W} = 1 \times 10^{-4} \Msol \pyr$, giving a bubble of half the temperature of the low mass loss rate simulation."
 Mass and energv are added to cells within r=3107cm at cach timestep., Mass and energy are added to cells within $r = 3 \times 10^{18} \cm$ at each timestep.
 The ambient. medium is assumed to be uniform and totally ionised. with a total number density of no=10cm," The ambient medium is assumed to be uniform and totally ionised, with a total number density of $n_{0}= 10 \pcc$."
 We shall only consider simulatedROSAP PSPC (Position Sensitive Proportional Counter) data., We shall only consider simulated PSPC (Position Sensitive Proportional Counter) data.
 Although the eas proportional counters spectral resolution of ALfleQ0.436ο7 (EWLIAL with E measured in keV) is low compared to a mission such asον wind-blown bubbles are soft X-ray sources and has more sensitivity than at low energies.," Although the gas proportional counter's spectral resolution of $\Delta E/E \approx 0.43 (E/0.93)^{-0.5}$ (FWHM, with E measured in $\keV$ ) is low compared to a mission such as, wind-blown bubbles are soft X-ray sources and has more sensitivity than at low energies."
 For a detailed: discussion. of theAT satellite: sec “TheROSAT user's handbook (Briel 1994)., For a detailed discussion of the satellite see `The user's handbook' (Briel 1994).
 As the low and high mass loss rate simulations only cüller in the density and temperature of the shocked. wind. we shall concentrate on describing the low mass loss rate simulation below.," As the low and high mass loss rate simulations only differ in the density and temperature of the shocked wind, we shall concentrate on describing the low mass loss rate simulation below."
 Section 3.4. describes how the results of the high mass loss rate bubble ciller (rom those given below., Section \ref{sec:res_cool_bub} describes how the results of the high mass loss rate bubble differ from those given below.
 There are three definable stages of bubble growth. seen in the simulation., There are three definable stages of bubble growth seen in the simulation.
 Phese are i) before the shell cools. ii) during shell cooling and collapse. and iii) self-similar growth after shell collapse with a thin cold (1— 10I) shell.," These are i) before the shell cools, ii) during shell cooling and collapse, and iii) self-similar growth after shell collapse with a thin cold $T \sim 10^{4}
\K$ ) shell."
 The 1-dimensional analytic solutions for the first and last of these stages are presented in detail in Castor (1975) and Weaver (1977)., The 1-dimensional analytic solutions for the first and last of these stages are presented in detail in Castor (1975) and Weaver (1977).
 Initially the swept up LSAL is shock-heatec to logZ'(Ix)= 6.0.," Initially the swept up ISM is shock-heated to $5.5 \ltsimm \log T {\rm
(K)} \ltsimm 6.0 $ ."
 The shell is thick (see Figs., The shell is thick (see Figs.
 1 and 2)). and a strong emitter of extreme Ultraviolet EUV) radiation and soft N-ravs. as can be seen from the 0.1-2.4keV luminosity (lig. 3)).," \ref{fig:dens_4t} and \ref{fig:temp_4t}) ), and a strong emitter of extreme Ultraviolet (EUV) radiation and soft X-rays, as can be seen from the $0.1$ $2.4 \keV$ luminosity (Fig. \ref{fig:lx}) )."
 The major coolant is radiation in the UV-IZUM rather than N-ravs. the UW-EUY luminosity being of order a magnitude greater than the soft. X-ray luminosity before shell collapse.," The major coolant is radiation in the UV-EUV rather than X-rays, the UV-EUV luminosity being of order a magnitude greater than the soft X-ray luminosity before shell collapse."
 The Iuminositv rises rapidly with time. as he bubble sweeps up and. heats more ISM.," The luminosity rises rapidly with time, as the bubble sweeps up and heats more ISM."
 Phe rate of increase of luminosity decreases after /zz4000vr. the X-ray uminosity peaking at Lx—1.9.107eresL| at £z6300vr as the shell begins to cool.," The rate of increase of luminosity decreases after $t \approx 4000 \yr$, the X-ray luminosity peaking at $L_{\rm X} = 1.9 \times 10^{36} \ergps$ at $t \approx 6300
\yr$ as the shell begins to cool."
 The UV-EUY luminosity peaks ater. at zzSLOOvr as the shell cools further out of the X-ray ud and becomes denser.," The UV-EUV luminosity peaks later, at $\approx 8100 \yr$ as the shell cools further out of the X-ray band and becomes denser."
 The peak UV-IZUVM. luminosity of 9.110?eres briefly exceeds the wind energy input.," The peak UV-EUV luminosity of $9.1 \times 10^{37}
\ergps$ briefly exceeds the wind energy input."
 After shell collapse N-rav. luminosities are approximately wo orders of magnitude below the UV-IEUV. luminosity. xh remaining essentially constant for the duration of the simulation.," After shell collapse X-ray luminosities are approximately two orders of magnitude below the UV-EUV luminosity, both remaining essentially constant for the duration of the simulation."
 Following shell collapse shell densities are. typically several hundred. to ai few thousand particles per cubic centimetre. with Jos10119.," Following shell collapse shell densities are typically several hundred to a few thousand particles per cubic centimetre, with $T \approx
10^{4} \K$."
 In the absence of heat conduction and evaporation olf the cool shell. we would. simiplistically expect. the bubble interior to have a uniform low density and high temperature.," In the absence of heat conduction and evaporation off the cool shell, we would simplistically expect the bubble interior to have a uniform low density and high temperature."
 In this case the shocked-wind material has the same pressure as predicted by Castor (1975) and Weaver (1977). »it the temperature determined. by the reverse shock (the ermination shock of the freely expanding wind).," In this case the shocked-wind material has the same pressure as predicted by Castor (1975) and Weaver (1977), but the temperature determined by the reverse shock (the termination shock of the freely expanding wind)."
 1n practice. the clensity rises in as the shell is approached. and the temperature drops. although not. to 10 extent expected for true conduction.," In practice, the density rises in as the shell is approached, and the temperature drops, although not to the extent expected for true conduction."
 This can be seen in both the 2D images of Figs. 1-, This can be seen in both the 2D images of Figs. \ref{fig:dens_4t}-
-2 and the radial profiles X Pie. 4.., \ref{fig:temp_4t} and the radial profiles of Fig. \ref{fig:weaver}.
 Phis is due to οτος and swirling motions along 1e shell-bubble interface mixing material from the dense shell into the hot bubble interior., This is due to eddies and swirling motions along the shell-bubble interface mixing material from the dense shell into the hot bubble interior.
 The outward. velocity in 1ο shocked-wind is higher than the velocity at which the shell expands into the LSAT (Fig. 4)).," The outward velocity in the shocked-wind is higher than the velocity at which the shell expands into the ISM (Fig. \ref{fig:weaver}) ),"
 and coupled with the ‘orrugations seen on the inside surface of the shell. shear motions arise between the faster bubble interior and. the gaell. leading to swirling motions along the interface and the introduction of cooler. denser material into the hot bubble.," and coupled with the corrugations seen on the inside surface of the shell, shear motions arise between the faster bubble interior and the shell, leading to swirling motions along the interface and the introduction of cooler, denser material into the hot bubble."
 In a perfectly spherically svnunetric bubble. the lack of a tangential velocity component between the faster-expanding bubble interior ancl the shell would prevent such stripping of material olf the shell.," In a perfectly spherically symmetric bubble, the lack of a tangential velocity component between the faster-expanding bubble interior and the shell would prevent such stripping of material off the shell."
 la our simulations. the bubble-shell interlace is corrugated by instabilities [roni early on in the simulation. presenting [aces not totally perpendicular to the Dow in the bubble interior. and Ieading to mixing.," In our simulations, the bubble-shell interface is corrugated by instabilities from early on in the simulation, presenting faces not totally perpendicular to the flow in the bubble interior, and leading to mixing."
 The instabilities of the shell seen in Figs., The instabilities of the shell seen in Figs.
 1 and 2. have important consequences as they lead to the introduction of cooler. denser material into the bubble interior. hence mocifving the X-ray emitting properties of the bubble.," \ref{fig:dens_4t} and \ref{fig:temp_4t} have important consequences as they lead to the introduction of cooler, denser material into the bubble interior, hence modifying the X-ray emitting properties of the bubble."
 The shell should. be stable against. Havleigh-Taylor instabilities. as it is constantly clecelcrating. sugecsting that the instabilities are Vishniac instabilities (Vishniac 1983).," The shell should be stable against Rayleigh-Taylor instabilities, as it is constantly decelerating, suggesting that the instabilities are Vishniac instabilities (Vishniac 1983)."
 The initial seed. perturbation is numerical artifact. arising when the forward shock first appears at the start of the simulation. and is due to the orthogonal nature of the computational erid and the finite size of the enerey injection reeion.," The initial seed perturbation is numerical artifact, arising when the forward shock first appears at the start of the simulation, and is due to the orthogonal nature of the computational grid and the finite size of the energy injection region."
towards the west.,towards the west.
 At MJD 52779.4. just 13 davs after ILUr43s X-ray flare. their separation was 166420 mas.," At MJD 52779.4, just 13 days after H1743's X-ray flare, their separation was $166 \pm 20$ mas."
 Later. on MJD 52782.4 and MJD 52786.4 the separations were 256+20 mas and 288+20 mas. respectively.," Later, on MJD 52782.4 and MJD 52786.4 the separations were $256 \pm 20$ mas and $288 \pm 20$ mas, respectively."
 The majority of the jet data considered in our analvsis are taken [rom Tables 1 and 3 of Corbeletal.(2005)., The majority of the jet data considered in our analysis are taken from Tables 1 and 3 of \citet{Corbel_2005}.
. These tables provide jet-source separation measurements for radio and X-ray observations which were conducted from 6 months onward following II1743's flare., These tables provide jet-source separation measurements for radio and X-ray observations which were conducted from 6 months onward following H1743's jet-launching flare.
 The X-ray data consist of three ~30 ksChandra X-ray observations in which both jets were detected., The X-ray data consist of three $\sim30$ ks X-ray observations in which both jets were detected.
 In radio. Corbeletal.(2005) report on five observations from the Australian Telescope Compact Array (ATCA).," In radio, \citet{Corbel_2005} report on five observations from the Australian Telescope Compact Array (ATCA)."
 The eastern jet was present in each image. but the western jet was detected only in the final observation.," The eastern jet was present in each image, but the western jet was detected only in the final observation."
" These X-ray and radio observations were carried out between MJD 52955 and MJD 53092. when the jet-source separations were in the range ~4""—7""."," These X-ray and radio observations were carried out between MJD 52955 and MJD 53092, when the jet-source separations were in the range $\sim4\arcsec-7\arcsec$."
 The substantially larger angular separations of the eastern jet indicate (hat it is approaching and (he western jet is receding., The substantially larger angular separations of the eastern jet indicate that it is approaching and the western jet is receding.
" In determining the spin of 11742. we analvze the full set of PCU-2 ""standard 27 data obtained during (he 2003 outburst. with the spectra binned into 170 hal[-day intervals."," In determining the spin of H1743, we analyze the full set of PCU-2 “standard 2” data obtained during the 2003 outburst, with the spectra binned into 170 half-day intervals."
 These spectra have been modeled in detail bv McClintockοἱal.(2009) and (2009).. ancl we use (he same data reduction procedures here.," These spectra have been modeled in detail by \citet{JEM_H1743} and \citet{Steiner_2009}, and we use the same data reduction procedures here."
 Briefly. all the data are corrected. background subtracted. and analvzed with the inclusion of a svstematic uncertainty (Jahodaetal.2006).," Briefly, all the data are dead-time corrected, background subtracted, and analyzed with the inclusion of a systematic uncertainty \citep{Jahoda_2006}."
. We stanclardize all detector calibrations to the Seward(1974). values for the Crab using a custom model which adjusts both the overall [τιν normalization and the spectral shape (see Steinerοἱal. 2010))., We standardize all detector calibrations to the \citet{Toor_Seward} values for the Crab using a custom model which adjusts both the overall flux normalization and the spectral shape (see \citealt{Steiner_lmcx3}) ).
" During the early weeks of the outburst evele. RATEss pointing was offset by 0.32"" [rom I11743."," During the early weeks of the outburst cycle, s pointing was offset by $0.32^{\circ}$ from H1743."
 We have corrected the fhixes to the full collimator transmission by assuming a (Giangular response will FWIIM = 1° (see Steinerοἱal. 2009)., We have corrected the fluxes to the full collimator transmission by assuming a triangular response with FWHM = $1^{\circ}$ (see \citealt{Steiner_2009}) ).
 Our jet model. which is based on one developed by Wangetal. (2003)... was [ist applied in describing gamma-rav-bursts.," Our jet model, which is based on one developed by \citet{WDL_2003}, , was first applied in describing gamma-ray-bursts."
 Here. we consider a pair of svyiumetric jets. each ejected with an initial kinetic energy. Ly and Lorentz [actor D.," Here, we consider a pair of symmetric jets, each ejected with an initial kinetic energy $E_0$ and Lorentz factor $\Gamma_0$."
 During their expansion. the jets decelerate as they sweepup gas in their paths.," During their expansion, the jets decelerate as they sweepup gas in their paths."
 Assuming acdiabatic expansion. the evolution of each jet is governed by:," Assuming adiabatic expansion, the evolution of each jet is governed by:"
removes the hieh frequency sigual from a stellar spectrum. which makes the region coutaiuing most ol the RV information more seusitive to the choice of OPD as the power spectrum clistribution becomes narrower due loss of high p component.,"removes the high frequency signal from a stellar spectrum, which makes the region containing most of the RV information more sensitive to the choice of OPD as the power spectrum distribution becomes narrower due loss of high $\rho$ component."
 h theory. a spectrograph with au idinitely high resolution would ye able to extract all the RV iuολατοι coutained in a stellar spectrum.," In theory, a spectrograph with an infinitely high resolution would be able to extract all the RV information contained in a stellar spectrum."
 However. in pratice. it is üunpossible to completely 'écove: the RV information with a specrograph with a finite spectral 'esolutiou whose spectral 'espolse [function drops at the high spaial [requeicy end.," However, in pratice, it is impossible to completely recover the RV information with a spectrograph with a finite spectral resolution whose spectral response function drops at the high spatial frequency end."
 Although te power spectrum of the derivaive of the stellar spectrum is shitec to the low frequency region where most o“the RV iufornlation is carried. the power spectri ds still xoad iu the spatial {'equeney. (p) doijalhi (see Fie. 1 .," Although the power spectrum of the derivative of the stellar spectrum is shifted to the low frequency region where most of the RV information is carried, the power spectrum is still broad in the spatial frequency $\rho$ ) domain (see Fig. \ref{fig:Rho_Power_PSF}) )."
 Thus. high. A can help to extract more RV information.," Thus, high $R$ can help to extract more RV information."
 Iu a waveleneth coverage from SOO ln to 1350 um. we calculate Qvalues lor stellar spectra with Vsiu£ of 0.2 .5 and 10 kil-S1 al ciffereit & (5.000 to 150.000 wi haste» of 5.000) i1 order to investigate tle depeudeuce oCQouk (Fig. τὴ).," In a wavelength coverage from 800 nm to 1350 nm, we calculate $Q$values for stellar spectra with $V \sin{i}$ of 0 ,2 ,5 and 10 $\rm{km\cdot s}^{-1}$ at different $R$ (5,000 to 150,000 with a step of 5,000) in order to investigate the dependence of $Q$ on $R$ (Fig. \ref{fig:Q_Res}) )."
 We find that more RV information (higler Q factor) ean be extracted as 2 ine‘eases., We find that more RV information (higher $Q$ factor) can be extracted as $R$ increases.
 (Q factors for DEDI and DE conve‘oe at hieh BR becase the spectral response tuuction is wide enough in the p domain to cover the region rk hin RV iformation. not affectecl yy the power spectrum shifting involved in DEDI.," $Q$ factors for DFDI and DE converge at high $R$ because the spectral response function is wide enough in the $\rho$ domain to cover the region rich in RV information, not affected by the power spectrum shifting involved in DFDI."
 In acdition. he Q [acto “ata given 2 increases as Tey drops from 31001. to 2100]x. which is largely due to stroneer moleclar absorption features 1ithe LY αμα J ας (see Fig. 3)).," In addition, the $Q$ factor at a given $R$ increases as $T_{\rm{eff}}$ drops from 3100K to 2400K, which is largely due to stronger molecular absorption features in the I, Y and J bands (see Fig. \ref{fig:Wav_Flux}) )."
 We divide & into three reglous. low resolution (5.000 to 20.000). ineitu resolutiou (20.000 o 20.000) and high resolution (50.000 to 120.000).," We divide $R$ into three regions, low resolution (5,000 to 20,000), medium resolution (20,000 to 50,000) and high resolution (50,000 to 150,000)."
 We use a power law to fit Q lor both DEDI aud DE as a function of 2., We use a power law to fit $Q$ for both DFDI and DE as a function of $R$ .
 The power tudices X of three regious for Teg2100 are yesented in Table 2.., The power indices $\chi$ of three regions for $T_{\rm{eff}}=2400K$ are presented in Table \ref{tab:PowerLawR}.
 At low R region. y remains roughly a constant for O lans cgslu’ <5 kines B ut τί drops for stars with Vsiu of 10 kines! indicating stellar absorption lines begiu to be 'esolved even at low £A.," At low $R$ region, $\chi$ remains roughly a constant for 0 $\rm{km\cdot s}^{-1}$ $\leq V \sin i\leq$ 5 $\rm{km\cdot s}^{-1}$ , but it drops for stars with $V \sin i$ of 10 $\rm{km\cdot s}^{-1}$ indicating stellar absorption lines begin to be resolved even at low $R$."
 At hieher A regions. x decreases as Vsiui? increases. a reduced value of x imiplies diminishing benelit brought by iicreasine 2. Stellar absorption lines are broadened by stellar 'Otatico. and they are resolve at a certain 2 beyond wlich iicreasiug 4? does not siguificautly ealn Doppler sensitivity.," At higher $R$ regions, $\chi$ decreases as $V \sin i$ increases, a reduced value of $\chi$ implies diminishing benefit brought by increasing $R$ Stellar absorption lines are broadened by stellar rotation, and they are resolved at a certain $R$ beyond which increasing $R$ does not significantly gain Doppler sensitivity."
 Over:ill. \ fo: DE is larger thar tha ol DEDI. especially for low and neciuiu. AU," Overall, $\chi$ for DE is larger than that of DFDI, especially for low and medium $R$."
 In other words. (2Dipl is ess sensitive to a clanee of A. aud the DEDI instrument can extract relatively more Doppler information at low or 1rectitu spectral resolution than the DE nethod.," In other words, $Q_{\rm{DFDI}}$ is less sensitive to a change of $R$, and the DFDI instrument can extract relatively more Doppler information at low or medium spectral resolution than the DE method."
 For example. for slow ‘otators (Vo sinz-2 kin-s1 al de low R region (R=5.000-20.000). Qprpix{ιο," For example, for slow rotators $V \sin i$ =2 $\rm{km\cdot s}^{-1}$ ) at the low $R$ region (R=5,000-20,000), $Q_{\rm{DFDI}}\propto R^{0.63}$."
" Doppler sensitivity οως is Inversely. proportional to two factors: Q and YN. ACCOLCling to Equation (6 (6)) ane (13)). where ,V, is the toal potou count collected by the CCD cletectOr."," Doppler sensitivity $\delta v_{rms}$ is inversely proportional to two factors: $Q$ and $\sqrt{N_{e^-}}$ according to Equation \ref{eq:overall_Doppler}) ) and \ref{eq:overall_Doppler_2d}) ), where $N_{e^-}$ is the total photon count collected by the CCD detector."
" JN,x(S/N)7 INgisar. where S/N is the average signal to noise ratio per pixel. aud Vyixel is total πια of pixels."," $N_{e^-}\propto(S/N)^2\cdot N_{\rm{pixel}}$ , where $S/N$ is the average signal to noise ratio per pixel, and $N_{\rm{pixel}}$ is total number of pixels."
" Note that ΑΝ,x Rifthe wavelenet coverage. S/N per pixel andthe resolution sampling are fixed."," Note that $N_{e^{-}}\propto R$ ifthe wavelength coverage, S/N per pixel andthe resolution sampling are fixed."
"Therefore. 6v,x10.630.5Hk113 for DEDL","Therefore, $\delta v_{rms}\propto R^{-0.63-0.5}=R^{-1.13}$ for DFDI."
 In comparison. NUprsx21.51 for DE given the same waveleneth coverage aix S/N per pixel.," In comparison, $\delta v_{rms}\propto R^{-1.57}$ for DE given the same wavelength coverage and S/N per pixel."
 The power law is, The power law is
the flat-spectrum emission in (νο N-3 and GRS 1915|105 appears to have approximately the same luminosity.,the flat-spectrum emission in Cyg X-3 and GRS 1915+105 appears to have approximately the same luminosity.
 GX 339-4 is a persistent. black hole candidate X- binary with similar radio properties to Cve δ-] (Llannikainen et al., GX 339-4 is a persistent black hole candidate X-ray binary with similar radio properties to Cyg X-1 (Hannikainen et al.
 1998: Fender et al., 1998; Fender et al.
 1999 and references therein)., 1999 and references therein).
 In particular the source displays at em wavelengths a flat spectrum with comparable luminosity to that of Cvg X-1., In particular the source displays at cm wavelengths a flat spectrum with comparable luminosity to that of Cyg X-1.
 We fully expect. therefore. that. sulliciently sensitive observations should also detect a flat spectrum through the (sub)mm regime [rom this source., We fully expect therefore that sufficiently sensitive observations should also detect a flat spectrum through the (sub)mm regime from this source.
 Additionally. as CX 4 is believed. to be a low mass X-ray binary with a less luminous companion star than in the (νο X-1 system. we may have more chance of detecting the Uat spectrum at near-infrared wavelengths.," Additionally, as GX 339-4 is believed to be a low mass X-ray binary with a less luminous companion star than in the Cyg X-1 system, we may have more chance of detecting the flat spectrum at near-infrared wavelengths."
 ‘Tables 1.3 and Figs 1 3 sumnmarise the observations of (νο NX-1 for 1997 Auge 4 ancl 1998 May. 11-20., Tables 1–3 and Figs 1 3 summarise the observations of Cyg X-1 for 1997 Aug 4 and 1998 May 11-20.
 The source is clearly clisplaving a flat spectrum through the radiomm regimes at both epochs., The source is clearly displaying a flat spectrum through the radio–mm regimes at both epochs.
 While radio emission from X-ray binaries is generally. assumed to be svnchrotron in origin (see e.g. Ljellming 1988: Hjellming Han 1995). in the case of Cvenus N-1 we do not have direct observational evidence for this.," While radio emission from X-ray binaries is generally assumed to be synchrotron in origin (see e.g. Hjellming 1988; Hjellming Han 1995), in the case of Cygnus X-1 we do not have direct observational evidence for this."
" Even the most rapid variability observed at 15 6112 does not require a brightness temperature in excess of 10"" Ix. and there is no direct measurement of linear polarisation."," Even the most rapid variability observed at 15 GHz does not require a brightness temperature in excess of $10^9$ K, and there is no direct measurement of linear polarisation."
 So. while some form of selt-absorbed. svnehrotron emission remains a possible origin for the [lat spectral component. other emissive mechanisms must also be considered.," So, while some form of self-absorbed synchrotron emission remains a possible origin for the flat spectral component, other emissive mechanisms must also be considered."
 The observed Luminosity of a Hat-spectrum source is directly woportional to the total bandwidth., The observed luminosity of a flat-spectrum source is directly proportional to the total bandwidth.
 In the case of (νο X-1l. the emmam flat spectral component corresponds to a radiative luminosity of 5210 erg 5 (2«1073 NC).," In the case of Cyg X-1, the cm–mm flat spectral component corresponds to a radiative luminosity of $\geq 2 \times 10^{31}$ erg $^{-1}$ $2 \times 10^{24}$ W)."
 Le he emission arises in an outflow in which non-racliative (c.g. acliabatic expansion) losses dominate (which seems likely to »' the case for relativistic jets from X-ray transients. see e.g. IIjellming Llan 1995) then even the integrated: radiative uminosityv is only a lower limit on the total power (i.e. it neglects c.g. electron acceleration ancl bulk kinetic energy) required to maintain the jet.," If the emission arises in an outflow in which non-radiative (e.g. adiabatic expansion) losses dominate (which seems likely to be the case for relativistic jets from X-ray transients, see e.g. Hjellming Han 1995) then even the integrated radiative luminosity is only a lower limit on the total power (i.e. it neglects e.g. electron acceleration and bulk kinetic energy) required to maintain the jet."
 Beyond the mm regime. in he infrared. thermal emission from the companion. stellar wind and accretion disc begin to dominate the spectrum. of the system (Fig.," Beyond the mm regime, in the infrared, thermal emission from the companion, stellar wind and accretion disc begin to dominate the spectrum of the system (Fig."
 3) and it may be very οΠοια to ever measure any high-frequency. limit to the flat spectral component emission., 3) and it may be very difficult to ever measure any high-frequency limit to the flat spectral component emission.
 Note that there is strong observational evidence that the [lat-spectrum oscillations observed. from GRS 1915|105 are dominated by adiabatic expansion losses. based. upon the similarity of the oscillation decay. rates at cm and infrared wavelengths (Fender et al.," Note that there is strong observational evidence that the flat-spectrum oscillations observed from GRS 1915+105 are dominated by adiabatic expansion losses, based upon the similarity of the oscillation decay rates at cm and infrared wavelengths (Fender et al."
 LOOT: Fender Pooley 1998)., 1997; Fender Pooley 1998).
" lt is easy to craw parallels between the Blat-spectrum radiomm emission. [rom (νο N-1. (anc also €Cvg X-3 and ο 1915]105: see above) and the “Hat-speetrun,’ extragalactic radio sources.", It is easy to draw parallels between the flat-spectrum radio–mm emission from Cyg X-1 (and also Cyg X-3 and GRS 1915+105; see above) and the `flat-spectrum' extragalactic radio sources.
 These systems are generally radio-loud AGN in which the flat-spectrum component corresponds to the ‘core’ or base of the jet., These systems are generally radio-loud AGN in which the flat-spectrum component corresponds to the `core' or base of the jet.
 s pointed out by Cotton et al. (, As pointed out by Cotton et al. (
"1980) it would appear to require a ""cosmic conspiracy of superposition of individual self-absorbed svnchrotron components in order to. produce a composite Hat spectrum.",1980) it would appear to require a `cosmic conspiracy' of superposition of individual self-absorbed synchrotron components in order to produce a composite flat spectrum.
 Alarscher Gear (1985) and. O'Dell οἱ al. (, Marscher Gear (1985) and O'Dell et al. (
1988) showed that vou can more comfortably reproduce ‘flat-spectrumy variability via shocks in conical jets.,1988) showed that you can more comfortably reproduce `flat-spectrum' variability via shocks in conical jets.
 A conical jet mocel for radio emission from X-ray binarics was presented by Hjellming Johnston (1988)., A conical jet model for radio emission from X-ray binaries was presented by Hjellming Johnston (1988).
 Ciiovanoni lIxazanas (1990) suggested. that energy transport by relativistic neutrons naturally explained. the combination of electron. spectrum. density and magnetic field profiles required to produce an observed Hat svachrotron spectrum.," Giovanoni Kazanas (1990) suggested that energy transport by relativistic neutrons naturally explained the combination of electron spectrum, density and magnetic field profiles required to produce an observed flat synchrotron spectrum."
 Alternatively. Wang ct al. (," Alternatively, Wang et al. ("
1997) have suggested that the Uat-spectrum emission. is optically thin [rom a. [lattened electron. energy distribution.,1997) have suggested that the flat-spectrum emission is optically thin from a flattened electron energy distribution.
 However. there are problems with the application of most. possibly all. of these models to the Lat radiomam(infrared) spectra observed (rom Cvg," However, there are problems with the application of most, possibly all, of these models to the flat radio–mm(–infrared) spectra observed from Cyg"
Studies of active objects at IR aud ταν wavelengths indicate that formation and ACN activity may be related (Fadda et al.,Studies of active objects at IR and X-ray wavelengths indicate that star-formation and AGN activity may be related (Fadda et al.
 2002)., 2002).
 The trigecr mechanis for both phenomena could be the iuteractiou or the mereine of eas-rich galaxies., The trigger mechanism for both phenomena could be the interaction or the merging of gas-rich galaxies.
 This ecuerates fast compression of the available eas iu the inner galactic regions. causing both the onset of a major starburst aud the fueling of a ceutral black hole raising the ACN activity.," This generates fast compression of the available gas in the inner galactic regions, causing both the onset of a major starburst and the fueling of a central black hole raising the AGN activity."
" However the concomitant AGN and starburst activity is expected to happen iu a hiel-deusity medium (Ny,2107?2h 2)) characterized by high dust extinction of the UV-optical flux aud strong photoclectric absorption of the soft X-rays (e.g. Fabian et al."," However the concomitant AGN and starburst activity is expected to happen in a high-density medium $N_H \geq 10^{23-24}\:$ ), characterized by high dust extinction of the UV-optical flux and strong photoelectric absorption of the soft X-rays (e.g. Fabian et al."
 1998)., 1998).
 Thus the study of these active pliases in ealaxies becomes very difficult: optical aud even mid-/far-IR spectroscopy may not be sufficient to diseutaugle starburst activity from ACN activity. which is actually best probed in the lard (E76 keV. in order to sample also the Fe Ίνα lie) X-ray euergv baud.," Thus the study of these active phases in galaxies becomes very difficult; optical and even mid-/far-IR spectroscopy may not be sufficient to disentangle starburst activity from AGN activity, which is actually best probed in the hard $E > 6\:$ keV, in order to sample also the Fe $\alpha$ line) X-ray energy band."
 To search for hidden ACNs iud to shed light ou the starburst-AGN conucction and its occurrence we have started a svstematie aud objective investigation m hard (E2 6keV) N-vavs ofgalecies., To search for hidden AGNs and to shed light on the starburst-AGN connection and its occurrence we have started a systematic and objective investigation in hard $E >6\:$ keV) X-rays of.
" The sample consists of 28 ealaxies selected from theQuasars (see lttp:/Aisdipac.caltecl.edu/] as having fü)ja,>50 Jy or fosjan>10 Jy.", The sample consists of 28 galaxies selected from the (see http://irsa.ipac.caltech.edu/) as having $f_{60\mum} > 50\:$ Jy or $f_{25\mum} > 10\:$ Jy.
" We stress here that no other selection criteria (6g. established presence of au ACN, Iuninosities. IR colours. ete.)"," We stress here that no other selection criteria (e.g. established presence of an AGN, luminosities, IR colours, etc.)"
 have been applied to the sample definition., have been applied to the sample definition.
"i) The geometry. (shape) of the svstem (i.e. all distances are proportional {ο a scale factor): i) The sublimation temperature of the material: ii) The spectral shape of the input radiation. AFA/F: iv) The shape of the cust absorption and scattering opacilies. ayΑλ and e3/65,. respectively. where Ap is the fiducial wavelenght: v) The dust scattering phase function(SPF): vi) The overall optical depth at the fiducial wavelength.","i) The geometry (shape) of the system (i.e., all distances are proportional to a scale factor); ii) The sublimation temperature of the material; iii) The spectral shape of the input radiation, $\lambda F_\lambda/F$; iv) The shape of the dust absorption and scattering opacities, $\kappa_\lambda/\kappa_{\lambda_0}$ and $\sigma_\lambda/\sigma_{\lambda_0}$ , respectively, where $\lambda_0$ is the fiducial wavelenght; v) The dust scattering phase function; vi) The overall optical depth at the fiducial wavelength."
 The above list can be regarded as a set ofrequirements., The above list can be regarded as a set of.
 Two models that have dillerent plysical parameters. but meet the above requirements exactly. are equivalent and have (he same SED.," Two models that have different physical parameters, but meet the above requirements exactly, are equivalent and have the same SED."
 For example. the physical dimensions of the dusty region can be freely changed (sav by changing the stellar huminositv. which increases the dust condensation radius) with no effect on the SED. as long as the overall optical depth and shape factors do nol change.," For example, the physical dimensions of the dusty region can be freely changed (say by changing the stellar luminosity, which increases the dust condensation radius) with no effect on the SED, as long as the overall optical depth and shape factors do not change."
 If an invariance requirement is violated. however. (hen the models are no longer equivalent and the SEDs are expected to be different.," If an invariance requirement is violated, however, then the models are no longer equivalent and the SEDs are expected to be different."
 For instance. if one were to change the grain size. the shape of (he opacity ancl the SPF would change. violating requirements (iil) and (v).," For instance, if one were to change the grain size, the shape of the opacity and the SPF would change, violating requirements (iii) and (v)."
 For a given class of astronomical objects (e.g... AGB stus). most of the quantities above are likely to be similar (geometry ancl grain composition). so the optical cleptl becomes the single most important parameter (hat controls the SED.," For a given class of astronomical objects (e.g., AGB stars), most of the quantities above are likely to be similar (geometry and grain composition), so the optical depth becomes the single most important parameter that controls the SED."
 It is reasonable to expect that if an invariance requirement is weakly violated. the SED will still be approximately the same.," It is reasonable to expect that if an invariance requirement is weakly violated, the SED will still be approximately the same."
 This was noted by IE97 who pointed out that if the erains are very sniall (about a tenth of wavelength of the peak of (he source spectrum) then the shape of the opacities are very. similar and the grain size is irrelevant for the problem., This was noted by IE97 who pointed out that if the grains are very small (about a tenth of wavelength of the peak of the source spectrum) then the shape of the opacities are very similar and the grain size is irrelevant for the problem.
 For a 3000 Ix source. for example. the upper limit for the grain size is about 0.0570.," For a 3000 K source, for example, the upper limit for the grain size is about $0.05 \mu \rm m$."
 Grains larger (han Chis upper limit will. according to IE9T. significantly alter the results.," Grains larger than this upper limit will, according to IE97, significantly alter the results."
 In the case of optically thin envelopes. (his condition on the masini grain size can be further relaxed. because the star completely dominates the optical SED while the grains produce the IR. SED.," In the case of optically thin envelopes, this condition on the maximum grain size can be further relaxed, because the star completely dominates the optical SED while the grains produce the IR SED."
 For the IR SED to remain similar between models with different erain racii. itis evident that both the shape of the IR emissivity and the reprocessedIuminosity," For the IR SED to remain similar between models with different grain radii, it is evident that both the shape of the IR emissivity and the reprocessedluminosity"
where C is a normalization constant that should be determined [rom matching this solution with that in the regionp.,where C is a normalization constant that should be determined from matching this solution with that in the region.
.. However. the function e depends on the solution itself.," However, the function v depends on the solution itself."
 Fortunately. (this quantity can be caleulated prior to determining and therefore. this solution may. be written in a closed. form.," Fortunately, this quantity can be calculated prior to determining and therefore, this solution may be written in a closed form."
 To illustrate this. let us consider a particularly simple case ofpyay.. and we will turn to (he general case alterwiurds.," To illustrate this, let us consider a particularly simple case of, and we will turn to the general case afterwards."
 Clearly. means u4.," Clearly, means."
. Evidently. we max replace C in eq.(34)) bv so that for vip) we have (C1 p/puas)).," Evidently, we may replace G in \ref{v:def}) ) by so that for (p) we have ( )."
 Thus. from eq.(35)) we obtain the following shape of the cut-off near In the vest of the vr. p--domain where yf) and p is not close (opas. We Way assume (hat the CR cliffision coefficient is close (ο its Dohm value.," Thus, from \ref{G0:sol}) ) we obtain the following shape of the cut-off near In the rest of the x,p -domain where (p) and p is not close to, we may assume that the CR diffusion coefficient is close to its Bohm value."
 Indeed. in contrast to the phase space region ury(p) al each given vr.p there are waves generated along the entire characteristic of eq.(14)) passing through this point of the phase space and occupying an extended reeion of the CR precursor. Figure 2..," Indeed, in contrast to the phase space region (p) at each given x,p there are waves generated along the entire characteristic of \ref{wke2}) ) passing through this point of the phase space and occupying an extended region of the CR precursor, Figure \ref{fig:ph:plane}."
 We may use then the asvaptotic hieh Mach number solution found in (Malkov.1997) llere is numerically small (ivpically 1/6)) and Chis solution without 3--term manifests the balance between the diffusion ancl convection terms on the Ilis.," We may use then the asymptotic high Mach number solution found in \citep{m97a}
 Here is numerically small (typically 1/6 ) and this solution without -term manifests the balance between the diffusion and convection terms on the l.h.s."
 of eq.(13)) which is more accurate approximation far upstream where the flow modification (r.h.s.), of \ref{dc2}) ) which is more accurate approximation far upstream where the flow modification (r.h.s.)
 is weak., is weak.
 The flow profile depends on the form of &(p) and for p in the internal part of the shock (ransition uc) behaves linearly with 2.2 Adopting this solution to (the regionrj. we may write so that for e we have," The flow profile depends on the form of (p) and for p in the internal part of the shock transition u(x) behaves linearly with x. Adopting this solution to the region, we may write so that for v we have"
is also possible that there is more than one episode of AGNjet activity in the course of a merger.,is also possible that there is more than one episode of AGN/jet activity in the course of a merger.
 For example. ⋜⋯↕↓↕∠⇂↕∖⇁↕∠⇂⇂⇂⋜↧⇂⊳∖∙∖⇁⊳∖↿⋖⋅⊔↓⊔↓⋜↧∙∖⇁⊔⊔∠⇂∢⊾↓⋅⋏∙≟∪⋜↧↓≻⋜↧↓⋅∣↕≼⇍⇂⇂↓⋜↧↓⋅⇂∙∖⇁↓≻∪∖∖⊽∢⋅↓⋅⇂⋅ ⊳− phase of AGN/jet activity as the nuclei coalesce. close to the peak of starburst activity (the coalescence phase above). but the AGN/jet activity may also be triggered (or re-trigecrecl) earlier or later in the merger sequence. depending on the details of the radial eas Lows in the merger.," For example, an individual system may undergo a particularly powerful phase of AGN/jet activity as the nuclei coalesce, close to the peak of starburst activity (the coalescence phase above), but the AGN/jet activity may also be triggered (or re-triggered) earlier or later in the merger sequence, depending on the details of the radial gas flows in the merger."
" Hndeed. as noted above. six of the starburst radio galaxies in our sample show evidence for re-triggered radio source activity in the form of high surface brightness inner radio structures anc more cdilluse and extended outer radio structures (36213.1. 3€218. Con AX.3€236.of 3€293.""n PIXS1345|12)."," Indeed, as noted above, six of the starburst radio galaxies in our sample show evidence for re-triggered radio source activity in the form of high surface brightness inner radio structures and more diffuse and extended outer radio structures (3C213.1, 3C218, Cen A, 3C236, 3C293, PKS1345+12)."
 Moreover. in the particular case 3CM O'Deaetal.(2001). have argued. for multiple phases of jet activity. based. on both the double-double morphology of its radio source. ancl the evidence for two major epochs of star formation in its host galaxy.," Moreover, in the particular case of 3C236, \citet{odea01} have argued for multiple phases of jet activity, based on both the double-double morphology of its radio source, and the evidence for two major epochs of star formation in its host galaxy."
 We emphasise that not all of the starburst radio galaxies can be readily accommocdated: within the merger. scheme outlined. above: there are. some prominent misfits., We emphasise that not all of the starburst radio galaxies can be readily accommodated within the merger scheme outlined above; there are some prominent misfits.
" Most notably. the central cluster galaxies 3€C218 and PINS0620-52 do not show clear morphological signs of major mergers. and 3€218 is also unusual in having relatively voung VSP (£4, 0.05 Gyr) but low emission line ancl far-LR luminosities (the age the YSP in PINSOQC20-52 is not well-determined)."," Most notably, the central cluster galaxies 3C218 and PKS0620-52 do not show clear morphological signs of major mergers, and 3C218 is also unusual in having relatively young YSP $t_{ysp}\sim$ 0.05 Gyr) but low emission line and far-IR luminosities (the age of the YSP in PKS0620-52 is not well-determined)."
 PINSO023-26a may also represent an ambiguous case in the sense it lies at the heart of a rich. cluster of galaxies. and has a peculiar amorphous outer envelope that is difficult to classify in terms of the merger sequence (although a merger cannot be entirely ruled out).," PKS0023-26 may also represent an ambiguous case in the sense that it lies at the heart of a rich cluster of galaxies, and has a peculiar amorphous outer envelope that is difficult to classify in terms of the merger sequence (although a merger cannot be entirely ruled out)."
 In the case of the one system in our sample that shows evidence For large-scale. and relatively settled. gaseous disk NGC612 it djs not clear whether the AGN/jet activity has been triggered by interactions with massive galaxies on a large scale in the wider galaxy group (as may be evidenced bv the LIE observations). or bv recent. mergers/interactions with closer companion galaxies (seeI5montsctal.2008a).," In the case of the one system in our sample that shows evidence for large-scale, and relatively settled, gaseous disk — NGC612 — it is not clear whether the AGN/jet activity has been triggered by interactions with massive galaxies on a large scale in the wider galaxy group (as may be evidenced by the HI observations), or by recent mergers/interactions with closer companion galaxies \citep[see][]{emonts08a}."
. The latter seems. more likely given that δές019 shows an optical shell structure. along with YSP that are much vounger (Fui0.1 Gyr) than the estimated. time since its closest. approach to the the massive companion galaxy NGCG619 l Gyr)," The latter seems more likely given that NGC612 shows an optical shell structure, along with YSP that are much younger $t_{ysp} < 0.1$ Gyr) than the estimated time since its closest approach to the the massive companion galaxy NGC619 $\ge 1$ Gyr)."
 Clearly. the possibility. of multiple interactions and mergers in galaxy groups can complicate the interpretation of the galaxy morphologies and YSP ages in terms of a simple merger sequence.," Clearly, the possibility of multiple interactions and mergers in galaxy groups can complicate the interpretation of the galaxy morphologies and YSP ages in terms of a simple merger sequence."
 Finally we note that. while the properties of the most luminous starburst radio galaxies in particular. the ULLIBCG-like coalescence svstenis are consistent. with triggering in major (similar mass) gas-rich mergers. i60 is dillicult to rule out minor mergers as the trigger for some of the other svstems.," Finally we note that, while the properties of the most luminous starburst radio galaxies – in particular, the ULIRG-like coalescence systems – are consistent with triggering in major (similar mass) gas-rich mergers, it is difficult to rule out minor mergers as the trigger for some of the other systems."
 Indeed. it has been argued that the large-scale. features of Centaurus A are consistent with a minor (1:10) merger between the host radio galaxy and a smaller clisk galaxy (Malin.Quinn&Graham1983)..," Indeed, it has been argued that the large-scale features of Centaurus A are consistent with a minor (1:10) merger between the host radio galaxy and a smaller disk galaxy \citep{malin83a}."
 As noted in the previous section. some of the starburst radio galaxies that are situated close to the centres of rich clusters of galaxies are not readily. accommocdated. in. the nmiergor sequence.," As noted in the previous section, some of the starburst radio galaxies that are situated close to the centres of rich clusters of galaxies are not readily accommodated in the merger sequence."
 For such objects cooling associated with the hot N-ray. emitting gas may provide an alternative rigecrine mechanism for both the visible star formation and the AGN/jet activity., For such objects cooling associated with the hot X-ray emitting gas may provide an alternative triggering mechanism for both the visible star formation and the AGN/jet activity.
 This mechanism is supported by he irregular morphologies anc kinematics of the emission inc σας in central cluster galaxies hosting radio sources Tadhunter.Fosbury&Quinn1989:Daum.HeckmanvanDreugel1992).," This mechanism is supported by the irregular morphologies and kinematics of the emission line gas in central cluster galaxies hosting radio sources \citep{tadhunter89,baum92}."
. Moreover. while the high velocity clispersions of the galaxies in the centres of massive galaxy clusters can junper gas acecretion via major. gas-rich mergers. there is no such problem for the accretion of warm/hot gas via cooling lows. since the cooling eas will naturally fall towards the centre of the the cluster potential well.," Moreover, while the high velocity dispersions of the galaxies in the centres of massive galaxy clusters can hamper gas accretion via major, gas-rich mergers, there is no such problem for the accretion of warm/hot gas via cooling flows, since the cooling gas will naturally fall towards the centre of the the cluster potential well."
 Although star formation has been discussed. as a rotential sink for the cooling eas. until recentlv the apparently large dilferences between the estimated hot. gas cooling rates and the star formation rates in the central cluster galaxies suggested that only a small fraction of the the cooling gas ends up forming stars (e.g.MeNamara&OConnell1989) ," Although star formation has been discussed as a potential sink for the cooling gas, until recently the apparently large differences between the estimated hot gas cooling rates and the star formation rates in the central cluster galaxies suggested that only a small fraction of the the cooling gas ends up forming stars \citep[e.g.][]{mcnamara89}."
- However. spectroscopic observations with the new generation of X-ray. satellites have [ed to a major downward revision in estimates of the hot gas cooling rates. so that they are now much closer to the star formation rates.," However, spectroscopic observations with the new generation of X-ray satellites have led to a major downward revision in estimates of the hot gas cooling rates, so that they are now much closer to the star formation rates."
 Therefore. it is plausible that a significant fraction of the cooling gas does in [act form stars (seeRallertyetal.2006).," Therefore, it is plausible that a significant fraction of the cooling gas does in fact form stars \citep[see][]{rafferty06}."
. In the cases of the three cooling How caneliclates in our sample. based on their infrared. luminosities ancl the relation of Ixennicutt.(1998) star formation rates are 5. dand 25 M. ve+ for PIKS0620-52. 3C218(Hvdra A) and WSO023-36 respectively.," In the cases of the three cooling flow candidates in our sample, based on their infrared luminosities and the relation of \citet{kennicutt98}, the star formation rates are 5, 4 and 25 $_{\odot}$ $^{-1}$ for PKS0620-52, 3C218(Hydra A) and PKS0023-36 respectively."
 OL these three. only 3€218 has published high. quality X-ray observations. and its hot gas cooling rate of 10ΕΕ M. + (Rallertyetal.2006) proves o be within a [actor of 4 of the star formation rate.," Of these three, only 3C218 has published high quality X-ray observations, and its hot gas cooling rate of $16\pm4$ $_{\odot}$ $^{-1}$ \citep{rafferty06} proves to be within a factor of 4 of the star formation rate."
 Given the agreement between its star. formation and 100 gas cooling rates. 3€218 is one of the best. candidates or an object in which the activity has been triggered by the wirm/cool gas condensing out. of a cooling How.," Given the agreement between its star formation and hot gas cooling rates, 3C218 is one of the best candidates for an object in which the activity has been triggered by the warm/cool gas condensing out of a cooling flow."
 However. even in the case of 3€218 we cannot entirely rule out the idea hat the activity has been triggered in a galaxy merger or interaction. since it is possible for central cluster galaxies to undergo mergers that could. potentially. form star forming eascous disks similar to that associated with the dust [ane in the central regions of 3€218 (RamosAlmeidaetal. 2010).," However, even in the case of 3C218 we cannot entirely rule out the idea that the activity has been triggered in a galaxy merger or interaction, since it is possible for central cluster galaxies to undergo mergers that could, potentially, form star forming gaseous disks similar to that associated with the dust lane in the central regions of 3C218 \citep{ramos10}."
. Although such disks (and the merging galaxies that produced them) may represent only a small fraction of the total masses of the cD galaxies. the associated eas infall rates may be sullicient to fuel the AGN/jet activity in the nuclei on the requisite timescales.," Although such disks (and the merging galaxies that produced them) may represent only a small fraction of the total masses of the cD galaxies, the associated gas infall rates may be sufficient to fuel the AGN/jet activity in the nuclei on the requisite timescales."
 Note that the large-scale morphological signatures of galaxy. mergers (c.g. tidal tails. fans. shells) are more dillicult to detect against the light of the massive stellar haloes of the central cluster galaxies than they are in lower mass elliptical galaxies: the tidal features are also likely to be erased on a relatively short timescale by ongoing tidal interactions between all the galaxies in the dense central regions of the galaxy. clusters.," Note that the large-scale morphological signatures of galaxy mergers (e.g. tidal tails, fans, shells) are more difficult to detect against the light of the massive stellar haloes of the central cluster galaxies than they are in lower mass elliptical galaxies; the tidal features are also likely to be erased on a relatively short timescale by ongoing tidal interactions between all the galaxies in the dense central regions of the galaxy clusters."
eraius for which Θωρμα)21 and (Q1pan.230IN))zz0.05 (Draine&Lee198[).,"grains for which $Q_{abs}(Ly\alpha) \approx 1$ and $\left< Q(1\ \mu{\rm m},230\ {\rm K})\right> \approx 0.05$ \citep{dra84}."
. Then. We then fiud where we have taken 5»;=(1—\Jny (since the fraction of atoms in excited. bound states is negligible).," Then, We then find where we have taken $n_i = (1-\chi)n_H$ (since the fraction of atoms in excited, bound states is negligible)."
 To obtain the numerical value given above we used the recombination coefficient in he density bounded case with T=10! 1 (Osterbrock1980)., To obtain the numerical value given above we used the recombination coefficient in the density bounded case with $T = 10^4$ K \citep{ost89}.
. Assuming Lyian-a radiation is orimarily removed through absorption by dust (in which case erl(ie)/iw is replaced by2472/z).{99f/ hen ο unless X exceeds 2.1%10.7. which is unlikely (Osterbrock1989).," Assuming $\alpha$ radiation is primarily removed through absorption by dust (in which case ${\rm erf}(w)/w$ is replaced by$2\sqrt{2/\pi}$ ), then $S < S_{crit}$ unless $\chi$ exceeds $2.4 \times 10^{-3}$, which is unlikely \citep{ost89}."
. It should also be ioted that other sources. such as contiuuuim radiation. are likely importantin heating the dust. hereby reduciug further the implied value of γα (aud therefore S$).," It should also be noted that other sources, such as continuum radiation, are likely importantin heating the dust, thereby reducing further the implied value of $n_{Ly\alpha}$ (and therefore $S$ )."
" Finally. the 1jan grain size assiiued hereis probably an overestimate iiplyiug that (7,4 is also overestimatect."," Finally, the $1\ \mu$ m grain size assumed hereis probably an overestimate implying that $n_{Ly\alpha}$ is also overestimated."
 Evideutly. the 25 state is overpopulated relative to 2p. thus the fine structure trausitious will proceed [rom 254;5 10 2pa4;5 via absorption aud to 2p4;5 via stimulated. emission.," Evidently, the $2s$ state is overpopulated relative to $2p$ , thus the fine structure transitions will proceed from $2s_{1/2}$ to $2p_{3/2}$ via absorption and to $2p_{1/2}$ via stimulated emission."
 Although the 2p populations are probably negligible. they will be included in the radiative trausfer calculation.," Although the $2p$ populations are probably negligible, they will be included in the radiative transfer calculation."
 The distribution of 2p states between 2p4;5 aud 2p4;5 maydeviate somewhat from the statistical weights (1/3 and 2/3. respectively). in part because the separate collisional rates (rom 2s are uot proportional to the statistical weights.," The distribution of $2p$ states between $2p_{1/2}$ and $2p_{3/2}$ maydeviate somewhat from the statistical weights $1/3$ and $2/3$, respectively), in part because the separate collisional rates from $2s$ are not proportional to the statistical weights."
 Since the 2p population is most likely negligible. a detailed calculation of the distribution between 2p4;» and 2p4;5 states will not be carried outhere.," Since the $2p$ population is most likely negligible, a detailed calculation of the distribution between $2p_{1/2}$ and $2p_{3/2}$ states will not be carried outhere."
" Rather. the fractional populations of 2p4;5 aud 2p4;5 will be parameterized as 9,/3 and 2955/3. respectively."," Rather, the fractional populations of $2p_{1/2}$ and $2p_{3/2}$ will be parameterized as $\beta_a/3$ and $2\beta_b
/ 3$, respectively."
 If;=ον1. then these states are populated according to their statistical weights.," If $\beta_a = \beta_b = 1$, then these states are populated according to their statistical weights."
" There is the obvious constraint that 2,/3+295/53=1.", There is the obvious constraint that $\beta_a/3 + 2\beta_b / 3 = 1$.
" The fine structure transitions are allowed electric dipole trausitions and the correspouding rates may be computed in a stralghitforward manuuer (Bethe&Salpeter1957).. giving Ay,=1.597x109 | (25,,5-2p4,5) and Ay—ST8x10.* 1 (2p4,5-25,,5)."," The fine structure transitions are allowed electric dipole transitions and the corresponding rates may be computed in a straightforward manner \citep{bet57}, giving $A_a = 1.597 \times 10^{-9}$ $^{-1}$ $2s_{1/2}$ $2p_{1/2}$ ) and $A_b = 8.78 \times 10^{-7}$ $^{-1}$ $2p_{3/2}$ $2s_{1/2}$ )."
 The absorption coellicient (valid for either transition) is where g is the degeneracy of the final state (2 for 2p4;5. E for 2pa;s: gas= 2).," The absorption coefficient (valid for either transition) is where $g$ is the degeneracy of the final state (2 for $2p_{1/2}$ , 4 for $2p_{3/2}$ ; $g_{2s} = 2$ )."
" The —sien corresponds to trausitions to ρω the + sien to trausitions to 2p4;5. and 3 is either 6, or 5). respectively."," The $-$sign corresponds to transitions to $2p_{1/2}$ ; the $+$ sign to transitions to $2p_{3/2}$ , and $\beta$ is either $\beta_a$ or $\beta_b$, respectively."
 Either final state quickly decays to the grouud state via Lyiman-a with rate σι., Either final state quickly decays to the ground state via $\alpha$ with rate $A_{21}$ .
 The, The
"We have simulated for a wide range of parameters as Ey,=10-10?! eres. ΑΝ=10- 0. P= 100-1000. and R=100-1055 em.","We have simulated for a wide range of parameters as $E_{\rm tot}=10^{48}$ $10^{54}$ ergs, $N=10$ $1000$, $\Gamma=100$ $1000$ , and $R=10^{13}$ $10^{15}$ cm."
" Of course. larger £4, and smaller 2 are favorable for neutrino production. and as DermerandAtovan(2003) showed. very luminous bursts are required to detect neutrinos on the Earth."," Of course, larger $E_{\rm tot}$ and smaller $R$ are favorable for neutrino production, and as \citet{der03} showed, very luminous bursts are required to detect neutrinos on the Earth."
 Therefore. we show only oue representative example in this Letter.," Therefore, we show only one representative example in this Letter."
" The parameter values are Zi=10?! eres. N=1000 (Ej,=10°! eres). P=100. and R=LOM em."," The parameter values are $E_{\rm tot}=10^{54}$ ergs, $N=1000$ $E_{\rm sh}=10^{51}$ ergs), $\Gamma=100$, and $R=10^{13}$ cm."
" The corresponding variability timescale A4,7R/T?~30 ms. which is not so far [rom the typical observed timescale mm>0.1/(1+2) s. since the allowed region of the GRB parameters is wide. there may be both optically (hin ancl thick sources to Thomson scattering (MészárosandRees2000)."," The corresponding variability timescale $t_{\rm var} \simeq R/\Gamma^2 \sim 30$ ms, which is not so far from the typical observed timescale $\gtrsim 0.1/(1+z)$ s. Since the allowed region of the GRB parameters is wide, there may be both optically thin and thick sources to Thomson scattering \citep{mes00}."
 Our example would imply (hat. fe is close to the photosphere., Our example would imply that $R$ is close to the photosphere.
 When it is assumed (hat the energy density of protons is the same as the photon energy density and (he average proton energy is mildly relativistie (Eο Πρ). the proton munber density in the comoving [rame is obtained as 2Fy/(20rRmc).," When it is assumed that the energy density of protons is the same as the photon energy density and the average proton energy is mildly relativistic $\lesssim 5 m_{\rm p} c^2$ ), the proton number density in the comoving frame is obtained as $\gtrsim E_{\rm sh}/(20 \pi R^3 m_{\rm p} c^2)$."
 This means that the optical depth for the Thomson scattering is ~1 for our parameter set., This means that the optical depth for the Thomson scattering is $\sim 1$ for our parameter set.
 Photon scatterings do not sufficiently affect (he eamama-ray spectrum for this marginal optical depth., Photon scatterings do not sufficiently affect the gamma-ray spectrum for this marginal optical depth.
 From our simulation. we obtain spectra of created mesons as is shown in Figure 2.," From our simulation, we obtain spectra of created mesons as is shown in Figure 2."
 One-half of neutral kaons are Nj. while the rest. are NS.," One-half of neutral kaons are $K^0_{\rm L}$, while the rest are $K^0_{\rm S}$."
 Since the cross sections of kaon production are smaller Chan (hose of pion production. the munber of kaons is niuch less (han pions.," Since the cross sections of kaon production are smaller than those of pion production, the number of kaons is much less than pions."
 Llowever. the highest energv charged mesons will cool down belore (μεν decay. into neutrinos.," However, the highest energy charged mesons will cool down before they decay into neutrinos."
 Our results lor other parameter setsm):agree wilh (he condition of ultra high energy cosmic- production obtained by Asano(20ἱ R>10Mby/10tere)? cem., Our results for other parameter sets agree with the condition of ultra high energy cosmic-ray production obtained by \citet{asa05}: $R \gtrsim 10^{14} (E_{\rm sh}/10^{51} {\rm erg})^{1/2}$ cm.
 In the case of Figure 2. also protons above 1012 eV cool down belore they escape from the shell.," In the case of Figure 2, also protons above $10^{15}$ eV cool down before they escape from the shell."
 We follow the behavior of pions aud kaons until they decay into positrons (electrons) ancl neuirinos using the same method as in Asano (2005).., We follow the behavior of pions and kaons until they decay into positrons (electrons) and neutrinos using the same method as in \citet{asa05}. .
 Svncehrotron and inverse Compton emissions are taken into account., Synchrotron and inverse Compton emissions are taken into account.
" Charged kaons have six decay modes: A/——jipy, )). ππ (214).-xmm (654)) ποm (D). πμv, (35)). and &""x"" (250). while Nj Fivewill decavinto x€£ (394)).*p v,(TA)). wean! )). and xz"" (134))."," Charged kaons have six decay modes; $K^+ \to \mu^+ \nu_\mu$ ), $\pi^+ \pi^0$ ), $\pi^+ \pi^+ \pi^-$ ), $\pi^0 e^+ \nu_e$ ), $\pi^0 \mu^+ \nu_\mu$ ), and $\pi^+ \pi^0 \pi^0$ ), while $K^0_{\rm L}$ will decay into $\pi^+ e^- \bar{\nu_e}$ ), $\pi^+ \mu^- \bar{\nu_\mu}$ ), $\pi^0 \pi^0 \pi^0$ ), and $\pi^+ \pi^- \pi^0$ )."
 In 3. total neutrino spectra emitted from this example are shown.," In Figure 3, total neutrino spectra emitted from this example are shown."
" Although there are fewer kaons than pions. the highest οποίον neutrinos originate [rom kaons around e,LOY eV. since. very high. [lux is. required. (o detect neutrinos. Irom. GRBs by a Ikm> neutrino. detector suchas IeeCube (DermerandAtovan2003).. we consider an optimistic case: a GRB occurs al 30 Mpc. and the detection efficiency. of upward-going neutrinos wilh enerev €, is assumedto be LOtHe,ο for e,>10!! eV. although it may be difficult to"," Although there are fewer kaons than pions, the highest energy neutrinos originate from kaons around $\epsilon_\nu \sim 10^{18}$ eV. Since very high flux is required to detect neutrinos from GRBs by a $1 {\rm km}^2$ neutrino detector suchas IceCube \citep{der03}, we consider an optimistic case: a GRB occurs at 30 Mpc, and the detection efficiency of upward-going neutrinos with energy $\epsilon_\nu$ is assumedto be $10^{-4} (\epsilon_\nu/10^{14} {\rm eV})^{1/2}$ for $\epsilon_\nu >10^{14}$ eV, although it may be difficult to"
is complex (with five components).,is complex (with five components).
 Interestingly. all three LAEs observed in the HUDF are identified as single component svstems will asymmetric. extended emission.," Interestingly, all three LAEs observed in the HUDF are identified as single component systems with asymmetric, extended emission."
 However. (his morphology. which is observable on both the GOODS and the ΗΙΟΕ images. is probably not representative of the overall LAE population.," However, this morphology, which is observable on both the GOODS and the HUDF images, is probably not representative of the overall LAE population."
 In [act. as shown by 2.. of the ? LAEs are unresolved even al GOODS depth.," In fact, as shown by \citet{Bond09}, of the \citet{GronwallLAE} LAEs are unresolved even at GOODS depth."
 We note (hat our ability to detect multiple components is seusilive to strvev depth., We note that our ability to detect multiple components is sensitive to survey depth.
" This is evident from a comparison of our results from sGOODS subset of GOODS: while 12 of the 17 individual components are unresolved in sGOODS images only 4 of thesel? ""point sources"" remain unresolved al the GOODS depth.", This is evident from a comparison of our results from sGOODS subset of GOODS: while 12 of the 17 individual components are unresolved in sGOODS images only 4 of these12 “point sources” remain unresolved at the GOODS depth.
 It is therefore likely that a deeper imaging survevs would resolve more of the LAE components. and reveal extended emission that is undetectable with current data.," It is therefore likely that a deeper imaging surveys would resolve more of the LAE components, and reveal extended emission that is undetectable with current data."
" In fact. (he majority of the multiple rest-frame UV ""clumps detected on the Tames are probably individual star-forming regions within a single. larger system. or possibly (he result of an ongoing merger."," In fact, the majority of the multiple rest-frame UV “clumps"" detected on the frames are probably individual star-forming regions within a single, larger system, or possibly the result of an ongoing merger."
 For the analvsis presented below. we analvze the morphology of each component individually as well as the LAE svstem as a whole.," For the analysis presented below, we analyze the morphology of each component individually as well as the LAE system as a whole."
 Like the LAE svstems. approximately half of the observed LAE components are unresolved: (his includes 50 of the 95 components in GEMS and 15 of the 31 components in GOODS.," Like the LAE systems, approximately half of the observed LAE components are unresolved: this includes 50 of the 95 components in GEMS and 15 of the 31 components in GOODS."
 Tests performed in Paper I suggest that a S/N230 within a fixed hall-light radius is required. (o robustly determine (he size of a galaxy., Tests performed in Paper I suggest that a $S/N > 30$ within a fixed half-light radius is required to robustly determine the size of a galaxy.
 Thus. in this paper. we measure ihe concentration index and present the results of our GALFIT fits for ellipticites. Sérrsic profiles. aad. hall-light radii for each component with S/N>30. as well as for each LAE svstem with S/N>30.," Thus, in this paper, we measure the concentration index and present the results of our GALFIT fits for ellipticites, Sérrsic profiles, and half-light radii for each component with $S/N > 30$, as well as for each LAE system with $S/N>30$."
 This signal-to-noise eut corresponds to Ve:<28.5 for UDF. Vorc26.8 lor GOODS. and Voy<26.5 lor GEMS.," This signal-to-noise cut corresponds to $V_{GF} <28.5$ for HUDF, $V_{GF} <26.8$ for GOODS, and $V_{GF} < 26.5$ for GEMS."
 The Concentraton. Asvanmetry. ChunupinesS svstem. or CAS (?)— was developed to estimate (he morphology of distant galaxies quantitatively.," The Concentraton, Asymmetry, ClumpinesS system, or CAS \citep{CAS} was developed to estimate the morphology of distant galaxies quantitatively."
 We did not make extensive tests of measuring clumpiness for our sample as the majority of (he sample did not show multiple clumps on visual inspection and those that did only showed a few clumps at most., We did not make extensive tests of measuring clumpiness for our sample as the majority of the sample did not show multiple clumps on visual inspection and those that did only showed a few clumps at most.
 We were also concerned that the low surface brightness and small angular extent of the LAEs would preclude an accurate measurement of clumpiness., We were also concerned that the low surface brightness and small angular extent of the LAEs would preclude an accurate measurement of clumpiness.
 We note that (?) did not attempt measuring clunmpiness for their sample while thev did measure concentration and asvimnmnetry., We note that \citep{Pirzkal07} did not attempt measuring clumpiness for their sample while they did measure concentration and asymmetry.
 We did attempt. to measure asvnuuelry for our sample and found Chat the asymmetry parameter routinelv was derived to bequite small (4< 0.2) even for objects that are, We did attempt to measure asymmetry for our sample and found that the asymmetry parameter routinely was derived to bequite small $A \lesssim 0.2$ ) even for objects that are
its one dav alias.),its one day alias.)
 The CLEANed X-ray power spectra. of confirmed. intermediate polars generally show clear signals at the system periods (see for example Norton ct al 1997. Beardmore et al 1998). and that is not the case here.," The ed X-ray power spectra of confirmed intermediate polars generally show clear signals at the system periods (see for example Norton et al 1997, Beardmore et al 1998), and that is not the case here."
 There are no significant signals at either of the previously reported. periods of this object., There are no significant signals at either of the previously reported periods of this object.
 At periods of both 6.06 hr and 1.996 hr. the power is less than about LO-τ ο.2 2 corresponding to a limiting amplitude in the light curve of <6⋅101 Cs + (," At periods of both 6.06 hr and 1.996 hr, the power is less than about $10^{-7}$ $^{2}$ $^{-2}$, corresponding to a limiting amplitude in the light curve of $<6 \times 10^{-4}$ c $^{-1}$. ("
Nb.,Nb.
 The amplitude is equal to twice the square root of the CLEANed power.), The amplitude is equal to twice the square root of the ed power.)
 The upper limit to any mocdulation in the X-ray light curve at these periods is therefore0., The upper limit to any modulation in the X-ray light curve at these periods is therefore.
34... Furthermore. there are no other significant periods detected either the Nyquist frequency of the time series is around 32«107 Iz and there are no significant," Furthermore, there are no other significant periods detected either – the Nyquist frequency of the time series is around $3 \times 10^{-3}$ Hz and there are no significant"
or each giveu mass accretion rate until we achieved a best-fit.,for each given mass accretion rate until we achieved a best-fit.
" The bes fit of al] was achieved when he WD temperature was T,=27.000lx aud the accretion rate wew LOΛΙ, /vr."," The best fit of all was achieved when the WD temperature was $_{wd} = 27,000$ K and the accretion rate was $10^{-9.5} M_{\odot}$ /yr."
 The 4Z for his fit was 3.11., The $\chi^{2}_{\nu}$ for this fit was 3.41.
 Tve white dwarf contributes only of the [lux aid the disk contributes (see figure 7)., The white dwarf contributes only of the flux and the disk contributes (see figure 7).
 However. the distance obtained was far too large (de 10t) pe) ane| the mass accretion ate rather large lor quiescence.," However, the distance obtained was far too large $\sim$ 460 pc) and the mass accretion rate rather large for quiescence."
 Therefore. again. we rejected this resut.," Therefore, again, we rejected this result."
 As ¢an be seen [rom the igure. the fit is beter in theJUVE range tha it theFUSE rauge. whe'e the model fux is basically 20 too low.," As can be seen from the figure, the fit is better in the range than in the range, where the model flux is basically 20 too low."
 Finally. we coiipared the results of the above fitting atemplts witli a two-teriperature WD uodel fit.," Finally, we compared the results of the above fitting attempts with a two-temperature WD model fit."
 A WD is expected to show a temperature variatjou with latitude if ac‘cretion occurs xeferentially at the equator from a disk or if a WD is maguetie aud accretes p'eferentially at uagnetie poles., A WD is expected to show a temperature variation with latitude if accretion occurs preferentially at the equator from a disk or if a WD is magnetic and accretes preferentially at magnetic poles.
 hi the former case. we elvision the possibility of a hot accretion belt. or hot inner disk rine.," In the former case, we envision the possibility of a hot accretion belt, or hot inner disk ring."
 We ran a series of models in which a WD is cooler aud roates more slowly at ugher latitudes ard has a fast spinuiug (uear Ixepleriau speed) lot atinosplere belt., We ran a series of models in which a WD is cooler and rotates more slowly at higher latitudes and has a fast spinning (near Keplerian speed) hot atmosphere belt.
 We tried a ange of combinations of cooler WD aud belts of different enperatures al(| lower gravity., We tried a range of combinations of cooler WD and belts of different temperatures and lower gravity.
 For he WD we kept logg=8.3 coustai| and searched for a best fit consistei with a distauce of dzz 180pc., For the WD we kept $\log{g}=8.3$ constant and searched for a best fit consistent with a distance of $d \approx 186$ pc.
 The white cwarf plus aceretion belt combination which vielded the best fit has a WD with Tyg=25. 000Ix. aud an accretion belt with Ἐν=10. 000. logg=6 wih solar abundances.," The white dwarf plus accretion belt combination which yielded the best fit has a WD with $_{wd} = 25,000$ K, and an accretion belt with $_{belt} = 40,000$ K, $\log{g} = 6$ with solar abundances."
 The cooler portion of the WD contrib1ues of the FUV flux while the accretion belt contributes of the FUV flux., The cooler portion of the WD contributes of the FUV flux while the accretion belt contributes of the FUV flux.
 This is the same best fit combiued model as the one .or the1 spectrum alone., This is the same best fit combined model as the one for the spectrum alone.
 The X79 value was 3.56op and the scale parameter led to a distance of 190pc., The $\chi^{2}_{\nu}$ value was 3.56 and the scale parameter led to a distance of 190pc.
 Our comparison of the two-temiperature (WD + aceretion belt) moclel is displayed iu figure 8., Our comparison of the two-temperature (WD + accretion belt) model is displayed in figure 8.
 The inproveunient of model fits that result (rom WDs las been reported elsewhere for a number of other systems (e.g. Szkody et al., The improvement of model fits that result from two-temperature WDs has been reported elsewhere for a number of other systems (e.g. Szkody et al.
 2003: Sion et al., 2003; Sion et al.
 20()3)., 2003).
 During the quiescence of WW Ceti. our fits to the combinedFUSE +IUE fluxes reveal hat a single-temperature with Tig26.000& 0001 cau account for the flux.," During the quiescence of WW Ceti, our fits to the combined + fluxes reveal that a single-temperature with $T_{wd} \sim 26,000 \pm 1000$ K can account for the flux."
 The white dwarf appears to have a rotatioua| velocity of GOO4100 s5., The white dwarf appears to have a rotational velocity of $600 \pm 100$ $~$ $^{-1}$.
 The error bars are choset1 here to be the size of the increments by which the parameters are varied., The error bars are chosen here to be the size of the increments by which the parameters are varied.
 For WW Cet. which is actualy a rather weak source forFUSE audIUE. these error bars are cousisteut with t1 inoceliugs. i.e. stnaller error bars/iucremenuts do not lead to a significant improvement in the fit.," For WW Cet, which is actually a rather weak source for and, these error bars are consistent with the modelings, i.e. smaller error bars/increments do not lead to a significant improvement in the fit."
 The best agreement with the observations is provided by a two-temperature white chwarl model wi ha‘ooler whie dwarf at Tig=29. 00018 providing ofthe FUV Mux aud a hotter region (accreion |elt or optically thick disk ring) with T = 10.000Ix contributing of the flux.," The best agreement with the observations is provided by a two-temperature white dwarf model with a cooler white dwarf at $T_{wd} = 25,000$ K providing of the FUV flux and a hotter region (accretion belt or optically thick disk ring) with T = 40,000K contributing of the flux."
 The fitting of tlie shape of the absorption lines in theFUSE range led to the following best fit abuudances: Carbo. 0.1 x solar. Nitrogen 2 x solar. and Silicon 0.3-0.5 x solar.," The fitting of the shape of the absorption lines in the range led to the following best fit abundances: Carbon 0.1 x solar, Nitrogen 2 x solar, and Silicon 0.3-0.5 x solar."
 La all the fitting models of the combined (FUSE+/UE) data we have kept the mass of the WD coustaut (374=O.SAL. correspoudiug to loggy= 8.3) and rejected, In all the fitting models of the combined ) data we have kept the mass of the WD constant $M=0.8M_{\odot}$ corresponding to $\log{g}=8.3$ ) and rejected
sun )) can be made with Lipparcos measurements based on motions of Cepheicls.,Sun ) can be made with Hipparcos measurements based on motions of Cepheids.
 Feast&Whitelock(1997) conelucle that the angular velocity of circular rotation at the Sun. (= Oort’s AD). is 27.19+ORT (218+Tkms flor Ry=8.0 kpe).," \citet{FW97} conclude that the angular velocity of circular rotation at the Sun, (= Oort's A–B), is $27.19 \porm 0.87$ $218 \porm 7$ for $\rnot=8.0$ kpc)."
 Our value of OLHR. obtaimed by removing the Solar Motion in longitude [rom thereflex of the motion of in longitude. is 29.450.15perkpe.," Our value of $\tnot/\rnot$, obtained by removing the Solar Motion in longitude from the of the motion of in longitude, is $29.45\pm0.15$."
. The difference between the VLBA and Hipparcos angular velocities is 2.2640.9perkpe:: these measurements are mareinally consistent., The difference between the VLBA and Hipparcos angular velocities is $2.26\pm0.9$; these measurements are marginally consistent.
 Neither measurements are sensitive to the value of.. as il is primarily used only {ο remove the small contribution of the solar Motion.," Neither measurements are sensitive to the value of, as it is primarily used only to remove the small contribution of the Solar Motion."
 Other measurements of ΟμHj. lor example [rom proper motions of halo stars relative to galaxies by Ialiraietal.(2004)... vielcl consistent. values with slightly greater observational uneertainty.," Other measurements of $\tnot/\rnot,$ for example from proper motions of halo stars relative to galaxies by \citet{Kal04}, yield consistent values with slightly greater observational uncertainty."
" Our value of Oo/itp is a true ""global measure of the angular rotation rate of the Galaxy. as opposed to those derived [rom Oorts constants. which indirectly. determine from the shear and vorticity in the velocity field of material in the solar neighborhood 1986).."," Our value of $\tnot/\rnot$ is a true “global” measure of the angular rotation rate of the Galaxy, as opposed to those derived from Oort's constants, which indirectly determine from the shear and vorticity in the velocity field of material in the solar neighborhood \citep{KL86}. ."
" The small difference between the local (AD) and global measures ol Ou/Ry suggests that local variations in Galactic dvnamies (4(0/R)/dH) are less (han a3kmsfh,", The small difference between the local (A–B) and global measures of $\tnot/\rnot$ suggests that local variations in Galactic dynamics $d(\Theta/R)/dR$ ) are less than $\approx3$.
 We now estimate the peculiar motion of in the direction of Galactic rotation bv subtracting the IHipparcos-based. angular rotation rate of the Galaxy from the VLBA angular motions ofÀ*., We now estimate the peculiar motion of in the direction of Galactic rotation by subtracting the Hipparcos-based angular rotation rate of the Galaxy from the VLBA angular motions of.
. After removing the current best estimate of the motion of the Sun around the Galactic Center of 223 (2184-5.25) (Feast&Whitelock&Dinnev1998) from our VLBA observation of 241.. we find the peculiar motion of is -I8+7 (lowarcl positive Galactic longitude (see Table 3).," After removing the current best estimate of the motion of the Sun around the Galactic Center of 223 $218+5.25$ ) \citep{FW97,DB98} from our VLBA observation of $241$, we find the peculiar motion of is $-18\pm7$ toward positive Galactic longitude (see Table 3)."
 This estimate of (he “in-plane” motion of comes from cdilferencing twoangular motions., This estimate of the “in-plane” motion of comes from differencing two motions.
 Since this difference is small. the uncertamiv in ddoes not strongly affect. this component of the peculiar motion ofA*.," Since this difference is small, the uncertainty in does not strongly affect this component of the peculiar motion of."
. It is unclear al this time whether or not the estimate of (his component of the peculiar motion of differs significantly from zero and. if so. if this indicates a difference between the elobal and local measures of (heangular rotation rate of (he Galaxy or a much larger peculiar," It is unclear at this time whether or not the estimate of this component of the peculiar motion of differs significantly from zero and, if so, if this indicates a difference between the global and local measures of theangular rotation rate of the Galaxy or a much larger peculiar"
"should not only be independent from the processes that govern the grain size distribution, but they should also be able to work on bigger amorphous grains.","should not only be independent from the processes that govern the grain size distribution, but they should also be able to work on bigger amorphous grains."
" Alternatively, the crystalline lattice should be able to keep itself regular during the coagulation of small crystalline dust to create big crystalline grains."," Alternatively, the crystalline lattice should be able to keep itself regular during the coagulation of small crystalline dust to create big crystalline grains."
" The correlation between the strength of the 10 jum feature and the mean grain size in disk surfaces, combined with the lack of correlation between crystallinity fraction and sioumpeak* supports the wide: usage of Speak10um as a proxy for dust size: in: literature: (vanBoekeletal.2003;Kessler-Silacci2006;Pas- 2009)."," The correlation between the strength of the 10 $\mu$ m feature and the mean grain size in disk surfaces, combined with the lack of correlation between crystallinity fraction and $S^{10\mu{\rm m}}_{{\rm
 peak}}$, supports the wide usage of $S^{10\mu{\rm m}}_{{\rm
 peak}}$ as a proxy for dust size in literature \citep{VB03,KE06,PA09}."
". Bouwmanetal.(2008) found a strong correlation between disk geometry and the strength of the 10 wm silicate feature for a very small sample of T Tauri stars (7 which points to flatter disks having shallower 10 µπι disks),features (ie., big grains in the disk surface)."," \citet{BO08} found a strong correlation between disk geometry and the strength of the 10 $\mu$ m silicate feature for a very small sample of T Tauri stars (7 disks), which points to flatter disks having shallower 10 $\mu$ m features (i.e., big grains in the disk surface)."
" Using results from similar decomposition procedures, Olofssonetal. and Juhászetal. confirm this trend for larger (2010)samples of T Tauri (2010)(58 disks) and Herbig Ae/Be stars (45 disks), respectively."," Using results from similar decomposition procedures, \citet{OF10} and \citet{JU10} confirm this trend for larger samples of T Tauri (58 disks) and Herbig Ae/Be stars (45 disks), respectively."
 Those trends are much weaker than that found by Bouwmanetal. showing a larger spread.," Those trends are much weaker than that found by \citet{BO08}, showing a larger spread."
" For the current even larger (2008),,sample (139 disks), no significant trend is seen, indicating that the earlier small sample trends may have been affected by a few outliers."," For the current even larger sample (139 disks), no significant trend is seen, indicating that the earlier small sample trends may have been affected by a few outliers."
 This result is similar to that found by Oliveiraetal. for a large YSO sample (e 200 objects) using the (2010)strength of the 10 pm silicate feature as a proxy for grain size (Figure 14 in that , This result is similar to that found by \citet{OL10} for a large YSO sample $\sim$ 200 objects) using the strength of the 10 $\mu$ m silicate feature as a proxy for grain size (Figure 14 in that paper).
"As discussed by Oliveira et al.,"," As discussed by Oliveira et al.,"
 the sedimentation paper).models of Dullemond&Dominik(2008) expect a strong correlation of larger grains in flatter disks that is not seen., the sedimentation models of \citet{DD08} expect a strong correlation of larger grains in flatter disks that is not seen.
 This means that sedimentation alone cannot be responsible for the distribution of mean grain sizes in the upper layers of protoplanetary disks around T 'Tauri stars., This means that sedimentation alone cannot be responsible for the distribution of mean grain sizes in the upper layers of protoplanetary disks around T Tauri stars.
" Furthermore, the lack of correlation between crystallinity fraction and disk geometry is not in support of the results of Watsonetal. and Sargentetal. (2009),, who find a link between (2009)increasing crystallinity fraction and dust sedimentation."," Furthermore, the lack of correlation between crystallinity fraction and disk geometry is not in support of the results of \citet{WA09} and \citet{ST09}, who find a link between increasing crystallinity fraction and dust sedimentation."
" As discussed in Oliveiraetal.(2010) for Serpens and Taurus, and confirmed by the addition of considerably older samples, there is no clear difference in the mean grain sizes in the disk surfaces with mean cluster age, which can be seen in Figure 11.."," As discussed in \citet{OL10} for Serpens and Taurus, and confirmed by the addition of considerably older samples, there is no clear difference in the mean grain sizes in the disk surfaces with mean cluster age, which can be seen in Figure \ref{f_grain}."
" This evidence supports the discussion in that paper that the dust population observed in the disk surface cannot be a result of a progressive, monotonic change of state from small amorphous grains, to large, more crystalline grains, or ‘grain growth and processing’."," This evidence supports the discussion in that paper that the dust population observed in the disk surface cannot be a result of a progressive, monotonic change of state from small amorphous grains, to large, more crystalline grains, or `grain growth and processing'."
" The fact that the distribution of grain sizes in the upper layers of disks does not change with cluster age implies that an equilibrium of the processes of dust growth and fragmentation must exist, which also supports the existence of small grains in disks that are millions of years old whereas dust growth is a rapid process (Weidenschilling1980;Dullemond&Dominik2005)."," The fact that the distribution of grain sizes in the upper layers of disks does not change with cluster age implies that an equilibrium of the processes of dust growth and fragmentation must exist, which also supports the existence of small grains in disks that are millions of years old whereas dust growth is a rapid process \citep{WE80,DD05}."
". That small dust is still seen in disks in older regions like Upper Sco and η Cha argues that this equilibrium of processes is maintained for millions of years, as long as the disks are optically thick, but independent of them having a flared or flatter geometry."," That small dust is still seen in disks in older regions like Upper Sco and $\eta$ Cha argues that this equilibrium of processes is maintained for millions of years, as long as the disks are optically thick, but independent of them having a flared or flatter geometry."
" Literature studies of disk fractions of different YSO clusters with different mean ages show a trend of decreasing disk fraction, i.e. disks dissipating with time, over some few millions of years (Haischetal.2001;Hernándezetal. 2008)."," Literature studies of disk fractions of different YSO clusters with different mean ages show a trend of decreasing disk fraction, i.e. disks dissipating with time, over some few millions of years \citep{HA01,HE08}."
. This decrease is clearly confirmed by the lower fraction of disks still present in the older regions studied here (Upper Sco and 7 Cha)., This decrease is clearly confirmed by the lower fraction of disks still present in the older regions studied here (Upper Sco and $\eta$ Cha).
" According to current planet formation theories, if giant planets are to be formed from gas rich disks, the optically thin, gas-poor disks in those older regions should already harbor (proto-)planets."," According to current planet formation theories, if giant planets are to be formed from gas rich disks, the optically thin, gas-poor disks in those older regions should already harbor (proto-)planets."
" Considering the evidence from small bodies in our own Solar System that suggest considerably higher crystallinity fractions than ISM dust (see Woodenetal.2007 and Pontoppidan&Brearley2010 for reviews of latest results), a crystallinity increase must occur."," Considering the evidence from small bodies in our own Solar System that suggest considerably higher crystallinity fractions than ISM dust (see \citealt{WO07} and \citealt{PB10} for reviews of latest results), a crystallinity increase must occur."
" In Figure 12,, the mean crystallinity fraction per region is plotted against two evolutionary parameters: disk fraction (left) and mean age (right)."," In Figure \ref{f_comp2}, the mean crystallinity fraction per region is plotted against two evolutionary parameters: disk fraction (left) and mean age (right)."
" Within the spread in individual fractions it is seen that, just as for grain sizes, there is no strong evidence of an increase of crystallinity fraction with either evolutionary parameter."," Within the spread in individual fractions it is seen that, just as for grain sizes, there is no strong evidence of an increase of crystallinity fraction with either evolutionary parameter."
 This implies that there is no evolution in grain sizes or crystallinity fraction for the dust in the surface of disks, This implies that there is no evolution in grain sizes or crystallinity fraction for the dust in the surface of disks
"along large-scale structure ""filaments"" — simulations have shown that large scale structure may contribute only about 10 per cent to the cluster surface mass density (Wambsganss. Bode. Ostriker 2005: Hilbert et al.","along large-scale structure “filaments” – simulations have shown that large scale structure may contribute only about 10 per cent to the cluster surface mass density (Wambsganss, Bode, Ostriker 2005; Hilbert et al."
 2007)., 2007).
 Instead. there is a real and large variation in the total-mass-to-optical-light ratio among clusters.," Instead, there is a real and large variation in the total-mass-to-optical-light ratio among clusters."
 The low mass-to-light ratio of RCS cluster cores may be caused by a bias in favour of line-of-sight mergers in the optical selection process. a prominent and spectacular example of which is CIO0244-24 (Czoske et al.," The low mass-to-light ratio of RCS cluster cores may be caused by a bias in favour of line-of-sight mergers in the optical selection process, a prominent and spectacular example of which is Cl0024+24 (Czoske et al."
 2002)., 2002).
 Indeed. extensive spectroscopic follow-up of RCS clusters has uncovered several cases of close projection effects of possibly physically associated systems as well as line-of-sight substructure (Gilbank et al.," Indeed, extensive spectroscopic follow-up of RCS clusters has uncovered several cases of close projection effects of possibly physically associated systems as well as line-of-sight substructure (Gilbank et al."
 2007: Cain et al., 2007; Cain et al.
 2008)., 2008).
 A further effect to consider is the question of whether X-ray selection may favour the inclusion of clusters that are in the process of merging., A further effect to consider is the question of whether X-ray selection may favour the inclusion of clusters that are in the process of merging.
 Torri et al. (, Torri et al. (
2004) have found that. during a merger. the lensing cross section is increased by a factor of 510 for a duration of a couple of hundred million years. while the X-ray luminosities of merging clusters are increased by a factors of 5.,"2004) have found that, during a merger, the lensing cross section is increased by a factor of $5-10$ for a duration of a couple of hundred million years, while the X-ray luminosities of merging clusters are increased by a factors of $\sim
 5$."
 A similar conclusion regarding the X-ray luminosity of clusters during mergers was reached by Randall. Sarazin. Ricker (2002).," A similar conclusion regarding the X-ray luminosity of clusters during mergers was reached by Randall, Sarazin, Ricker (2002)."
 If X-ray-selected cluster samples indeed have a larger fraction of merging clusters. one could thus expect a larger fraction of highly efficient lenses in those samples.," If X-ray-selected cluster samples indeed have a larger fraction of merging clusters, one could thus expect a larger fraction of highly efficient lenses in those samples."
 Thus. the masses of the X-ray-selected clusters may be systematically overestimated as well.," Thus, the masses of the X-ray-selected clusters may be systematically overestimated as well."
 An interesting question is whether comparable optically and X-ray-selected cluster samples at zc0.7 also differ in their arc production efficiencies., An interesting question is whether comparable optically and X-ray-selected cluster samples at $z > 0.7$ also differ in their arc production efficiencies.
 We did not find giant arcs in any of the high-redshift RCS clusters we analyzed. even though their optical luminosities are comparable to those of the RCS clusters at low and medium redshifts.," We did not find giant arcs in any of the high-redshift RCS clusters we analyzed, even though their optical luminosities are comparable to those of the RCS clusters at low and medium redshifts."
 This contrasts with the results of Gladders et al. (, This contrasts with the results of Gladders et al. (
2003) who found RCS clusters to be more efficient lenses at high redshift.,2003) who found RCS clusters to be more efficient lenses at high redshift.
 Finally. the many ares found in the MACS low- and medium-redshift subsamples provide a statistically improved handle on the angular distribution of ares in clusters.," Finally, the many arcs found in the MACS low- and medium-redshift subsamples provide a statistically improved handle on the angular distribution of arcs in clusters."
" Our results show that ares do form at large angular separations from cluster centres. at up to 60"". in some cases."," Our results show that arcs do form at large angular separations from cluster centres, at up to $60''$, in some cases."
 Thus. the large Einstein radius of Abell 1689 is probably not unique.," Thus, the large Einstein radius of Abell 1689 is probably not unique."
 We have conducted an algorithmically based search for lensed ares in 100 clusters observed with HST., We have conducted an algorithmically based search for lensed arcs in $\sim 100$ clusters observed with HST.
 Our cluster sample includes an X-ray selected subsample (XBACs: MACS) and an optically selected subsample (RCS). each in a range of redshifts.," Our cluster sample includes an X-ray selected subsample (XBACs; MACS) and an optically selected subsample (RCS), each in a range of redshifts."
 Our search for giant ares has produced 12. 17. and 13 ares (/w7 10) in the XBACs. MACS low-redshift. and MACS medium-redshift subsamples. respectively.," Our search for giant arcs has produced $12$, $17$, and $13$ arcs $l/w > 10$ ) in the XBACs, MACS low-redshift, and MACS medium-redshift subsamples, respectively."
 Only 2. 5. and zero ares were found in the low-. medium-. and high-redshift RCS subsamples.," Only $2$, $5$, and zero arcs were found in the low-, medium-, and high-redshift RCS subsamples."
 The are production efficiency of the MACS clusters is therefore higher by a factor of 5.IO than that of the RCS clusters., The arc production efficiency of the MACS clusters is therefore higher by a factor of $5-10$ than that of the RCS clusters.
 The typical Einstein radii of MACS clusters are several times larger than those of the relatively few RCS cluster that do display strong lensing., The typical Einstein radii of MACS clusters are several times larger than those of the relatively few RCS cluster that do display strong lensing.
 If. as we suspect. the HST sample of RCS clusters was pre-selected in a way hat favored strong lenses. then these conclusions would only be strengthened.," If, as we suspect, the HST sample of RCS clusters was pre-selected in a way that favored strong lenses, then these conclusions would only be strengthened."
 These results constitute direct evidence. based on strong-ensing statistics. that optically selected RCS clusters are an order of magnitude less massive than X-ray selected clusters. despite heir similar optical properties.," These results constitute direct evidence, based on strong-lensing statistics, that optically selected RCS clusters are an order of magnitude less massive than X-ray selected clusters, despite their similar optical properties."
 This conclusion is supported by the actor-100 higher space density of RCS clusters., This conclusion is supported by the factor-100 higher space density of RCS clusters.
 In the are statistics iterature to date. the observed statistics from X-ray and optical clusters have often been discussed together and interchangeably.," In the arc statistics literature to date, the observed statistics from X-ray and optical clusters have often been discussed together and interchangeably."
 We have demonstrated that X-ray and optically selected clusters ikely probe distinct parts of the cluster mass function. and should herefore not be mixed in this way.," We have demonstrated that X-ray and optically selected clusters likely probe distinct parts of the cluster mass function, and should therefore not be mixed in this way."
 In a forthcoming paper. we will address are statistics from a heoretical point of view.," In a forthcoming paper, we will address arc statistics from a theoretical point of view."
 We will present strong lensing statistics wredictions using clusters from several of the latest cosmological, We will present strong lensing statistics predictions using clusters from several of the latest cosmological
are also strong sources at COAL?PEL energies (von Montigny οἱ al.,are also strong sources at COMPTEL energies (von Montigny et al.
 1995. et al.," 1995, Mukherjee et al."
 1997)., 1997).
 estimates of the blazar contribution Alukherjeeo the GRB rely on theQuantitative assumptions on theταν evolution and vary between 20 andspeculative 90% (ο. Müccke & Pohl 1998. Sreekumar et al.," Quantitative estimates of the blazar contribution to the GRB rely on the speculative assumptions on the$\gamma$ -ray evolution and vary between $20$ and $90$ % (e.g. Müccke & Pohl 1998, Sreekumar et al."
 1997 ancl references therein)., 1997 and references therein).
 In the present paper we discuss 5-ràv emission by luminous in the far-infrared (PLR) domain and the »»tential rolegalaxies of these objects for the GRB., In the present paper we discuss $\gamma$ -ray emission by galaxies luminous in the far-infrared (FIR) domain and the potential role of these objects for the GRB.
 We caleulate the of ο αν by cosmic ray electrons scattering low energy. photons., We calculate the production of $\gamma$ -rays by cosmic ray electrons scattering low energy photons.
production electrons produce the 7-rays: aJ with the interstellarEnergetic medium (Bremsstrahlung) and scatteringinteracting olf the cosmic microwave background (CAIB) and intrinsic galaxy6) radiation via the inverse Compton (1C), Energetic electrons produce the $\gamma$ -rays: interacting with the interstellar medium (Bremsstrahlung) and scattering off the cosmic microwave background (CMB) and intrinsic galaxy radiation via the inverse Compton (IC)
of sight velocities to velocities relative to the Local Group (Yahiletal.1977).,of sight velocities to velocities relative to the Local Group \cite{yah77}.
.. We have calculated stellar masses from the absolute ue magnitude of the galaxy assuming a stellar mass to light ratio of one (as suggested by Stavelv-Smith et al. (, We have calculated stellar masses from the absolute blue magnitude of the galaxy assuming a stellar mass to light ratio of one (as suggested by Stavely-Smith et al. (
1990) and de Blok et al. (,1990) and de Blok et al. (
1996) for gas rich galaxies) and an absolute Xue magnitude for the Sun of AIF=5.4 (Banksοἱal. 1999).,1996) for gas rich galaxies) and an absolute blue magnitude for the Sun of $M_{\odot}^{B}=5.4$ \cite{ban99}.
. We found that all of the photographie magnitudes isted in (Alorshicli-Esslinecretal.L999a) were significantly zünter than the available CCD magnitudes Listed in NED., We found that all of the photographic magnitudes listed in \cite{mor99a} were significantly fainter than the available CCD magnitudes listed in NED.
 We have used the NED CCD photometry wherever possible and made the photographic magnitudes. brighter by the mean of the CCD correction where this was not. possible., We have used the NED CCD photometry wherever possible and made the photographic magnitudes brighter by the mean of the CCD correction where this was not possible.
 The mean correction was 0.2 magnitudes., The mean correction was 0.2 magnitudes.
 1n figure 6 we show the distribution of (Mgi/Lg). for our sample., In figure \ref{fig:histml} we show the distribution of $(M_{HI}/L_{B})_{\odot}$ for our sample.
 Lt is quite clear that this sample consists of ealaxies with extraordinary values of (Mgi/Lg). (all of the S galaxies with (Magi1).<1 have (Aly;/Le).> 0.3)., It is quite clear that this sample consists of galaxies with extraordinary values of $(M_{HI}/L_{B})_{\odot}$ (all of the 8 galaxies with $(M_{HI}/L_{B})_{\odot}<1$ have $(M_{HI}/L_{B})_{\odot}>0.3$ ).
 The relative eas mass of these galaxies is lar higher than that of a typical’ spiral galaxy. (xnapp1900)., The relative gas mass of these galaxies is far higher than that of a 'typical' spiral galaxy \cite{kna90}.
 1n figure 7 we show (Mgi/Lg). plotted against Alp., In figure \ref{fig:mvml} we show $(M_{HI}/L_{B})_{\odot}$ plotted against $M_{B}$.
 Although there is some scatter a clear trend exists for the fainter galaxies to have larger values of (Mg/Le).., Although there is some scatter a clear trend exists for the fainter galaxies to have larger values of $(M_{HI}/L_{B})_{\odot}$.
 A leas squares fit gives (Ala;/Le).LyUlULL an exponen very close to the value of -0.3|/-0.1 found by Stavelv-Smith et al. (, A least squares fit gives $(M_{HI}/L_{B})_{\odot} \propto L_{B}^{-0.4+/-0.1}$ an exponent very close to the value of -0.3+/-0.1 found by Stavely-Smith et al. (
1992) for a sample of LL rich chwarl galaxies. but the (Stavely-Smithctal.1992). sample was optically selecte and has much lower values of (Alay/Le). than this sample.,"1992) for a sample of HI rich dwarf galaxies, but the \cite{sta92} sample was optically selected and has much lower values of $(M_{HI}/L_{B})_{\odot}$ than this sample."
 In the same wav that selection at anv wavelength predominantly selects. objects that are bright at. tha wavelength. the combination of relatively faint optica sources with LIE selection has led to a sample with large values of (Alay/Le).., In the same way that selection at any wavelength predominantly selects objects that are bright at that wavelength the combination of relatively faint optical sources with HI selection has led to a sample with large values of $(M_{HI}/L_{B})_{\odot}$.
 Thus. although small. we have constructed a sample of galaxies that apparently have turned only a small fraction of their gas into stars.," Thus, although small, we have constructed a sample of galaxies that apparently have turned only a small fraction of their gas into stars."
The primary Cosmic Microwave Background (CMB) and especially its angular power spectrum provides us with powerful constraints on the content of the universe and its evolution.,The primary Cosmic Microwave Background (CMB) and especially its angular power spectrum provides us with powerful constraints on the content of the universe and its evolution.
 It is now well established that an accurate understanding of the primary CMB power spectrum requires a good understanding of the secondary CMB anisotropies resulting from the interaction of the CMB photons with the matter along the line of sight from the last scattering surface to the observer (see?.forareview)..., It is now well established that an accurate understanding of the primary CMB power spectrum requires a good understanding of the secondary CMB anisotropies resulting from the interaction of the CMB photons with the matter along the line of sight from the last scattering surface to the observer \citep[see][for a review]{Aghanimrevue08}.
 The great efforts to understand these secondary anisotropies. in order to best recover the primary CMB. also provide us with powerful independent cosmological probes when the secondary anisotropies are regarded as a source of information rather than contamination.," The great efforts to understand these secondary anisotropies, in order to best recover the primary CMB, also provide us with powerful independent cosmological probes when the secondary anisotropies are regarded as a source of information rather than contamination."
 Among those secondary CMB anisotropies. some result from the gravitational interaction of the CMB photons with the potential wells they cross.," Among those secondary CMB anisotropies, some result from the gravitational interaction of the CMB photons with the potential wells they cross."
 Integrated Sachs-Wolfe (SW) effect. they pass through large scale time evolving potential wells (2)..," Integrated Sachs-Wolfe (ISW) effect, they pass through large scale time evolving potential wells \citep{SachsWolfe1967}."
 Since a dark energy like component is expected to attect the growth of large scale structures.shallower.a detection of the ISW effect is an important probe establishing existence - the Universe is flat and general relativity is a correc description of gravity - and constraining the equation of state of such a component.," Since a dark energy like component is expected to affect the growth of large scale structures, detection of the ISW effect is an important probe establishing existence - the Universe is flat and general relativity is a correct description of gravity - and constraining the equation of state of such a component."
 Detection claims of the ISW effect arose as soon as the firs year WMAP data were released., Detection claims of the ISW effect arose as soon as the first year WMAP data were released.
 Those claims were based upon cross correlation analyses of WMAP CMB data andsurveys., Those claims were based upon cross correlation analyses of WMAP CMB data and.
 While most of the firs analyses were conducted in real space (ie.. by computing the angular cross-correlation function). subsequently new results basec upon Fourier/multipole and wavelet space were presented.," While most of the first analyses were conducted in real space (i.e., by computing the angular cross-correlation function), subsequently new results based upon Fourier/multipole and wavelet space were presented."
 The results from the WMAP team on the cross correlation of NRAO Very Large Sky Survey (NVSS) with WMAP data (2?) were soon followed by other analyses applied not only on NVSS data. but also on X-ray and optical based catalogs like HEAO. SDSS. APM or 2MASS. (22222)..," The results from the WMAP team on the cross correlation of NRAO Very Large Sky Survey (NVSS) with WMAP data \citep{Nolta2004} were soon followed by other analyses applied not only on NVSS data, but also on X-ray and optical based catalogs like HEA0, SDSS, APM or 2MASS, \citep{Boughn2003-4,Fosalba2003,Scranton2003,Fosalba2004,Afshordi2MASS}."
 As subsequent data releases from both the CMB and the SDSS side became public. new studies prompted further evidence for significant cross correlation between CMB," As subsequent data releases from both the CMB and the SDSS side became public, new studies prompted further evidence for significant cross correlation between CMB"
oue tend to have simaller magnification factors m the sense that the outer inage is much brighter.,one tend to have smaller magnification factors in the sense that the outer image is much brighter.
 Because of this there will be a bias toward cases where C is near oue., Because of this there will be a bias toward cases where $C$ is near one.
 To fit real leus svstemi a more complicated. asviunetrie lens models must be used aud C must be calculated for cach paix of mages separately.," To fit real lens system a more complicated, asymmetric lens models must be used and $C$ must be calculated for each pair of images separately."
 This quantity can be evaluated at the ceuter of a jet ππαρσο or at a kink iu a jet image to determine if the bed is consistent with au iutrinsic feature iu the jet itself or requires substructure as an explanation., This quantity can be evaluated at the center of a jet image or at a kink in a jet image to determine if the bend is consistent with an intrinsic feature in the jet itself or requires substructure as an explanation.
 Two explanations for the apparent bend in image B of D1152|199 will be explored., Two explanations for the apparent bend in image B of B1152+199 will be explored.
 One is that image A actually has a small undetected curvature which is magnified in iiage D svhere it is detected., One is that image A actually has a small undetected curvature which is magnified in image B where it is detected.
 The second explanation is that image A is straight aud image D is bent by the influence of a substructure near it., The second explanation is that image A is straight and image B is bent by the influence of a substructure near it.
 Tuvestigating both of these hypothesis requires fitting a host lens model to the positious of the images and he center of the lens., Investigating both of these hypothesis requires fitting a host lens model to the positions of the images and the center of the lens.
 Since there are ouly two images in this case a complicated host lens 1odol is not well constrained by the positions alone. (, Since there are only two images in this case a complicated host lens model is not well constrained by the positions alone. \markcite{2002MNRAS.330..205R}{ (
2002) ft to cach VLBI image a poit source for the core aud a Gaussian for the jet: these positious are used as coustraints.,2002) fit to each VLBI image a point source for the core and a Gaussian for the jet; these positions are used as constraints.
 We choose to use a simple SIS model with a ckeround shear - ας)---Lrcost20.)|c?siut20.j|. a?j(x)25Lrx20.)4?cost20.j|.," We choose to use a simple SIS model with a background shear - $\alpha^1({\bf x}) = \gamma \left[ x^1\cos(2\theta_\gamma) +
x^2\sin(2\theta_\gamma)\right]$, $\alpha^2({\bf x}) = \gamma \left[
x^1\sin(2\theta_\gamma) - x^2\cos(2\theta_\gamma) \right]$."
 The shear xeaks the azimuthal sxauuetry of the host leus which is necessary for it to fit the observed leus position., The shear breaks the azimuthal symmetry of the host lens which is necessary for it to fit the observed lens position.
 No attempt is made to incorporate the possible cavarf companion of the Ίος galaxy that appears as à very aut snudge in the TST image., No attempt is made to incorporate the possible dwarf companion of the lens galaxy that appears as a very faint smudge in the HST image.
 We do uot expect that this object is large enough to significantly chanec he surface potential except in its near vicinity and the umaees are well separated from it., We do not expect that this object is large enough to significantly change the surface potential except in its near vicinity and the images are well separated from it.
 In addition. the quality of the ft discussed in 23.2.1. gives us confidence that the model accurately reproduces the local nagnification matrix at the positions of the images which is the ouly thing needed here.," In addition, the quality of the fit discussed in \ref{sec:with-no-substructure}
 gives us confidence that the model accurately reproduces the local magnification matrix at the positions of the images which is the only thing needed here."
" With the reported redshifts the critical deusity for this leus is X,=2.65«10956;AL.kpe"," With the reported redshifts the critical density for this lens is $\Sigma_c = 2.65\times 10^9
h_{65}\msun\kpc^{-2}$."
 A smooth model is fit to the positions of the lens ealaxy. the radio cores of the images and the ceuter of the jet images.," A smooth model is fit to the positions of the lens galaxy, the radio cores of the images and the center of the jet images."
 A model is found that fits all the positions to better than 0.1 1illi-aresecoud., A model is found that fits all the positions to better than 0.1 milli-arcsecond.
 In addition. the imaeuification ratio of the radio core agrees with the observed one to better than despite this not being used as α coustraiut ou the model.," In addition, the magnification ratio of the radio core agrees with the observed one to better than despite this not being used as a constraint on the model."
 This signifies that the local magnification matrix. A. is being accurately reproduced by the model.," This signifies that the local magnification matrix, ${\bf\tilde{A}}$, is being accurately reproduced by the model."
 The velocity dispersion of the leus is thos=217dans!1 aud the backerouud shear is y=0.102., The velocity dispersion of the lens is $\sigma_{\rm host}=247\kms$ and the background shear is $\gamma=0.102$.
 This velocity dispersion is not unusual for a leus galaxy., This velocity dispersion is not unusual for a lens galaxy.
 The estimated circular velocity is Vane=V20 yar.," The estimated circular velocity is $V_{\rm
circ}=\sqrt{2}\sigma_{\rm host}$ ."
 The maeuificatious at the positions of the radio cores are joy=3.8 and pp=1.5 a negative maenification indicates a one dimensional parity flip in the image., The magnifications at the positions of the radio cores are $\mu_A=3.8$ and $\mu_B=-1.5$ – a negative magnification indicates a one dimensional parity flip in the image.
 This model gives a curvature naenification factor of C=L9 at the ceuter of the jet with image D beiug the more curved of the two images as observed., This model gives a curvature magnification factor of $C=4.9$ at the center of the jet with image B being the more curved of the two images as observed.
 If the jet in image A has a curvature of 1/C times the curvature in image D and it is in the right direction then the observations can be explained without substructure., If the jet in image A has a curvature of $1/C$ times the curvature in image B and it is in the right direction then the observations can be explained without substructure.
 Figure 2. shows some attempts o model the jet in this way., Figure \ref{fig:map_jets_nosub} shows some attempts to model the jet in this way.
 Frou visual inspection it appears that the jet in image A is not bent cuoueh to explain the bend in image D. The curve should follow the crest of the jets surface brightuess. but a jet that is heut enough requires the eud of the jet to be shifted by ~ο Linas from the crest of the straight jet.," From visual inspection it appears that the jet in image A is not bent enough to explain the bend in image B. The curve should follow the crest of the jet's surface brightness, but a jet that is bent enough requires the end of the jet to be shifted by $\sim 3-4$ mas from the crest of the straight jet."
 The joan ds large m this dimension. 3.6 mas. but a shift inthe crest should be detectable below this level.," The beam is large in this dimension, $3.6$ mas, but a shift inthe crest should be detectable below this level."
Observations of the infrared background. provide important information on the emission. of cosmic Luminous sources throughout the history of the Universe.,Observations of the infrared background provide important information on the emission of cosmic luminous sources throughout the history of the Universe.
 Lt has been suggested. (Santos. Momm hamionkowski 2003: SalvaterraSteopps Ferrara“oppeype 2003) that⊳ a⊳ Fsρου fraction οἱ the measured. —Near-Iafralted (1-10. pam) cosmic Background (NIIU) arises [rom redshifted Lye line photons and nebular emission produced by the first very massive nmetal-[ree stars.," It has been suggested (Santos, Bromm Kamionkowski 2003; Salvaterra Ferrara 2003) that a large fraction of the measured Near-InfraRed (1-10 $\mu$ m) cosmic Background (NIRB) arises from redshifted $\alpha$ line photons and nebular emission produced by the first very massive metal-free stars."
 ‘This ivpothesis. however. is very demanding in terms of the required conversion cllicicney of barvons into stars (Alacau Silk 2005).," This hypothesis, however, is very demanding in terms of the required conversion efficiency of baryons into stars (Madau Silk 2005)."
 X large NIRB contribution [rom such stars has more recently been rejected by the paucity CE 3) of z~10 candidate sources in Llubble Space Telescope ultra-deep observations (Salvaterra Ferrara 2005)., A large NIRB contribution from such stars has more recently been rejected by the paucity $\le 3$ ) of $z\sim 10$ candidate sources in Hubble Space Telescope ultra-deep observations (Salvaterra Ferrara 2005).
 Nevertheless. a more moclest contribution from very high redshift galaxies. whose clustering should leave a clistinet signature on small-scale angular Uuetuations of the background. light (Magliocchetti. Salvaterra Ferrara 2003: IWashlinsky et al.," Nevertheless, a more modest contribution from very high redshift galaxies, whose clustering should leave a distinct signature on small-scale angular fluctuations of the background light (Magliocchetti, Salvaterra Ferrara 2003; Kashlinsky et al."
 2004: Coorayv οἱ al., 2004; Cooray et al.
 2004). is still possible.," 2004), is still possible."
 ]xashlinsky et al. (, Kashlinsky et al. (
2005) have recently found significant NIRB Buctuations in deep exposure data obtained with Spitzer/IRAC (Fazio ct al.,2005) have recently found significant NIRB fluctuations in deep exposure data obtained with Spitzer/IRAC (Fazio et al.
 2004a. 2004b) in four channels (3.6. 4.5. 5.8. and 8 pim). after Galactic stars ancl galaxies bright enough to be individually resolved by he instrument have been carefully subtracted.," 2004a, 2004b) in four channels (3.6, 4.5, 5.8, and 8 $\mu$ m), after Galactic stars and galaxies bright enough to be individually resolved by the instrument have been carefully subtracted."
 With the only exception of the 8 pm channel. the shape and amplitude of the power spectrum cannot be reproduced by either contributions from intervening clustv. Galactic neutral hydrogen. gas (cirrus) or from local interplanetary dust. (zodiacal light).," With the only exception of the 8 $\mu$ m channel, the shape and amplitude of the power spectrum cannot be reproduced by either contributions from intervening dusty, Galactic neutral hydrogen gas (cirrus) or from local interplanetary dust (zodiacal light)."
 Ordinary galaxies. (2- £5) produce Iluctuations. due to their. clustering. and shot-noise., Ordinary galaxies $z \lsim 5$ ) produce fluctuations due to their clustering and shot-noise.
 The faint Lux limits (20.3 55) of Spitzer data allow to push 1reir residual clustering contribution below the level of the excess signal at. relatively large. (=50 arcsec) angular scaes (Ixashlinsky et al., The faint flux limits $\geq 0.3$ $\mu$ Jy) of Spitzer data allow to push their residual clustering contribution below the level of the excess signal at relatively large $\gsim 50$ arcsec) angular scales (Kashlinsky et al.
 2005)., 2005).
 The shot nolse component. estimated cirectly from galaxy counts. fits the observed Iuctilations at smaller angular scales. ancl rapiclly faces away at arger angles.," The shot noise component, estimated directly from galaxy counts, fits the observed fluctuations at smaller angular scales, and rapidly fades away at larger angles."
 The residual large scale signal has been ascrixxb by Washlinsky ct al. (, The residual large scale signal has been ascribed by Kashlinsky et al. (
2005) as coming [rom very disant (22- 5) sources. provided. their total lux contribution is lnWam 7 ο.,2005) as coming from very distant $z\ge 5$ ) sources provided their total flux contribution is $> 1$ nW $^{-2}$ $^{-1}$.
 ‘Phe aim of this Letter is to show that this is inceed the case., The aim of this Letter is to show that this is indeed the case.
 The lavout of the paper is as follows: in Section 2 we will briclly describe the adopted model. while in Section 3 we provide predictions for the NUIRB intensity and Buctuations and compare the latter ones with the results of Ixashlinsky et al. (," The layout of the paper is as follows: in Section 2 we will briefly describe the adopted model, while in Section 3 we provide predictions for the NIRB intensity and fluctuations and compare the latter ones with the results of Kashlinsky et al. ("
2005).,2005).
 Section 4 summarizes our conclusions., Section 4 summarizes our conclusions.
"where ay is the fine structure constant. r, is the classical electron radius. In.Ac20 is the Coulomb logarithm. οἱ~Cy: and we have neglected logarithmic corrections to the relativistic Iree-Iree emissivitv.","where $\alpha_f$ is the fine structure constant, $r_e$ is the classical electron radius, $\ln{\Lambda}\simeq20$ is the Coulomb logarithm, $c_s^2\simeq c_{sp}^2$, and we have neglected logarithmic corrections to the relativistic free-free emissivity."
" The subscript ""IV in ή denotes relativistic Bremsstrallune."," The subscript “R” in $Q_{\rm ff,R}$ denotes relativistic Bremsstrahlung."
" In the one-temperature regime. both protons aud electrons ave cool and non-relativistic. and have nearly the same temperature. hence cl,9~(i,πρ]κ and ος9£wae)265, "," In the one-temperature regime, both protons and electrons are cool and non-relativistic, and have nearly the same temperature, hence $c_{sp}^2\simeq(m_e/m_p)c_{se}^2$ and $c_s^2\approx 2c_{sp}^2$ ."
":The two energv equations (9)) and (11)) can be combined to vield a single energy. equation for the accreling gas: —cieο where the [ree-Iree cooling takes (he form Quxn|52:EL.23/28""Eq-νο2e.. where op is the Thompson cross-section. aud the subscript NI stands lor non-relativistic."," The two energy equations \ref{4}) ) and \ref{5}) ) can be combined to yield a single energy equation for the accreting gas: -c_s^2 where the free-free cooling takes the form q^-= ^2 }, where $\sigma_T$ is the Thompson cross-section, and the subscript ${\rm NR}$ stands for non-relativistic."
 As a result of high density of the gas in the BL. optically thin bremsstrahlung cooling dominates over sell-absorbed svnchrotron cooling: hence the latter may salely be neglected.," As a result of high density of the gas in the BL, optically thin bremsstrahlung cooling dominates over self-absorbed synchrotron cooling; hence the latter may safely be neglected."
 We therefore neglect svnchrotron emission in our analvsis., We therefore neglect synchrotron emission in our analysis.
 In our model. we neglect the effects of radiation pressure compared to the gas pressure (which is of order M/Mg<1 and. hence. neeligible at low accretion rates) and Comptonization (which must be important in a high-temperature region. but is not important closer to the stellar surface. where the eas temperature is low. see more discussion below).," In our model, we neglect the effects of radiation pressure compared to the gas pressure (which is of order $\dot M/\dot M_{\rm Edd}\ll1$ and, hence, negligible at low accretion rates) and Comptonization (which must be important in a high-temperature region, but is not important closer to the stellar surface, where the gas temperature is low, see more discussion below)."
 For simplicity. we neglect also thermal conduction.," For simplicity, we neglect also thermal conduction."
 Thus. our present model is similar to the models we used in our previous studies of hot accretion.," Thus, our present model is similar to the models we used in our previous studies of hot accretion."
 Our simplified hydrodynamic model is. therefore. very instructive.," Our simplified hydrodynamic model is, therefore, very instructive."
 It allows us to study the DL problem on (he same grounds. on which the other hot flow solutions have been treated. as a viscous. radiative. purely hvdrodynamic flow.," It allows us to study the BL problem on the same grounds, on which the other hot flow solutions have been treated, — as a viscous, radiative, purely hydrodynamic flow."
 Oncethe basic, Oncethe basic
this case. Tololo 0109-383 would probably be the most X-ray soft among classical Seyferts.,"this case, Tololo 0109-383 would probably be the most X–ray soft among classical Seyferts."
 At least part of the soft excess may however be due to several unresolved emission lines from photoionized plasma. similarly to what found in NGC 1068 (Kinkhabwala et al.," At least part of the soft excess may however be due to several unresolved emission lines from photoionized plasma, similarly to what found in NGC 1068 (Kinkhabwala et al."
 2002). weakening the case for a very steep continuum.," 2002), weakening the case for a very steep continuum."
 Unfortunately. the source is too faint for the RGS to be profitably used. and we cannot check directly this hypothesis.," Unfortunately, the source is too faint for the RGS to be profitably used, and we cannot check directly this hypothesis."
 Combining previous ASCA and BeppoSAX results (Collinge Brandt 2000; Iwasawa et al., Combining previous ASCA and BeppoSAX results (Collinge Brandt 2000; Iwasawa et al.
 2001) with the Chandra and XMM-Newton observations discussed here. we can conclude that at least two X-ray absorbers are present: one Compton—thick. obscuring the nucleus on à small scale. the other Compton-thin. obscuring the soft X-ray emission and possibly the extended emission.," 2001) with the $Chandra$ and $Newton$ observations discussed here, we can conclude that at least two X–ray absorbers are present: one Compton--thick, obscuring the nucleus on a small scale, the other Compton–thin, obscuring the soft X–ray emission and possibly the extended emission."
 The latter may be associated with the dust lanes observed by HST to obscure the central part of the galaxy (Malkan et al., The latter may be associated with the dust lanes observed by HST to obscure the central part of the galaxy (Malkan et al.
 1998)., 1998).
 This provides one more piece of evidence in favour of the co-existence of both Compton-thin and Compton-thick matter in the circumnuclear regions of Seyfert galaxies (see e.g. Matt Guainazzi 2002 and references therein)., This provides one more piece of evidence in favour of the co–existence of both Compton--thin and Compton–thick matter in the circumnuclear regions of Seyfert galaxies (see e.g. Matt Guainazzi 2002 and references therein).
 The problem here ts that the optical broad lines appears to be seen through the dust lane rather than the Compton-thick X-ray absorber., The problem here is that the optical broad lines appears to be seen through the dust lane rather than the Compton–thick X-ray absorber.
 Recently. it has become clear that a fraction of type | nuclei are absorbed in X-rays (e.g. Maiolino et al.," Recently, it has become clear that a fraction of type 1 nuclei are absorbed in X–rays (e.g. Maiolino et al."
 2001a: Fiore et al., 2001a; Fiore et al.
 2001) and. more generally. that X-ray absorption column densities are often much larger than would be expected from the amount of optical extinction (e.g. Granato et al.," 2001) and, more generally, that X–ray absorption column densities are often much larger than would be expected from the amount of optical extinction (e.g. Granato et al."
 1997)., 1997).
 While a fraction of hard X-ray spectrum. type | AGN may be explained in terms of à temporary off of the nucleus. which makes them for a while reflection—dominated and so apparently Compton-thick absorbed (Matt et al.," While a fraction of hard X-ray spectrum, type 1 AGN may be explained in terms of a temporary switching--off of the nucleus, which makes them for a while reflection--dominated and so apparently Compton–thick absorbed (Matt et al."
 2002). this is certainly not the case for Tololo 0109-383. whose nucleus is definitely absorbed by Compton-thick matter (Iwasawa et al.," 2002), this is certainly not the case for Tololo 0109-383, whose nucleus is definitely absorbed by Compton–thick matter (Iwasawa et al."
 2001)., 2001).
 Let us call this matter. for simplicity. the ‘torus’.," Let us call this matter, for simplicity, the `torus'."
 There are several possible solutions to this problem (see Matolino et al., There are several possible solutions to this problem (see Maiolino et al.
 2001b for a discussion)., 2001b for a discussion).
 First of all. the Broad Line Region (BLR) may be located outside the torus.," First of all, the Broad Line Region (BLR) may be located outside the torus."
 However. the typical BLR size. from reverberation mapping studies. is usually of the order of light-days or of light-weeks. at least for a moderately luminous source as Tololo 0109-383. while the inner surface of the torus is expected to have a size of a fractior of a pe or more (e.g. Bianchi et al.," However, the typical BLR size, from reverberation mapping studies, is usually of the order of light-days or of light-weeks, at least for a moderately luminous source as Tololo 0109-383, while the inner surface of the torus is expected to have a size of a fraction of a pc or more (e.g. Bianchi et al."
 2001 and references therein)., 2001 and references therein).
 Moreover. the ratio between the fluxes of the broad and narrow components is very low (less than 1. Murayama et al.," Moreover, the ratio between the fluxes of the broad and narrow components is very low (less than 1, Murayama et al."
 1998)., 1998).
 Therefore. even though we cannot exclude that the size of either the BLR or the torus. or both. are different than usual in this object (we do not have any direct measurement of them) this possibility seems unlikely.," Therefore, even though we cannot exclude that the size of either the BLR or the torus, or both, are different than usual in this object (we do not have any direct measurement of them) this possibility seems unlikely."
 Another possibility is that the dust-to-gas ratio of the absorber is very low. due to dust sublimation (Granato et al.," Another possibility is that the dust–to–gas ratio of the absorber is very low, due to dust sublimation (Granato et al."
 1997)., 1997).
 However. the Ηα/Ηβ ratio is similar for the narrow and broad components. suggesting that they are observed through the same absorber. which is probably the dust lane observed by HST.," However, the $\alpha$ $\beta$ ratio is similar for the narrow and broad components, suggesting that they are observed through the same absorber, which is probably the dust lane observed by HST."
 Therefore. dust sublimation must be almost complete. to avoid further extinction of the broad components. and the low fluxes of the broad components would remain unexplained.," Therefore, dust sublimation must be almost complete, to avoid further extinction of the broad components, and the low fluxes of the broad components would remain unexplained."
 An alternative solution. proposed by Matolino et al. (," An alternative solution, proposed by Maiolino et al. ("
2001b). is that the sizes of the dust grains are much larger than that in the ISM of our own Galaxy. changing dramatically the extinction curves.,"2001b), is that the sizes of the dust grains are much larger than that in the ISM of our own Galaxy, changing dramatically the extinction curves."
 This solution has the merit of explaining qualitatively the low fluxes of the broad line components. but the expected extinction. given the X-ray measured column density. would be too high to allow them to be observed at all (see the figures in Maiolino et al.," This solution has the merit of explaining qualitatively the low fluxes of the broad line components, but the expected extinction, given the X–ray measured column density, would be too high to allow them to be observed at all (see the figures in Maiolino et al."
 200109)., 2001b).
 All these problems can be avoided if the broad lines are actually seen in reflected. rather than direct. light.," All these problems can be avoided if the broad lines are actually seen in reflected, rather than direct, light."
 Indeed. the broad lines in this object are best seen in polarized light (Toran et al.," Indeed, the broad lines in this object are best seen in polarized light (Moran et al."
 2000) and therefore at least part of them must be reflected., 2000) and therefore at least part of them must be reflected.
 It is possible that the reflecting medium ts the same responsible for the soft X-ray excess and the O and Ne lines we observe in the Chandra and XMM-Newton spectra., It is possible that the reflecting medium is the same responsible for the soft X–ray excess and the O and Ne lines we observe in the $Chandra$ and $Newton$ spectra.
" One may wonder why in this source the reflected light would be so ntense to permit the detection of broad lies in direct light. ""Sespite the fact that the polarization degree is by no means exceptional."," One may wonder why in this source the reflected light would be so intense to permit the detection of broad lines in direct light, despite the fact that the polarization degree is by no means exceptional."
 This is possible if the covering factor and optical 0epth of the reflecting matter are large enough: it must be =recalled that a large covering factor implies a low polarization degree. for obvious geometrical reasons.," This is possible if the covering factor and optical depth of the reflecting matter are large enough; it must be recalled that a large covering factor implies a low polarization degree, for obvious geometrical reasons."
 Unfortunately. this hypothesis cannot be readily tested in X-rays because the evidence of a spectral break in the nuclear spectrum makes it difficult to estimate the nuclear-to-scattered flux ratio.," Unfortunately, this hypothesis cannot be readily tested in X–rays because the evidence of a spectral break in the nuclear spectrum makes it difficult to estimate the nuclear–to–scattered flux ratio."
 To summarize. the Chandra and XMM-Newton observations. together with previous X-ray and optical observations. suggest the following scenario for Tololo 0109-383.," To summarize, the $Chandra$ and $Newton$ observations, together with previous X–ray and optical observations, suggest the following scenario for Tololo 0109-383."
 The nucleus is absorbed by Compton-thick. material. the nuclear radiation being reflected by: a) cold material (probably the inner wall of the torus) giving rise to the Compton reflection component and the iron Ka line (and. by reprocessing. to the infrared emission. Matt et al.," The nucleus is absorbed by Compton–thick material, the nuclear radiation being reflected by: a) cold material (probably the inner wall of the torus) giving rise to the Compton reflection component and the iron $\alpha$ line (and, by reprocessing, to the infrared emission, Matt et al."
 2000 and Iwasawa et al., 2000 and Iwasawa et al.
 2001): b) ionized matter. responsible for the soft X-ray excess. and the oxygen and neon He-like lines.," 2001); b) ionized matter, responsible for the soft X–ray excess, and the oxygen and neon He–like lines."
 This tonized matter may coincide with that reflecting and polarizing the otherwise obscured BLR., This ionized matter may coincide with that reflecting and polarizing the otherwise obscured BLR.
 Further material. partly spatially resolved. is responsible for the HINER: the size of the spatially extended HINER emission ( | kpe) is similar to that of the extended soft X-ray emission. but probably the two regions are not associated with each other.," Further material, partly spatially resolved, is responsible for the HINER; the size of the spatially extended HINER emission $\sim$ 1 kpc) is similar to that of the extended soft X–ray emission, but probably the two regions are not associated with each other."
 Finally. all these components are seen through a dust lane. responsible for the Balmer decrement and the absorption of the soft X-ray emission.," Finally, all these components are seen through a dust lane, responsible for the Balmer decrement and the absorption of the soft X–ray emission."
We have detected a significant weak lensing signal for a saluple of 116 intermediate redshift galaxy groups.,We have detected a significant weak lensing signal for a sample of 116 intermediate redshift galaxy groups.
 From the lensing signal we estimate that galaxy eroups have a mean ML of L85428 hM ../Lp.. within 1 IN pec. aud that this M/L is constant as the distance from the group center increases.," From the lensing signal we estimate that galaxy groups have a mean M/L of $\pm$ 28 $_\odot$ $_{B\odot}$ within 1 $^{-1}$ Mpc, and that this M/L is constant as the distance from the group center increases."
 When the sample is split into subsets of rich. aud poor galaxy groups. there is a clear offset iu the mass-to-light ratios of the two subsets.," When the sample is split into subsets of rich and poor galaxy groups, there is a clear offset in the mass-to-light ratios of the two subsets."
 The increase in. the M/Li as à functionB. ofB mass is. in. ecucral agreement with other results. but is detected bere for the first time using weak lensing in the galaxy group mass regine.," The increase in the M/L as a function of mass is in general agreement with other results, but is detected here for the first time using weak lensing in the galaxy group mass regime."
 This analysis indicates that a weak leusime signal can indeed. be measured from! galaxy eroups., This analysis indicates that a weak lensing signal can indeed be measured from galaxy groups.
 Clearly. a larger siuuple with. well determined. ανασα.. properties- would be ideal for this sort of study.," Clearly, a larger sample with well determined dynamical properties would be ideal for this sort of study."
 The structure of the dark matter halos of ealaxy eroups are still poorly understood., The structure of the dark matter halos of galaxy groups are still poorly understood.
" By combining this group Ieusine result with salaxxy-otealaxyi lensing it should be ""possible to determune the size an exten- OFF gaalaxvUSS Srouper carsn matCrhalos, dados, ἩwhichHen willwi aid significantly in our wucderstancing of structure in the Universe and the nature of dark matter."," By combining this group lensing result with galaxy-galaxy lensing it should be possible to determine the size and extent of galaxy group dark matter halos, which will aid significantly in our understanding of structure in the Universe and the nature of dark matter."
 MIST. acknowledges support from NSERC aud a Premiers Research Excellence Award., MJH acknowledges support from NSERC and a Premier's Research Excellence Award.
To determine the abundances we need an initial estimation of the microturbulent velocity (6).,To determine the abundances we need an initial estimation of the microturbulent velocity $\xi$ ).
 For this estimation we have used the standard method., For this estimation we have used the standard method.
 We computed the abundances from the Fe lines for a range of possible values of ἕ satisfying two conditions: a) that the abundances of Fe lines were not dependent on the equivalent widths and b) that the rms errors were minima., We computed the abundances from the Fe lines for a range of possible values of $\xi$ satisfying two conditions: a) that the abundances of Fe lines were not dependent on the equivalent widths and b) that the rms errors were minima.
 To achieve the first condition the slope in the plot abundance vs £ must be zero., To achieve the first condition the slope in the plot abundance vs $\xi$ must be zero.
 We tried different & values to fulfill this requirement., We tried different $\xi$ values to fulfill this requirement.
 In this sense the abundance and microturbulent velocity determinations are recursive and simultaneous., In this sense the abundance and microturbulent velocity determinations are recursive and simultaneous.
 Once a & value has been fixed the corresponding abundances to all chemical species measured are determined using the WIDTH9 code., Once a $\xi$ value has been fixed the corresponding abundances to all chemical species measured are determined using the WIDTH9 code.
 The WIDTH? code requires the model atmosphere calculated by the ATLAS9 program. the equivalent width of each line as well as atomic constants such as oscillator strength (Log ef) values. excitation potentials. damping constants. ete.," The WIDTH9 code requires the model atmosphere calculated by the ATLAS9 program, the equivalent width of each line as well as atomic constants such as oscillator strength (Log gf) values, excitation potentials, damping constants, etc."
 In particular for the Log gf we used Puhretal.(1988) and Kuruez(1992)., In particular for the Log gf we used \citet{fuhr88} and \citet{kurucz92}.
. This code calculates the theoretical equivalent widths for an initial input abundance and compares these values with the measured equivalent widths., This code calculates the theoretical equivalent widths for an initial input abundance and compares these values with the measured equivalent widths.
 Then the code modifies the abundance to achieve a difference between theoretical and measured equivalent widths « 0.01mA., Then the code modifies the abundance to achieve a difference between theoretical and measured equivalent widths $<$ 0.01.
. The final values of the metallicities corresponding to the N93 and C97 calibrations. are listed in Table 2..," The final values of the metallicities corresponding to the N93 and C97 calibrations, are listed in Table \ref{width.metal}."
 We have included the number of lines used in each determination as well as the rms of the average., We have included the number of lines used in each determination as well as the rms of the average.
 To estimate errors for our WIDTH metallicities we consider the following facts., To estimate errors for our WIDTH metallicities we consider the following facts.
 The most significant contribution to the final uncertainties. probably. comes from the equivalent width measurements.," The most significant contribution to the final uncertainties, probably, comes from the equivalent width measurements."
 We assume a error due to the continuum level determination., We assume a error due to the continuum level determination.
 This translates into maximum uncertainties in the metallicity estimation., This translates into maximum uncertainties in the metallicity estimation.
 The atomic constants may also have uncertainties., The atomic constants may also have uncertainties.
 In particular we estimate that the oscillator strength values may cause differences of about in the calculated metallicity., In particular we estimate that the oscillator strength values may cause differences of about in the calculated metallicity.
" Finally to provide an estimation of ""typical"" errors introduced by the WIDTH method we increased the T. by 150 K and the Log ef by 0.15. and recalculated the metallicity value for each star."," Finally to provide an estimation of ”typical” errors introduced by the WIDTH method we increased the $_{\rm eff}$ by 150 K and the Log gf by 0.15, and recalculated the metallicity value for each star."
 We derived a median difference of 0.20 dex., We derived a median difference of 0.20 dex.
 The largest difference corresponds to HD 28978 (0.55 dex)., The largest difference corresponds to HD 28978 (0.55 dex).
 The WIDTH method is not practical when the number of stars is large., The WIDTH method is not practical when the number of stars is large.
 For each object. we need to identify and measure many spectral lines.," For each object, we need to identify and measure many spectral lines."
 An alternative would be to compare the observed spectra with a grid of synthetic ones corresponding to different values of the metallicites and choose from the grid the spectrum that better reproduces the observed data (Grayetal.2001)., An alternative would be to compare the observed spectra with a grid of synthetic ones corresponding to different values of the metallicites and choose from the grid the spectrum that better reproduces the observed data \citep{gray01}.
. This comparison has the advantage that the complete profiles of the lines and not only the equivalent widths are used in the metallicity determinations., This comparison has the advantage that the complete profiles of the lines and not only the equivalent widths are used in the metallicity determinations.
 In general synthetic spectra depend on four parameters: Ty. surface gravity (Log g). metallicity ([Fe/H|) and microturbulent velocity (£).," In general synthetic spectra depend on four parameters: $_{\rm eff}$ , surface gravity (Log g), metallicity ([Fe/H]) and microturbulent velocity $\xi$ )."
 Following Grayetal.(2001).. we applied a multidimensional Downhill Simplex technique. in which the observed spectrum Is compared to," Following \citet{gray01}, we applied a multidimensional Downhill Simplex technique, in which the observed spectrum is compared to"
be a precious information for people dealing with stellar evolution and. nucleosynthesis.,be a precious information for people dealing with stellar evolution and nucleosynthesis.
 Finally. we address the problem of +°O) destruction in stars.," Finally, we address the problem of $^{17}$ O destruction in stars."
 ο is destroyed in stellar. interiors hy proton captures through the reactions Ο ((p. α) FN and ορ. 5) 7F (this latter occurs only at high temperatures).," $^{17}$ O is destroyed in stellar interiors by proton captures through the reactions $^{17}$ (p, $\alpha$ $^{14}$ N and $^{17}$ (p, $\gamma$ $^{18}$ F (this latter occurs only at high temperatures)."
 MI the models previously discussed assume that all the 1 O injected into the ISM by stars of dilferent masses is newly. produced. i.c. all the /O present in the protostellar nebula is destroved in the hot stellar interior.," All the models previously discussed assume that all the $^{17}$ O injected into the ISM by stars of different masses is newly produced, i.e., all the $^{17}$ O present in the protostellar nebula is destroyed in the hot stellar interior."
 In Fig., In Fig.
 9a.b we compare results from Mocellines) to those [rom Aloclelfines}.," 9a,b we compare results from Model to those from Model."
 Model is the same as Model 32. except. for the fact that now all the pristine ο O is assumed to survive and to be returned back into the ISM at the death of the star.," Model is the same as Model , except for the fact that now all the pristine $^{17}$ O is assumed to survive and to be returned back into the ISM at the death of the star."
 A realistic situation should be probably something in between the two., A realistic situation should be probably something in between the two.
 Assuming that all the IO present in the gasout, Assuming that all the $^{17}$ O present in the gasout
"The procedure consists of taking several 7},- ατανε, sets. trving to reproduce the spectral variations with each of them.","The procedure consists of taking several $T_{t_0}$ $\alpha_{t_0}$ $a_{V t_0}$ sets, trying to reproduce the spectral variations with each of them."
 First. the simplest component (a power law) is fitted to the initial data. obtaining γι.," First, the simplest component (a power law) is fitted to the initial data, obtaining $\alpha_{t_0}$."
 From) this fit. we look for an explanation to the spectral variation. ic. performing changes of amplitude. of the slope. or both. in the diagnosis plane.," From this fit, we look for an explanation to the spectral variation, i.e., performing changes of amplitude, of the slope, or both, in the diagnosis plane."
 If the explanation is satistactory. for simplicity we consider that the OAL event is consistent with changes in this unique coupoucut.," If the explanation is satisfactory, for simplicity we consider that the OM event is consistent with changes in this unique component."
 Otherwise. an extra (thermal) colmpoucnt must be added to fit the initial data.," Otherwise, an extra (thermal) component must be added to fit the initial data."
 When this second (thermal) component is added to ft the data. two families of characterization curves are generated: one for thermal variations (as those showed in Fieure 2)) aud one for nou- variations (as those showed in Figure 3)).," When this second (thermal) component is added to fit the data, two families of characterization curves are generated: one for thermal variations (as those showed in Figure \ref{pks15vardos}) ) and one for non-thermal variations (as those showed in Figure \ref{pks15alfas}) )."
 Coutrasting the data with each famuly of curves. we can determine whether a spectral variation is consistent with a chanee in the thermal coiiponeut (ic. changes of T; aud/or nz;) or m the thermal one (1.0... changes of o; aud/or nj;).," Contrasting the data with each family of curves, we can determine whether a spectral variation is consistent with a change in the thermal component (i.e., changes of $T_t$ and/or $n_{T t}$ ) or in the non-thermal one (i.e., changes of $\alpha_t$ and/or $n_{n t}$ )."
" After comparing the data with all the curves.we find that we can set the initial spectral index of the nou-thermal compoucut. αρ, at ~1. Then. we take this fit as representative of cach familyof curves. as shown in Figures 7."," After comparing the data with all the curves,we find that we can set the initial spectral index of the non-thermal component, $\alpha_{t_0}$, at $\sim -1$, Then, we take this fit as representative of each familyof curves, as shown in Figures 7."
" For this reason. the estimations for the values of Ty. αι aud (sz, must be considered nore as a sugecstion than as the actual conditions for both disk aud jet."," For this reason, the estimations for the values of $T_{t_0}$, $\alpha_{t_0}$, and $a_{V t_0}$ must be considered more as a suggestion than as the actual conditions for both disk and jet."
 We have explored how these assuniptious cau affect the discerument of the OM origin., We have explored how these assumptions can affect the discernment of the OM origin.
 As it cau be appreciated in Figures 2-6.. depending ou their origin. the curves show a different behavior.," As it can be appreciated in Figures \ref{pks15vardos}- \ref{pks15z}, depending on their origin, the curves show a different behavior."
" With different asstuuptions of values of T,,. ay. aud ayy, the trajectories for uon-thermal variations are very simular."," With different assumptions of values of $T_{t_0}$, $\alpha_{t_0}$, and $a_{V t_0}$, the trajectories for non-thermal variations are very similar."
" For this reason. we believe that the choice of &4,=—1 has no influence on the description of the spectral variation."," For this reason, we believe that the choice of $\alpha_{t_0} = -1$ has no influence on the description of the spectral variation."
 Expression (S)) shows. in fact. that (vty. OF Og 0nd Dag have more influence on the behavior of the trajectories than oy.," Expression \ref{ec5.16}) ) shows, in fact, that $a_{V t_0}$ $\alpha_t -\alpha_{t_0}$ , and $n_{n t}$ have more influence on the behavior of the trajectories than $\alpha_{t_0}$."
 Additionally. it is easy to," Additionally, it is easy to"
sub-regimes in the Lall diffusivity domain and three in the Ohm domain.,sub-regimes in the Hall diffusivity domain and three in the Ohm domain.
" Lt also led to analytic estimates of the magnetically reduced: (clue to magnetic pressure-gracdient ""wqueezing) density scale height ancl of the location of the κος surface. (where the inflow turns into an outflow). as well as of other pertinent quantities."," It also led to analytic estimates of the magnetically reduced (due to magnetic pressure-gradient `squeezing') density scale height and of the location of the disc's surface (where the inflow turns into an outflow), as well as of other pertinent quantities."
 These results are summarized in Appendix A., These results are summarized in Appendix \ref{sec:appA}.
 In this paper we test. these predictions by constructing exact solutions of the cise equations., In this paper we test these predictions by constructing exact solutions of the disc equations.
 We concentrate on the Hall regime in view of its expected importance in the inner regions of real svstenis: we do not consider the low-ionization Ohm regime in this paper given that it may have limited relevance to winde-driving protostellar discs., We concentrate on the Hall regime in view of its expected importance in the inner regions of real systems; we do not consider the low-ionization Ohm regime in this paper given that it may have limited relevance to wind-driving protostellar discs.
 We characterize the solutions in terms of the conductivitv-tensor components (i.e. the Pedersen. all and Obm conductivities).," We characterize the solutions in terms of the conductivity-tensor components (i.e. the Pedersen, Hall and Ohm conductivities)."
 However. to facilitate the comparison with the analytic results of Paper L which were derived in the framework of the multilluid formulation. we assume that the ratios of these terms are constant with height in the disc and. more specifically. that they scale with the density and field amplitude as p/27. which implies that the matterfield coupling parameter (the Llsasscr number A) is also constant with height.," However, to facilitate the comparison with the analytic results of Paper I, which were derived in the framework of the multifluid formulation, we assume that the ratios of these terms are constant with height in the disc and, more specifically, that they scale with the density and field amplitude as $\rho/B^2$, which implies that the matter–field coupling parameter (the Elsasser number $\Lambda$ ) is also constant with height."
 We require. the derived solutions to cross the sonic critical surface but we do not continue the integration past that surface: this is sullicient for the comparison with the analytic results and greatly simplifies the caleulations., We require the derived solutions to cross the sonic critical surface but we do not continue the integration past that surface; this is sufficient for the comparison with the analytic results and greatly simplifies the calculations.
 However. as recapitulated below. we also demonstrate that these solutions can be matched to wind solutions that extend to laree distances (and. in particular. cross the Alfvénn critical surface).," However, as recapitulated below, we also demonstrate that these solutions can be matched to wind solutions that extend to large distances (and, in particular, cross the Alfvénn critical surface)."
 Our findings can be summarized as follows., Our findings can be summarized as follows.
 We also detail the procedure for obtaining. global (raclially self-similar) ‘cole’ wine solutions following the methodology. introduced by DPS2., We also detail the procedure for obtaining global (radially self-similar) `cold' wind solutions following the methodology introduced by BP82.
 We compute solutions of this twpe for a [aree range of values of the wind. mocel xuwanmeters s. À and £i (tbe normalized mass-to-Iux ratio. specific angular momentum and field-line inclination at the discs surface. respectively).," We compute solutions of this type for a large range of values of the wind model parameters $\kappa$, $\lambda$ and $\xi'_{\rm b}$ (the normalized mass-to-flux ratio, specific angular momentum and field-line inclination at the disc's surface, respectively)."
 Tables of these solutions are available on the Vizielt data base of astronomical catalogues (http://edsarc.u-strasbg.fr/)., Tables of these solutions are available on the VizieR data base of astronomical catalogues (http://cdsarc.u-strasbg.fr/).
 As our raclially localized. and ecometrically thin model cannot be used. to. follow. the propagation of the outflow far from the disc. we match our disc solution to a DPS2-tvpe wind solution by adjusting one of the dise model parameters (c) and iterating on the disc and wind caleulations until the full solution converges.," As our radially localized and geometrically thin model cannot be used to follow the propagation of the outflow far from the disc, we match our disc solution to a BP82-type wind solution by adjusting one of the disc model parameters $\epsilon$ ) and iterating on the disc and wind calculations until the full solution converges."
 We present illustrative solutions of this tvpe that demonstrate that matched. disk/wind configurations can be obtained [or parameter values that are very. similar to those of the merely transonic solutions employed in our parameter-space analysis., We present illustrative solutions of this type that demonstrate that matched disk/wind configurations can be obtained for parameter values that are very similar to those of the merely transonic solutions employed in our parameter-space analysis.
 The aceretion process in protostellar clises may involve a varicty of angular-momentun transport mechanisms. including. in particular. radial transport by gravitational torques and by MIV-induced: turbulence.," The accretion process in protostellar discs may involve a variety of angular-momentum transport mechanisms, including, in particular, radial transport by gravitational torques and by MRI-induced turbulence."
 In. this. paper we consider only vertical transport. by centrifugally driven winds in an attempt to model a radially localized: disc region where this mechanism may dominate. (, In this paper we consider only vertical transport by centrifugally driven winds in an attempt to model a radially localized disc region where this mechanism may dominate. (
Note. however. that both vertical transport ancl radial transport — notably AlBU-induced turbulence — could in principle operate at the same cise radius: see Salmeron.Woniel&Wardle 2007..),"Note, however, that both vertical transport and radial transport – notably MRI-induced turbulence – could in principle operate at the same disc radius; see \citealt*{SKW07}. .)"
 As discussed in Paper L the larec-scale. ordered. magnetic field. envisioned in this scenario could be either interstellar Ποιά adveeted by the aceretion [low or dynamo-generated field. produced in either the star or the disc.," As discussed in Paper I, the large-scale, ordered magnetic field envisioned in this scenario could be either interstellar field advected by the accretion flow or dynamo-generated field produced in either the star or the disc."
 In view of the strong evidence for strong outllows from the inner regions of protostellar disces. we also neglect. alternative modes of angular momentum transport that could. be mediated: by such a field. including magnetic braking. [πιο winds and non-steady phenomena.," In view of the strong evidence for strong outflows from the inner regions of protostellar discs, we also neglect alternative modes of angular momentum transport that could be mediated by such a field, including magnetic braking, `failed' winds and non-steady phenomena."
 Our treatment has been deliberately, Our treatment has been deliberately
To calculate the bremsstrahluneg emissivity a sensible model for quark interactions has to be used.,To calculate the bremsstrahlung emissivity a sensible model for quark interactions has to be used.
 We adopt the basic assumption that the details of the interaction should be relatively unimportant. as long as the overall strength is correct and the relevant svimnetries are respected.," We adopt the basic assumption that the details of the interaction should be relatively unimportant, as long as the overall strength is correct and the relevant symmetries are respected."
 Hence we use the well-known Nambu-Jona-Lasinio (\JL) model in its 90(2) version in our calculations1994)., Hence we use the well-known Nambu-Jona-Lasinio (NJL) model in its $SU(2)$ version in our calculations.
. It has the same svnuuetries as QCD and describes an effective pointlike interaction with a constant coupling strength., It has the same symmetries as QCD and describes an effective pointlike interaction with a constant coupling strength.
 Usually the NJL model and its extensions are used in mean field caleulations., Usually the NJL model and its extensions are used in mean field calculations.
 As a consequence of adopting this model all the aspects related to color superconductivitv are neglected., As a consequence of adopting this model all the aspects related to color superconductivity are neglected.
 Let doy be (he cross-section for a given process of scattering of charged particles. which mav be accompanied bv (he emission of a certain number of photons.," Let $d\sigma _{0}$ be the cross-section for a given process of scattering of charged particles, which may be accompanied by the emission of a certain number of photons."
For example. doy could reler to (he scattering of a quark by an other quark. wilh the possible emission of hard photons.,"For example, $d\sigma _{0}$ could refer to the scattering of a quark by an other quark, with the possible emission of hard photons."
 Together with this process one could consider another process which ciffers from it only in that one extra photon is emitted., Together with this process one could consider another process which differs from it only in that one extra photon is emitted.
" In (his case the total cross-section do can be represented as a product of (wo independent. factors. the cross-section da) and the probability dV, of emission of a single photon in the collision1982)."," In this case the total cross-section $d\sigma $ can be represented as a product of two independent factors, the cross-section $d\sigma _{0}$ and the probability $dW_{\gamma }$ of emission of a single photon in the collision."
. The enission of a soft photon is a quasi-classical process., The emission of a soft photon is a quasi-classical process.
 The probability of emission is (he sine as the classically caleulated number of quanta emitted in the collision. that is tlie same as (heclassical intensity (total energy) of emission df. divided bv the lvequencey of the radiation w1982).," The probability of emission is the same as the classically calculated number of quanta emitted in the collision, that is the same as theclassical intensity (total energy) of emission $dI$, divided by the frequency of the radiation $\omega $."
 Hence as a first step in obtaining the bremsstrahlung emissivity of quark matter we have to calculate the classical radiation intensity emitted by a quark moving in a dense medium in which many inter-particle collisions occur., Hence as a first step in obtaining the bremsstrahlung emissivity of quark matter we have to calculate the classical radiation intensity emitted by a quark moving in a dense medium in which many inter-particle collisions occur.
 The probability dM. of emitting a photon of energy w by any of the scattered quarks is dM.—difw., The probability $dW_{\gamma }$ of emitting a photon of energy $\omega $ by any of the scattered quarks is $dW_{\gamma }=dI/\omega $.
 Therefore the total cross section for the emission of soft bremsstrahlung photons is given by To caleulate the total emission we need (to know crosssectionforaphotonday. describing (he scattering of quarks.," Therefore the total cross section for the emission of soft bremsstrahlung photons is given by To calculate the total cross section for a photon emission we need to know $%
d\sigma _{0 describing the scattering of quarks."
 Within the 6wo-flavor NJL model the elementary quark-quark cross sections have been caleulated by (1995).The calculations have been done nonperturbativelv in (he coupling constant. using tlie so-called 1/:N. expansion. where ἂν. is of the thenumberofcolors.," Within the two-flavor NJL model the elementary quark-quark cross sections have been calculated by .The calculations have been done nonperturbatively in the coupling constant, using the so-called $1/N_{c}$ expansion, where $N_{c}$ is the number of colors."
Irfirstorder1/.N. expansion ancl for temperature 7.<1. where 7.=0.19 GeV. (he quark-quark scattering eross-sections have a remarkably simple dependence on the collision energv s. given by o5(5)~L/s 1995).," In the first order of the $%
1/N_{c expansion and for temperature $T<T_{c}$, where $T_{c}=0.19$ GeV, the quark-quark scattering cross-sections have a remarkably simple dependence on the collision energy $s$ , given by $\sigma _{0}(s)\sim 1/s$ ."
. The cross section decreases with increasing energv., The cross section decreases with increasing energy.
 For a more general paranmetrization of (he, For a more general parametrization of the
"]t is. well known that elliptical⊀⊀ galaxies osin clusters. of ealaxies. have the tight. correlation. between their. colours and magnitudes. (so-called ""colour.magnitude. relation⋠↼: hereafter⋅ CALR).",It is well known that elliptical galaxies in clusters of galaxies have the tight correlation between their colours and magnitudes (so-called `colour–magnitude relation'; hereafter CMR).
 For example.1 the rms scatter about the mean CAL is tvpically≜∣ ~ 0.04 mag in mMVirgo and ]Coma elusters of. galaxies.," For example, the rms scatter about the mean CMR is typically $\sim$ 0.04 mag in Virgo and Coma clusters of galaxies."
: Lt is: à comparable size: to observational. errors (Bower. Lucey Ellis 1992).," It is a comparable size to observational errors (Bower, Lucey Ellis 1992)."
 Because it is believed that this tight relation rellects the formation and evolution processes of elliptical ealaxies.. many people have studied. this. relationship. ..in order to understand the galaxy formation. process.," Because it is believed that this tight relation reflects the formation and evolution processes of elliptical galaxies, many people have studied this relationship in order to understand the galaxy formation process."
 sPhe traditional- scenario. of the formation⋅⊀ of ⋅∢∢ellipticals is that a monolithic protogalactie cloud collapses and then forms⋅ stars for⋅ a short time-scale. until. blowing. the galactic. wind (Larson 1974: Arimoto Yoshit 1986. 1987).," The traditional scenario of the formation of ellipticals is that a monolithic protogalactic cloud collapses and then forms stars for a short time-scale until blowing the galactic wind (Larson 1974; Arimoto Yoshii 1986, 1987)."
 In this framework. theconventionad interpretation of the CALR of ellipticals is considered: as follows.," In this framework, the interpretation of the CMR of ellipticals is considered as follows."
 Colour of cach galaxy, Colour of each galaxy
"In Figures ]-a and 2-a we present the 1.3 mm continuum and (2-1) emission of 122198 and A5142 down to 0.4"".",In Figures 1-a and 2-a we present the 1.3 mm continuum and (2–1) emission of I22198 and A5142 down to $0.4''$.
" For 122198 we detected one strong source. MM? (following the nomenclature in Sánnchez-Monge 22010). extended in the southeast-northwest direction. with a faint extension at 60>, MM2-S. 0.8” to the southeast."," For I22198 we detected one strong source, MM2 (following the nomenclature in Sánnchez-Monge 2010), extended in the southeast-northwest direction, with a faint extension at $\sigma$ , MM2-S, $0.8''$ to the southeast."
 The overall extended emission of MM2 is perpendicular to the direction of outflow A (Sánnchez-Monge 22010)., The overall extended emission of MM2 is perpendicular to the direction of outflow A (Sánnchez-Monge 2010).
 We fitted in the we-plane an elliptical Gaussian to MM2 and obtained residual emission at the position of MM2-S. indicating an additional point. source.," We fitted in the $uv$ -plane an elliptical Gaussian to MM2 and obtained residual emission at the position of MM2-S, indicating an additional point source."
 The coordinates determined for MM2-S are. (J2000): 22:21:26.807. 63:51:37.14. and the flux density is 18.2+0.8 my. which corresponds to a mass of 0.1-0.6M... assuming a dust temperature of 10-30 K. a gas-to-dust mass ratio of 100. and a dust mass opacity coefficient at 1.3 mm of 0.899 απ’ σσ”! (agglomerated grains with thin ice mantles for densities ~10° cm. Ossenkopf Henning 1994).," The coordinates determined for MM2-S are (J2000): 22:21:26.807, 63:51:37.14, and the flux density is $18.2\pm0.8$ mJy, which corresponds to a mass of 0.1–0.6, assuming a dust temperature of 10–30 K, a gas-to-dust mass ratio of 100, and a dust mass opacity coefficient at 1.3 mm of 0.899 $^2$ $^{-1}$ (agglomerated grains with thin ice mantles for densities $\sim10^6$ $^{-3}$, Ossenkopf Henning 1994)."
 The uncertainty in the masses is estimated to be a factor of 2., The uncertainty in the masses is estimated to be a factor of 2.
 As for 122198-MM2. the deconvolved size is 500x300 AU at PA.=-357. the peak intensity and flux density are 91.8+0.9beam!.. and 246+3 mJy. and the mass is ~1M: ((Table 1).," As for I22198-MM2, the deconvolved size is $500\times300$ AU at $=-35^\circ$, the peak intensity and flux density are $91.8\pm0.9$, and $246\pm3$ mJy, and the mass is $\sim1$ (Table 1)."
 Towards A5142 the millimeter emission Is dominated by two partially extended and strong sources. MMI and MM2. which are surrounded by five faint point-like sources (Palaual... in prep.).," Towards A5142 the millimeter emission is dominated by two partially extended and strong sources, MM1 and MM2, which are surrounded by five faint point-like sources (Palau, in prep.)."
 The deconvolved sizes (from elliptical Gaussian fits in the we-plane) are 1200x900 AU at P.A.=-86° for MMI and 1000x400 AU at ΡΑ.Ξ+18° for MM2., The deconvolved sizes (from elliptical Gaussian fits in the $uv$ -plane) are $1200\times900$ AU at $=-86^\circ$ for MM1 and $1000\times400$ AU at $=+18^\circ$ for MM2.
 The peak intensities and flux densities are 38+3!.. 212€7 mJy for MMI. and 62x3!.. 151+4 mJy for MM2. yielding masses of ~4 ((Table 1).," The peak intensities and flux densities are $38\pm3$, $212\pm7$ mJy for MM1, and $62\pm3$, $151\pm4$ mJy for MM2, yielding masses of $\sim4$ (Table 1)."
" Regarding the CO((2-1) emission. we first caution that an important part of the emission is filtered out by the interferometer and we are only sensitive to compact knots. even after tapering the data to a final beam of 0.6""."," Regarding the (2–1) emission, we first caution that an important part of the emission is filtered out by the interferometer and we are only sensitive to compact knots, even after tapering the data to a final beam of $0.6''$."
 In 122198 we detected chains of knots possibly tracing the cavity walls of outflows A and B (Sánnchez-Monge 22010)., In I22198 we detected chains of knots possibly tracing the cavity walls of outflows A and B (Sánnchez-Monge 2010).
 As for AS142. the CO((2-1) emission is again very clumpy but showing chains of knots which match well the known outflows of the region ZZhang 22007).," As for A5142, the (2–1) emission is again very clumpy but showing chains of knots which match well the known outflows of the region Zhang 2007)."
 In Figure 3-left we show the PdBI spectrum towards 122198-MM2. A5142-MMI and A5142-MM2 for the entire observed bandwidth.," In Figure 3-left we show the PdBI spectrum towards I22198-MM2, A5142-MM1 and A5142-MM2 for the entire observed bandwidth."
 Line identification was. performed following the methodology described below., Line identification was performed following the methodology described below.
 First. we searched for molecules with >5 transitions detectable within the observed frequency range and compared preliminar synthetic spectra (see below) for these molecules to the observed spectra.," First, we searched for molecules with $\geq5$ transitions detectable within the observed frequency range and compared preliminar synthetic spectra (see below) for these molecules to the observed spectra."
 We found that. for the three sources. aand aalone couldaccount for about half of the detected transitions (Fig.," We found that, for the three sources, and alone couldaccount for about half of the detected transitions (Fig."
 3-left)., 3-left).
 Second. we computed rotational diagrams (Fig.," Second, we computed rotational diagrams (Fig."
 4) with the ttransitions. allowing us to determine the gas temperature," 4) with the transitions, allowing us to determine the gas temperature"
"spectrum of similar shape as the data, the uncertainty levels of the simulation can be used to estimate the uncertainty of each histogram bin of the data.","spectrum of similar shape as the data, the uncertainty levels of the simulation can be used to estimate the uncertainty of each histogram bin of the data."
" The standard deviation in the simulation in each bin was on the order of the Poisson statistic, justifying the Poisson approximation used in Sect. 3.."," The standard deviation in the simulation in each bin was on the order of the Poisson statistic, justifying the Poisson approximation used in Sect. \ref{sec:ana_all}."
" The distribution of the simulated lightcurve describes the measured histogram almost as well as the best-fit Gaussian (y?—216 in 84 bins,see Fig. 2))."," The distribution of the simulated lightcurve describes the measured histogram almost as well as the best-fit Gaussian $\chi^2 = 216$ in 84 bins,see Fig. \ref{fig:lcratehistall.ps}) )."
" We presented the first analysis of the statistical distribution of the brightness of Vela X-1, with special regard to the flaring behavior."," We presented the first analysis of the statistical distribution of the brightness of Vela X-1, with special regard to the flaring behavior."
 Our main results are: From our results obtained by analyzing 3.6MMsec of ISGRI data we can calculate the probability to measure extremely bright flares in that time range., Our main results are: From our results obtained by analyzing Msec of ISGRI data we can calculate the probability to measure extremely bright flares in that time range.
 Extremely bright flares are defined by their countrate being larger than ccountsssec! in the 20-60kkeV band., Extremely bright flares are defined by their countrate being larger than $^{-1}$ in the keV band.
" The calculated probability is ~0.011% for the best-fit Gaussian, while it is ~0.023% for the simulated lightcurve."," The calculated probability is $\sim$ for the best-fit Gaussian, while it is $\sim$ for the simulated lightcurve."
 These numbers agree well with ~0.02% measured directly from our dataset., These numbers agree well with $\sim$ measured directly from our dataset.
" We conclude that bright flares are rare, but not singular events."," We conclude that bright flares are rare, but not singular events."
 No separate process is required for their explanation., No separate process is required for their explanation.
" The origin of the constant flaring behavior lies in the accretion flow onto the neutron star, which is obviously not smooth, but highly structured."," The origin of the constant flaring behavior lies in the accretion flow onto the neutron star, which is obviously not smooth, but highly structured."
" Neglecting any absorption and scattering effects, the luminosity distribution gives us the opportunity of calculating the mass distribution in this flow in units of mass per time."," Neglecting any absorption and scattering effects, the luminosity distribution gives us the opportunity of calculating the mass distribution in this flow in units of mass per time."
 To calculate the mass accretion rate it is necessary to know the absolute luminosity of the X-ray source in the given energy band., To calculate the mass accretion rate it is necessary to know the absolute luminosity of the X-ray source in the given energy band.
" As the luminosity is strongly energy-dependent, we modeled the energy spectrum with a power law with Dirac-cutoff, using the parameters of Kreykenbohmetal.(2008),, who found these parameters for data from revolutions 137-141, which were part of our analysis."," As the luminosity is strongly energy-dependent, we modeled the energy spectrum with a power law with Fermi-Dirac-cutoff, using the parameters of \citet{kreykenbohm08a}, who found these parameters for data from revolutions 137–141, which were part of our analysis."
" The overall spectrum of our data could be described with these values, too."," The overall spectrum of our data could be described with these values, too."
" With those parameters and a distance of kkpc (Nagaseetal.,1986) the median absolute luminosity (Lx) is 5.1x1026 ssec"".", With those parameters and a distance of kpc \citep{nagase86a} the median absolute luminosity $\left<L_\text{x}\right>$ is $ 5.1\times10^{36}$ $^{-1} $.
 During the accretion process only part of the potential energy is converted to X-rays., During the accretion process only part of the potential energy is converted to X-rays.
" We assume an accretion efficiency of 7=0.3 and calculate the median accretion rate The multiplicative standard deviation is &,,—1.9 in the kkeV energy range.", We assume an accretion efficiency of $\eta = 0.3$ and calculate the median accretion rate The multiplicative standard deviation is $\tilde \sigma_{\dot M} = 1.9$ in the keV energy range.
 Figure 6 shows the distribution of inferred M values and the corresponding luminosity., Figure \ref{fig:histo_mdost.eps} shows the distribution of inferred $\dot M$ values and the corresponding luminosity.
" To show the log-normal behavior of the distribution more clearly, we used a linear scaled x-axis."," To show the log-normal behavior of the distribution more clearly, we used a linear scaled $x$ -axis."
" We included data not only from the kkeV energy band, but also from the softer kkeV band and the harder kkeV band in this figure."," We included data not only from the keV energy band, but also from the softer keV band and the harder keV band in this figure."
 All three bands show the same behavior., All three bands show the same behavior.
" Slight deviations between the different energy bands are due to the rough spectral fit, because only the limited energy range of ISGRI was available for fitting."," Slight deviations between the different energy bands are due to the rough spectral fit, because only the limited energy range of ISGRI was available for fitting."
 Nonetheless it is evident that all three curves follow a log-normal distribution in a range of M very well expected for this kind of object regarding the mass loss rate of the optical companion and the assumed accretion efficiency., Nonetheless it is evident that all three curves follow a log-normal distribution in a range of $\dot{M}$ very well expected for this kind of object regarding the mass loss rate of the optical companion and the assumed accretion efficiency.
" If we assume that the neutron stars accretes directly from the wind, we can infer the mass distribution in the stellar wind from the accretion rate."," If we assume that the neutron stars accretes directly from the wind, we can infer the mass distribution in the stellar wind from the accretion rate."
" Strong density variations in the wind comparable to the variations in the accretion rate can be described by a model of a clumped stellar wind, a model based on instabilities in the the line-driven acceleration mechanism (see,e.g.,Feldmeieretal.,2003;Dessart&Owocki,2005; 2007).."," Strong density variations in the wind comparable to the variations in the accretion rate can be described by a model of a clumped stellar wind, a model based on instabilities in the the line-driven acceleration mechanism \citep[see, e.g.,][]{feldmeier03a, dessart05a, oskinova07a}. ."
 A clumpy wind is also supported by observations., A clumpy wind is also supported by observations.
" Based on data,"," Based on data,"
(see. e.g.. De Marco et 2005 and Mathieu Geller 2009).,"(see, e.g., De Marco et 2005 and Mathieu Geller 2009)."
 While. therefore. the high rates of rotation may make mass measurements more difficult. it is also possible that they are providing a signal (hat mass (transfer did occur.," While, therefore, the high rates of rotation may make mass measurements more difficult, it is also possible that they are providing a signal that mass transfer did occur."
 The advent of a new observing capability has almost always led to unanticipated discoveries., The advent of a new observing capability has almost always led to unanticipated discoveries.
 The two binaries. ΙΟΤΕ and NOLS]. may be examples.," The two binaries, KOI-74 and KOI-81, may be examples."
 Nevertheless. we mar cuestion whether the evolutionary models presented above are expected to occur commonly enough thatAepler should have been able to discover these interesting binary svstems by monitoring only ~150.000 stars.," Nevertheless, we may question whether the evolutionary models presented above are expected to occur commonly enough that should have been able to discover these interesting binary systems by monitoring only $\sim 150,000$ stars."
 We therefore conducted a first-principles study of binaries in a stellar population to predict the numbers of transiting svstems we expect (o be comprised of a main sequence star orbited by a white dwarf (hat has emerged [rom an episode of mass transfer., We therefore conducted a first-principles study of binaries in a stellar population to predict the numbers of transiting systems we expect to be comprised of a main sequence star orbited by a white dwarf that has emerged from an episode of mass transfer.
 Although these calculations were suggested bv the discoveries of IXOI-74 and they do not rely on the interpretation of these svstems.," Although these calculations were suggested by the discoveries of KOI-74 and KOI-81, they do not rely on the interpretation of these systems."
 We note that the evolution of interacting binaries consisting of a main sequence accretor and a subgiant or giant. donor is complex and involves a wide range of physical processes., We note that the evolution of interacting binaries consisting of a main sequence accretor and a subgiant or giant donor is complex and involves a wide range of physical processes.
 There are uncertainties. including the results of common envelope evolution. the fraction of incoming matter that can be retained by the aceretor. and (he angular momentum evolution of the svstem.," There are uncertainties, including the results of common envelope evolution, the fraction of incoming matter that can be retained by the accretor, and the angular momentum evolution of the system."
 Nevertheless. by paranmeterizing the effects of these processes we are able to derive a robust conclusion.," Nevertheless, by parameterizing the effects of these processes we are able to derive a robust conclusion."
 TheAepler team selected targets [rom roughly half a million stars in its field brighter than 16th magnitude., The team selected targets from roughly half a million stars in its field brighter than 16th magnitude.
 More than 90% of the targets were selected based on signal-to-noise considerations that suggested the possibility of detecüng terrestrial-size planets., More than $90\%$ of the targets were selected based on signal-to-noise considerations that suggested the possibility of detecting terrestrial-size planets.
 While a small fraction of the targets are selected (o pursue a range of other science opportunities. including.," While a small fraction of the targets are selected to pursue a range of other science opportunities, including."
 e.g.. eclipsing binaries ancl high-proper-motion stars. the majority of the targets ( 90.000) are G-tvpe stars on or near (he main sequence (Batalha et 2010).," e.g., eclipsing binaries and high-proper-motion stars, the majority of the targets $\sim 90,000$ ) are G-type stars on or near the main sequence (Batalha et 2010)."
 The presence of IKOI-T4 and IXOI-31 illustrates the presence of more massive main-secuence stars as well., The presence of KOI-74 and KOI-81 illustrates the presence of more massive main-sequence stars as well.
 In (he calculations described below. we compute (he fraction of monitored stars Gansited bv dwarls per year.," In the calculations described below, we compute the fraction of monitored stars transited by s per year."
 The number of systems in which (transits can be detected is the product of this faction and the number of monitored stars with high enough signal-to-noise that the Gransits of dwarls can be detected., The number of systems in which s can be ed is the product of this fraction and the number of monitored stars with high enough signal-to-noise that the s of s can be ed.
 Massive cwarls have radii comparable to the radius of the Earth., Massive s have radii comparable to the radius of the Earth.
 Transits of many of the {target stars by objects of this size should be detectable., Transits of many of the target stars by objects of this size should be detectable.
 The rracdius increases with decreasing mass (see., The radius increases with decreasing mass (see.
 e.g.. Parsons et 2010). so the less massive dwarls which are expected to be common among pproducts should also produce detectable transits when (hey pass in front of 90% of (targets.," e.g., Parsons et 2010), so the less massive s which are expected to be common among products should also produce able s when they pass in front of $90\%$ of targets."
 White chwarls that have not vet had a chance to cool are even larger., White dwarfs that have not yet had a chance to cool are even larger.
 We therefore, We therefore
55C [lux is stronger than the IC/CMD flux in 5-avs.,SSC flux is stronger than the IC/CMB flux in $\gamma$ -rays.
 Note that the IC/CMD his is almost in the Thomson regime. while the SSC flux is largely affected by the Ixlein-Nishina effect.," Note that the IC/CMB flux is almost in the Thomson regime, while the SSC flux is largely affected by the Klein-Nishina effect."
" We fit the data with the parameters jj=0.005. 544,=7-0x107. 54,=6.0xLO”. min=LOx107. p,= 1.5. and po=2.5."," We fit the data with the parameters $\eta = 0.005$ , $\gamma_{\rm max} = 7.0 \times 10^9$ , $\gamma_{\rm b} = 6.0 \times 10^5$ , $\gamma_{\rm min} = 1.0 \times 10^2$, $p_1 = 1.5$ , and $p_2 = 2.5$."
 The fraction parameter 7 governs the absolute values of the fluxes and the flux ratio of the inverse Compton scattering to the svuchrotvon radiation. as discussed in more detail in section ??..," The fraction parameter $\eta$ governs the absolute values of the fluxes and the flux ratio of the inverse Compton scattering to the synchrotron radiation, as discussed in more detail in section \ref{explanation}."
 The NC model derived 6«1 [rom the viewpoint of (he curent dynamical structure of the Crab Nebula. while we determine 7«1 from the viewpoint of the spectral evolution.," The KC model derived $\sigma \ll 1$ from the viewpoint of the current dynamical structure of the Crab Nebula, while we determine $\eta \ll 1$ from the viewpoint of the spectral evolution."
" The parameters 5,45. th. Pp. and ps are lixed to reproduce the observed svuchrotvon spectral shape. such as the spectral breaks and the photon indices. while 7,4, should be regarded as an upper limit to reproduce (he radio Εαν at the lowest frequency."," The parameters $\gamma_{\rm max}$, $\gamma_{\rm b}$, $p_1$, and $p_2$ are fixed to reproduce the observed synchrotron spectral shape, such as the spectral breaks and the photon indices, while $\gamma_{\rm min}$ should be regarded as an upper limit to reproduce the radio flux at the lowest frequency."
 In section ??.. these fitted parameters characterizing particle injection are discussed in detail.," In section \ref{other}, these fitted parameters characterizing particle injection are discussed in detail."
 In our caleulation. the current magnetic field strength of the Crab Nebula turns out to be Buoy=85j(G. which is smaller than ~3005 used by Atovan&Aharonian(1996).," In our calculation, the current magnetic field strength of the Crab Nebula turns out to be $B_{\rm now} = 85\mu \rm{G}$, which is smaller than $\sim 300 \mu \rm{G}$ used by \citet{aa96}."
.. This difference of the magnetic field strength can be explained as follows., This difference of the magnetic field strength can be explained as follows.
 Atovan& adopted Byeceὀθθμα from the NC model and adjusted the particle number to reproduce the observations., \citet{aa96} adopted $B_{\rm KC} \sim 300 \mu \rm G$ from the KC model and adjusted the particle number to reproduce the observations.
 Thev applied roughly hall a spin-down power compared with the NC model to reproduce the spectrum and (hus the other half is missing., They applied roughly half a spin-down power compared with the KC model to reproduce the spectrum and thus the other half is missing.
 On theother hand. all the injected spin-down power is divided between the magnetic field and the particle energies in our model.," On theother hand, all the injected spin-down power is divided between the magnetic field and the particle energies in our model."
 If. sve adopt Buoy=Dyee300jG. the svnehrotron flux and also the SSC flux increase by about an order of magnitude.," If we adopt $B_{\rm now} = B_{\rm KC} \sim 300 \mu \rm G$, the synchrotron flux and also the SSC flux increase by about an order of magnitude."
" Note that the relativistic MIID simulation byVolpietal.(2008) alsoindicatesasmaller value of the spatially averaged magnetic field strength  100jC. which is close to our value D,= 85jG."," Note that the relativistic MHD simulation by\citet{vet08} alsoindicatesasmaller value of the spatially averaged magnetic field strength $ \sim 100 \mu G$ , which is close to our value $B_{\rm now} = 85\mu \rm{G}$ ."
velocity as the ejecta expand.,velocity as the ejecta expand.
 Although the 800 nm [feature fades with the approach to maxinmumn elt. i always remains distinct [rom the photospheric Ca IL IR Giplet feature as (he latter gains in strength.," Although the 800 nm feature fades with the approach to maximum light, it always remains distinct from the photospheric Ca II IR triplet feature as the latter gains in strength."
 The 800 nm [feature does not seem to evolve substantially in velocity space., The 800 nm feature does not seem to evolve substantially in velocity space.
 In addition. the post-maxinnun pholospheric component of the Ca II Ih triplet in SN 2001el has a sharp blue edge at velocity around -12.500 ((see the last (wo spectra at +16 and +38 d in Figure 2).," In addition, the post-maximum photospheric component of the Ca II IR triplet in SN 2001el has a sharp blue edge at velocity around -12,500 (see the last two spectra at +16 and +38 d in Figure 2)."
 This implies (that there is a sharp density drop of Ca at that velocity., This implies that there is a sharp density drop of Ca at that velocity.
 On the other haud. (he pre-inaximunm line profiles of the 500 nm feature show rather sharp edges (hat. if identified with Ca IL. correspond to a velocity of about 19.000 20.000 oon the red side and about 26.000 oon the blue side (see the top spectrum in Figure 2).," On the other hand, the pre-maximum line profiles of the 800 nm feature show rather sharp edges that, if identified with Ca II, correspond to a velocity of about 15,000 – 20,000 on the red side and about 26,000 on the blue side (see the top spectrum in Figure 2)."
 The polarization data accentuates this delineation— in velocity space as shown in the first panels corresponding to the data on 26 sept in Figure 3., The polarization data accentuates this delineation in velocity space as shown in the first panels corresponding to the data on 26 Sept in Figure 3.
 In anv case. the red edge of the high velocity feature al the top of Figure 2 does not overlap with the blue edge of the low velocity Ca II feature al the bottom of Figure 2.," In any case, the red edge of the high velocity feature at the top of Figure 2 does not overlap with the blue edge of the low velocity Ca II feature at the bottom of Figure 2."
 If both of these absorptions are due to the Ca II UR triplet. this implies that the high-velocity filament or shell also has rather well-defined geometrical boundaries.," If both of these absorptions are due to the Ca II IR triplet, this implies that the high-velocity filament or shell also has rather well-defined geometrical boundaries."
 The polarization data suggest that the lowest velocity matter in this hieh-velocity [eature might be at about 17.000. if the feature is Ca LL," The polarization data suggest that the lowest velocity matter in this high-velocity feature might be at about 17,000 if the feature is Ca II."
 The polarization shows that this feature also has a different geometrical orientation than the geometry that defines the dominant axis of the photosphere., The polarization shows that this feature also has a different geometrical orientation than the geometry that defines the dominant axis of the photosphere.
 Taken together. (he velocity separation. (he large amplitude of the polarization. aud the different polarization augle all imply (hat (his feature is a kinematically and geometrically separate hieh-velocitv component (hat is enriched in caleiuni.," Taken together, the velocity separation, the large amplitude of the polarization, and the different polarization angle all imply that this feature is a kinematically and geometrically separate high-velocity component that is enriched in calcium."
 It is difficult to see whether something like (his separate high velocity feature exists in other SN Ia. The kinematic boundaries of the hieh-velocity calcium are impossible to discern in the data of SN 1994D. since the feature is so much weaker.," It is difficult to see whether something like this separate high velocity feature exists in other SN Ia. The kinematic boundaries of the high-velocity calcium are impossible to discern in the data of SN 1994D, since the feature is so much weaker."
" A search of published SN Ia spectra show that most of the SN Ia with pre-naximnunm spectra covering the 800 nm area reveal weak spectral features similar to that of SN 1994D. As for SN 1994D. however. it is difficult to discern whether this high velocity component is geometrically ""detached"" from (he photospheric structure when the line is weak."," A search of published SN Ia spectra show that most of the SN Ia with pre-maximum spectra covering the 800 nm area reveal weak spectral features similar to that of SN 1994D. As for SN 1994D, however, it is difficult to discern whether this high velocity component is geometrically “detached"" from the photospheric structure when the line is weak."
 In addition. the presence of Fe HH absorption al about 800 nm can obscure (he nature of this feature.," In addition, the presence of Fe II absorption at about 800 nm can obscure the nature of this feature."
 The curent SN Ia data base does seeni lo suggest that the verv strong feature. definitely clisplaced from (he photospheric Ca 1I IH. triplet. is rather special to SN 2001el.," The current SN Ia data base does seem to suggest that the very strong feature, definitely displaced from the photospheric Ca II IR triplet, is rather special to SN 2001el."
 We address the possible physical origin of this leature in 85., We address the possible physical origin of this feature in 5.
appears to be dominated by the jet svnchrotron spectrum: in particular the spectral turn-over between self-absorbed and optically thin svnehrotron emission has been detected in the near-infared in GX 4 (Corbel Fender 2002).,appears to be dominated by the jet synchrotron spectrum; in particular the spectral turn-over between self-absorbed and optically thin synchrotron emission has been detected in the near-infared in GX $-$ 4 (Corbel Fender 2002).
 ]nstead. the infrared. emission of NTE 318 during these epochs appears to be dominated by the accretion disc and/or irradiated companion star in agreement with the disc instability model (e.g. Lasota 2001).," Instead, the infrared emission of XTE $-$ 318 during these epochs appears to be dominated by the accretion disc and/or irradiated companion star in agreement with the disc instability model (e.g. Lasota 2001)."
 Vhe X-rav/radio behaviour over the course of the outburst bears a number of similarities to NTIS J1859|226., The X-ray/radio behaviour over the course of the outburst bears a number of similarities to XTE J1859+226.
" Following the initial lowhard state. which in both sources asted ~ days. the radio and soft) X-rays reached. a ocak, quasi-simultaneousky: in the case of NTE J1859|226 here was a hard X-rav peakprior to this."," Following the initial low/hard state, which in both sources lasted $\sim$ days, the radio and soft X-rays reached a peak quasi-simultaneously; in the case of XTE J1859+226 there was a hard X-ray peak to this."
 Similarly ATE sali318 reaches a soft. X-ray peak approximately coincide (or just before) the radio peak: indeed Fender. Belloni Gallo(in prep.)," Similarly XTE $-$ 318 reaches a soft X-ray peak approximately coincidentally (or just before) the radio peak; indeed Fender, Belloni Gallo (in prep.)"
 suggest that the radio peak and soft X-ray peak are closely related (see also Corbel et al., suggest that the radio peak and soft X-ray peak are closely related (see also Corbel et al.
 2004 who reach a similar conclusion)., 2004 who reach a similar conclusion).
 The steady jet of the initial ονπαν state is thought to persist through the period. of spectral softening and give way to an optically thin racio ejection just after the soft X-ray peak., The steady jet of the initial low-hard state is thought to persist through the period of spectral softening and give way to an optically thin radio ejection just after the soft X-ray peak.
 After the initial maximum of NTIS J1859|226. it was the hard N-ravs that were more closely linked with the racio behaviour.," After the initial maximum of XTE J1859+226, it was the hard X-rays that were more closely linked with the radio behaviour."
 As the X-ray source decaved. from. outburst it underwent a series of temporary hardenings. superimposed on a general softening. and cach of these hardenings was associated with a new radio ejection. (Brocksopp οἱ al.," As the X-ray source decayed from outburst it underwent a series of temporary hardenings, superimposed on a general softening, and each of these hardenings was associated with a new radio ejection (Brocksopp et al."
 2002)., 2002).
 While the S/N of the hardness ratio is poor. XTE 318 also appears to show some level of both intensity and spectral variability during its decay. including least two glitches.," While the S/N of the hardness ratio is poor, XTE $-$ 318 also appears to show some level of both intensity and spectral variability during its decay, including at least two glitches."
 Comparison with other sources such as NPE J1859|226. GRO 40 and NTE 264 (Brocksopp et al.," Comparison with other sources such as XTE J1859+226, GRO $-$ 40 and XTE $-$ 264 (Brocksopp et al."
 2002 and references therein) might suggest that simultaneous radio ejections would have been expected., 2002 and references therein) might suggest that simultaneous radio ejections would have been expected.
 Given that both elitehes and radio ejections are associated with temporary hardenings in these other sources dU is tempting to suggest that they are all different manifestations ofa jet event., Given that both glitches and radio ejections are associated with temporary hardenings in these other sources it is tempting to suggest that they are all different manifestations of a jet event.
 Unfortunately the NTE 318 data co not allow us to confirm this. perhaps due to inadequate S/N and/or time resolution.," Unfortunately the XTE $-$ 318 data do not allow us to confirm this, perhaps due to inadequate S/N and/or time resolution."
 However the data of Nagata et al. (, However the data of Nagata et al. (
2003). do indeed show simultaneous X-rav/infrared events which coincide with radio non-detections.,2003) do indeed show simultaneous X-ray/infrared events which coincide with radio non-detections.
 Again. this may be due to the S/N and time-resolution or it may confirm. previous suggestions that some elitches may. be events taking place in the accretion disc (e.g. Lasota 2001) and. independently of the jet.," Again, this may be due to the S/N and time-resolution or it may confirm previous suggestions that some glitches may be events taking place in the accretion disc (e.g. Lasota 2001) and independently of the jet."
 We also note that. contrary NTE 31s. the three sources. listed above were in the very high state (or steep power-law state) at. the time of the simultaneous elitehes/ejections: it is probable that the spectral state of the accretion disc is significant in determining whether or not ejections take place (e.g. Fender. Belloni Gallo in prep.).," We also note that, contrary to XTE $-$ 318, the three sources listed above were in the very high state (or steep power-law state) at the time of the simultaneous glitches/ejections; it is probable that the spectral state of the accretion disc is significant in determining whether or not ejections take place (e.g. Fender, Belloni Gallo in prep.)."
 We have presented radio observations of the 2003 outburst of the X-ray transient NTE 318., We have presented radio observations of the 2003 outburst of the X-ray transient XTE $-$ 318.
 Phe radio source was unresolved. ancl reached a peak of ~5 my., The radio source was unresolved and reached a peak of $\sim 5$ mJy.
 Study. of the spectral index showed that. while the source was optically thin throughout. there was some significant variability which we interpret as partially scllabsorbecl emission at the onset of two cliscrete ejection events.," Study of the spectral index showed that, while the source was optically thin throughout, there was some significant variability which we interpret as partially self-absorbed emission at the onset of two discrete ejection events."
 Following a period of non-detection. the radio source switched on again contemporancously with the transition to the low/hard state.," Following a period of non-detection, the radio source switched on again contemporaneously with the transition to the low/hard state."
 The broadband spectrum showed that the infrared. emission was significantly brighter than the racio svnchrotron spectrum: we suggest that the infrared emission was dominated bv the accretion disc and/or irracdation of the companion star as expected for a csolt” X-ray transient event., The broadband spectrum showed that the infrared emission was significantly brighter than the radio synchrotron spectrum; we suggest that the infrared emission was dominated by the accretion disc and/or irradiation of the companion star as expected for a “soft” X-ray transient event.
 The variability curing the decay was reminiscent. of that of NTI 1550|226: in the case of NPE 11859|226 this variability was interpreted as a sequence of jet ejections., The variability during the decay was reminiscent of that of XTE J1859+226; in the case of XTE J1859+226 this variability was interpreted as a sequence of jet ejections.
 This emphasizes the need. for high S/N X-ray hardness and racio monitoring during the ~elitches” which are often observed superimposed on the decay of the X-ray source., This emphasizes the need for high S/N X-ray hardness and radio monitoring during the “glitches” which are often observed superimposed on the decay of the X-ray source.
 We are very grateful to Jean Swank. who kindly shared some RANTE/PCA results with us prior to publication.," We are very grateful to Jean Swank, who kindly shared some /PCA results with us prior to publication."
 We are also eratelul for the quick-look results provided by the IXTE/NSM toam., We are also grateful for the quick-look results provided by the /ASM team.
 Phe Australia Telescope is funded. by the Commonwealth of Australia for operation as a National Facility managed by CSIRO., The Australia Telescope is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO.
 DCI is a Fellow of the Finnish Academy., DCH is a Fellow of the Finnish Academy.
of the dillerent “wing” sizes present in real images. in the next section we have modelled the Molfat. PSEs using three cilferent values of 2j: j—5 (to simulate the turbulence prediction). 2.5 (the default value of the LRAL package). and 1.5 (to model a large “wine” in the PSI).,"of the different “wing” sizes present in real images, in the next section we have modelled the Moffat PSFs using three different values of $\beta$ : $\beta$ =5 (to simulate the turbulence prediction), 2.5 (the default value of the IRAF package), and 1.5 (to model a large “wing” in the PSF)."
 The equations that we have shown in the previous section are general results for Mollat seeing., The equations that we have shown in the previous section are general results for Moffat seeing.
 For practical purposes with applications to real galaxies. we are going to focus on the Sérrsic profile.," For practical purposes with applications to real galaxies, we are going to focus on the Sérrsic profile."
" In the particular case of i"". proliles. the surface brightness. distribution is given (in elliptical coordinates) by: where Z(0) is the central intensity anc ο the cllective radius of the profile."," In the particular case of $r^{1/n}$ profiles, the surface brightness distribution is given (in elliptical coordinates) by: where $I(0)$ is the central intensity and $r_{\rm e}$ the effective radius of the profile."
" The constant b, is chosen such that half the total luminosity predicted by the law comes from £<r.", The constant $b_n$ is chosen such that half the total luminosity predicted by the law comes from $\xi<r_{\rm e}$.
 ὃν can be well approximated by the relation θ.δόδη0.142 (Caon et al., $b_n$ can be well approximated by the relation $b_n=0.868n-0.142$ (Caon et al.
 1993)., 1993).
 To studs the effect. of seeing on the central intensity we make use of Eqs. (, To study the effect of seeing on the central intensity we make use of Eqs. (
8) and (14).,8) and (14).
 Figure 3 shows this effect for different values of the ellipticity in the intrinsic light profile. assuming 7 = 1.5. 2.5. 5 and for the Gaussian case (ic. 3 cox as well.," Figure 3 shows this effect for different values of the ellipticity in the intrinsic light profile, assuming $\beta$ = 1.5, 2.5, 5 and for the Gaussian case (i.e. $\beta \to \infty$ ) as well."
" na As 3 increases (Le. as the size of the PSE ""wings? decrease) we asymptotically recover the ellect on the central intensity produced by a Gaussian PSE.", na As $\beta$ increases (i.e. as the size of the PSF “wings” decrease) we asymptotically recover the effect on the central intensity produced by a Gaussian PSF.
 From this figure it follows that the ellect of seeing on the central intensity increases as the size of the “wings” increase., From this figure it follows that the effect of seeing on the central intensity increases as the size of the “wings” increase.
 This is easily understood. as broader “wines” increase the probability that a photon will hit the imaging device at a point further olfset from. where it would have hit in the absence of a seeing., This is easily understood as broader “wings” increase the probability that a photon will hit the imaging device at a point further offset from where it would have hit in the absence of a seeing.
 For a typical value of 7 (e.g. 3 = 2.5). the dillerence [rom a Caussian PSE is ~LOM%..," For a typical value of $\beta$ (e.g. $\beta$ = 2.5), the difference from a Gaussian PSF is $\sim$."
 For a given secingdisc tthe relative size of the EWHLIAL to etlective radius along the xaxis of Fig., For a given seeing–disc the relative size of the FWHM to effective radius along the x–axis of Fig.
 3. FFig.," 3, Fig."
 4 and Fig., 4 and Fig.
 5). the central intensity of profiles with [larger values of η is more alfected than for low n as is expected because of the higher central light concentration of these profiles.," 5), the central intensity of profiles with larger values of $n$ is more affected than for low $n$ as is expected because of the higher central light concentration of these profiles."
 As noted in Section 2.2. the elfect of secing on the central intensity of the object is also dependent. on the intrinsic ellipticitv of the object: the central intensity of galaxies with larger ellipticities are more alfected by seeing.," As noted in Section 2.2, the effect of seeing on the central intensity of the object is also dependent on the intrinsic ellipticity of the object: the central intensity of galaxies with larger ellipticities are more affected by seeing."
" The seeing cllect on elfective radius can be obtained by solving for ro from the conservation of luminosity by the convolution Lo το). where Lr) is the luminosity of the source inside r, and £°0r5) is the luminosity. obtained [rom the object alfected by secing. measured. inside its ellective radius."," The seeing effect on effective radius can be obtained by solving for $r_e^c$ from the conservation of luminosity by the convolution $L^{\rm c}(r_{\rm
e}^{\rm c})$ $L(r_{\rm e})$, where $L(r_e)$ is the luminosity of the source inside $r_{\rm e}$ and $L^{\rm c}(r_{\rm e}^{\rm c})$ is the luminosity obtained from the object affected by seeing, measured inside its effective radius."
 Figure 4 shows this effect for cüllerent values of ellipticitv. with ;3 = 1.5. 2.5. 5 and for the Gaussian PSE.," Figure 4 shows this effect for different values of ellipticity, with $\beta$ = 1.5, 2.5, 5 and for the Gaussian PSF."
 Here. the presence of significant “wings” in the PSE produces an effective radius larger than what is expected from Gaussian secing.," Here, the presence of significant “wings” in the PSF produces an effective radius larger than what is expected from Gaussian seeing."
 As n increases theconmvotlved elective radius also increases., As $n$ increases the effective radius also increases.
 “Phe ellipticity cllect is also shown., The ellipticity effect is also shown.
 Greater ellipticities result in greater effective radii. and these dilferences are more important for greater values of n.," Greater ellipticities result in greater effective radii, and these differences are more important for greater values of $n$ ."
 This result is as expected. due to the diminution of the central intensity with Iarger ellipticity., This result is as expected due to the diminution of the central intensity with larger ellipticity.
 For the values of 2 typically present in real images (2.5 «3< 4: sce Saglia et al., For the values of $\beta$ typically present in real images (2.5 $<\beta<$ 4; see Saglia et al.
 1993) we obtain deviations [rom Gaussian secing in the range (bigger deviations are obtained for smaller values of the ratio reZENTHIM)., 1993) we obtain deviations from Gaussian seeing in the range (bigger deviations are obtained for smaller values of the ratio $r_{\rm e}^{\rm c}/{\rm FWHM}$ ).
 Lt should be noted that our measurement of the effective radius has been obtained over the semi-major axis., It should be noted that our measurement of the effective radius has been obtained over the semi-major axis.
 Some authors use as racial distance the magnitude i—eh: in this case. the effective racius ofthe object allected by seeing is MNgiven by iaV/l]ς ctrz). where εί} can be obtained. using Equation (10).," Some authors use as radial distance the magnitude $r^*=\sqrt{ab}$; in this case, the effective radius ofthe object affected by seeing is given by $r_{\rm e}^{\rm c*}=
r_{\rm e}^{\rm c}\sqrt{1-\epsilon (r_{\rm e}^{\rm c})}$, where $\epsilon(r_{\rm e}^{\rm c})$ can be obtained using Equation (10)."
 ‘To quantify the ellect of seeing on the shape parameter n we use the parameter 4(£) CEO1)., To quantify the effect of seeing on the shape parameter $n$ we use the parameter $\eta(\xi)$ (T01).
 g(£) is equivalent. locally. to the parameter 2 of the Sérrsic profile.," $\eta(\xi)$ is equivalent, locally, to the parameter $n$ of the Sérrsic profile."
 his. parameter can be understood as a measure of the slope of the profile., This parameter can be understood as a measure of the slope of the profile.
 In Fig., In Fig.
 5 we show the effects of secing on this parameter (evaluated. at 15) for dillerent. sizes of the PSE “wines” ancl different intrinsic ellipticities., 5 we show the effects of seeing on this parameter (evaluated at $r_{\rm e}^{\rm c}$ ) for different sizes of the PSF “wings” and different intrinsic ellipticities.
 Lt is easy to see how the real value of n is recovered. asymptotically when the ratio rf.EWHIT increases.," It is easy to see how the real value of $n$ is recovered asymptotically when the ratio $r_{\rm e}^{\rm
c}/{\rm FWHM}$ increases."
 Ehe effect of secing on lr.) is bigger for smaller values of 2., The effect of seeing on $\eta(r_{\rm e}^{\rm c})$ is bigger for smaller values of $\beta$.
" This means that bigger PSE ""wings? produce a stronger ellect on the slope of the profile.", This means that bigger PSF “wings” produce a stronger effect on the slope of the profile.
 The increase in the intrinsic ellipticity of the profile has a similar ellect. but the influence is not as important as it was for the previous parameters.," The increase in the intrinsic ellipticity of the profile has a similar effect, but the influence is not as important as it was for the previous parameters."
 Note that seeing elfects always produce a surface brightness profile with a smaller value of n than the actual one., Note that seeing effects always produce a surface brightness profile with a smaller value of $n$ than the actual one.
 It is expected that any procedure to recover the structural parameters of the profile that does not take into account the elfect of seeing will obtain lower values of η than the actual ones., It is expected that any procedure to recover the structural parameters of the profile that does not take into account the effect of seeing will obtain lower values of $n$ than the actual ones.
 Lower values of 2 will also be expected if the intrinsic ellipticities of the objects are not taken into account during the recovery process., Lower values of $n$ will also be expected if the intrinsic ellipticities of the objects are not taken into account during the recovery process.
 Εις is crucial in the study of highz galaxies., This is crucial in the study of high–z galaxies.
" As was shown for the central intensity and effective radius. the presence of ""wings"" in the PSEcauses the values of the profile parameters to deviate from the prediction mace correcting for Gaussian secing."," As was shown for the central intensity and effective radius, the presence of “wings” in the PSFcauses the values of the profile parameters to deviate from the prediction made correcting for Gaussian seeing."
 In the case of the yO.) parameter. thesedeviations are in the range of 2054.," In the case of the $\eta(r_{\rm e}^{\rm c})$ parameter, thesedeviations are in the range of ."
.. From the previous results it is now easy to study the cllect, From the previous results it is now easy to study the effect
This quautity is bilinear in the two functions £(7) and Gr). which are completely ivbitrary. except that thev aud their derivatives are assiunied to vanish sufficiently rapidly at the limuts of integration (0.ox).,"This quantity is bilinear in the two functions $F(\nu)$ and $G(\nu)$, which are completely arbitrary, except that they and their derivatives are assumed to vanish sufficiently rapidly at the limits of integration $(0,\infty)$."
 We write the basic RD F-P result (2)) for an arbitrary fiction F(r). = Poy ο," We write the basic RD F-P result ) for an arbitrary function $F(\nu)$, ) = ) ) ]."
 Using this for the integral over v! in (A}). we have v) Poly αυ |ον.," Using this for the integral over $\nu'$ in ), we have ) ) = ) ]."
 luteeratiugC»c» by parts twice. assiumunec» zero contributions at the limits. [m]eives Fae) ο ου ddw.. where we have interchanged v<>»/ in the double iutegral ou the left hand side.," Integrating by parts twice, assuming zero contributions at the limits, gives ) ) ) ] , where we have interchanged $\nu \leftrightarrow \nu'$ in the double integral on the left hand side."
" Interchaneine F<>G vields yRO”""..v)\F \) = Μο”. ον.", Interchanging $F \leftrightarrow G$ yields ) ) = ) ].
 Subtracting this equation from CÀ)) gives This relatiou is true for arbitrary G() and Fv}. which implies. )) R(/..v)). that is. the RD F-P approximation maintains the exact svuuuectry of the redistribution function on which itwas based.," Subtracting this equation from ) gives This relation is true for arbitrary $G(\nu)$ and $F(\nu)$, which implies, ) ), that is, the RD F-P approximation maintains the exact symmetry of the redistribution function on which itwas based."
 , 	 
to Al-normal stars (the Al ratio) aud the ratio of blue to red HB stars (the HB ratio.,to Al-normal stars (the Al ratio) and the ratio of blue to red HB stars (the HB ratio).
]t is the ¢Ma rtuudalces as a [fiuction of maguituce that is critical iu determining 11e whether this correlation exists and what is cause might be., It is the of abundances as a function of magnitude that is critical in determining the whether this correlation exists and what its cause might be.
 This cur1611 work attempts o provide some fresh data ou clusters that have been Iistorically uudestudied., This current work attempts to provide some fresh data on clusters that have been historically under-studied.
 Both Mso (NCC 6093) aud (GC 67252 possess blue HBs relative to other clusters al sliilar metalicity., Both M80 (NGC 6093) and NGC 6752 possess blue HBs relative to other clusters at similar metallicity.
 For exampe. Ferraroetal.(1998). compared Hubble U.V. MSO p10tomeltry direcly with M13 aud M3. with he result tiat NSOs HB is very similar to M13. while M3 lacked he esxtended blue tail of tle oller two c1sters.," For example, \citet{Ferr98} compared Hubble $U,V$ M80 photometry directly with M13 and M3, with the result that M80's HB is very similar to M13, while M3 lacked the extended blue tail of the other two clusters."
 Meanwhile. Ciruudahlletal.(1999) presented exleusive StrOuugren photometry of NGC 0152. M13. aud M3. showing exteusive blue tails iu the Orhier two cOldpared with he latter.," Meanwhile, \citet{GCLSA1999} presented extensive Strömmgren photometry of NGC 6752, M13, and M3, showing extensive blue tails in the former two compared with the latter."
" Neiler M50 nor NGC 6722 had been studied exteusively or abuucdaik""es util Grattoretal.(2001) €and Cirüidahletal.(2002) determiuec [Al/Fe] for 39 Slars in NGC 6752 in total. exteudiug tle ca of Norris&DaCosta(1995)."," Neither M80 nor NGC 6752 had been studied extensively for abundances until \citet{GratAl2001} and \citet{GBNF2002}
 determined [Al/Fe] for 39 stars in NGC 6752 in total, extending the data of \citet{ND95}."
. These were significat1 'esults yecatise Chey probe less evolved stars near the main-sequence (πο. ο1 the subgiaut ranch. πα1 the base of the RGB. whic1 are below the point. that uiinine theories predict tliat aluilu Cal be produced.," These were significant results because they probed less evolved stars near the main-sequence turnoff, on the subgiant branch, and at the base of the RGB, which are below the point that mixing theories predict that aluminum can be produced."
 Using a non-LTE atalysis. Grattouetal.(2001) observed dwarls with [Al/Fe] as low as —0.76 dex from the AI I rexoinance lines. aud subgiauts witli [Al/Fe] as high as +0.86 dex rou the doublet a AASTT3/TLA.," Using a non-LTE analysis, \citet{GratAl2001} observed dwarfs with [Al/Fe] as low as $-0.76$ dex from the Al I resonance lines, and subgiants with [Al/Fe] as high as $+0.86$ dex from the doublet at ${\lambda}{\lambda}$ 8773/74."
 hi fact. the resonance line aialysis for the dwarls vieldedl [AL/Fe] —0.1840.15 (s.e.n.)," In fact, the resonance line analysis for the dwarfs yielded [Al/Fe] $=~-0.18~{\pm}~0.15$ (s.e.m.)"
 dex. while the subgiauts gave [Al/Fe] =+0.29+0.11 dex (s.ean.).," dex, while the subgiants gave [Al/Fe] $=~+0.29~{\pm}~0.11$ dex (s.e.m.)."
 The 'esults fkoe the cdwarls are uicertain ¢ue to the difficulty oL analyzing resouauce lines with noLLTE corrections of as mucl as +0.6 dex: vet. the results are still siuprisinely low.," The results for the dwarfs are uncertain due to the difficulty of analyzing resonance lines with non-LTE corrections of as much as $+0.6$ dex; yet, the results are still surprisingly low."
 What causes the clrastic cha199 roi tje 1nall sequence o tlie subgiant brauch: atuospheric effects. the different. choice of lines. (YE αμactual Ipsical phenomenon?," What causes the drastic change from the main sequence to the subgiant branch: atmospheric effects, the different choice of lines, or anactual physical phenomenon?"
 The Curuiclahletal.(2002)2) results are wore ln ine with lie 'esults of Norjs&DaCosta(1995). and wih other clusters., The \citet{GBNF2002} results are more in line with the results of \citet{ND95} and with other clusters.
 We discuss inplications of the altmiuuu data in NQC 6752 in inore detail iu section 6.2.. ALSO., We discuss implications of the aluminum data in NGC 6752 in more detail in section \ref{sec:final_look}.
 οι the othe: liaxl. has ix) pudlished abundances and is iu ieec of further iuvestigation. especially οἷν‘ell LS SILΠΕties to M13 1311 inetallicity and HB morphoogy.," M80, on the other hand, has no published abundances and is in need of further investigation, especially given its similarities to M13 in metallicity and HB morphology."
 The res ol the pa peris outlined acc'ordiug to tle [ollowiug: We begin with a description of the observatious In Sectio1 2.. Followed NOu data reduction techuiques 1 Section 3..," The rest of the paper is outlined according to the following: We begin with a description of the observations in Section \ref{sec:obs}, followed by our data reduction techniques in Section \ref{sec:ccd}."
 We then discuss membership criteria in Section L., We then discuss membership criteria in Section \ref{sec:rv}.
 After οιilliug the data. we show the results of our abuudauce analysis in Section 5 aud give our inal conclusions in Sectiou 6..," After culling the data, we show the results of our abundance analysis in Section \ref{sec:analysis} and give our final conclusions in Section \ref{sec:conclude}."
 Between 1999 and 2001 we used the CTIO Blauco fim telescope with the Hydra inulti-object spectrograph in the echelle mode to observe 105 stars near NGC 6752 aud £7 stars near N80., Between 1999 and 2001 we used the CTIO Blanco 4m telescope with the Hydra multi-object spectrograph in the echelle mode to observe 105 stars near NGC 6752 and 47 stars near M80.
 Iu the eud. ouly a subset of these spectra were of sufficient quality to determine both cluster membership and abundance information.," In the end, only a subset of these spectra were of sufficient quality to determine both cluster membership and abundance information."
 We cleseribe the observatious in the following two subsectious.aud discuss the radial velocity membership criteria in Section 1..," We describe the observations in the following two subsections,and discuss the radial velocity membership criteria in Section \ref{sec:rv}. ."
Observatory (180) wide band fillers at 12. 25. 60 ancl 100j/mun. (See Juraetal.(LO87) [or an early study of elliptical galaxies with ΗνΑΙ observations.,"Observatory (ISO) wide band filters at 12, 25, 60 and m. (See \citet{Jura} for an early study of elliptical galaxies with IRAS observations."
 See for a recent review of mid and far IHR emission [rom all tvpes of galaxies.), See \citet{MidIR View of Galaxies} for a recent review of mid and far IR emission from all types of galaxies.)
 In. addition to this thermal emission. broad line absorption or emission is observed in most AGB stars ad 10-12;anm and for oxvgen-rich. AGB stars. again at 20jmm (See Speckοἱal.(2000). for a recent study of AGB mid-IBR features).," In addition to this thermal emission, broad line absorption or emission is observed in most AGB stars at m and for oxygen-rich AGB stars, again at m (See \citet{Speck} for a recent study of AGB mid-IR features)."
 For older. low mass stars found in elliptical galaxies. ihe AGB star envelopes have a high oxveen content as opposed to the high carbon content present in high mass AGB envelopes. which occurs through inner laver carbon dredge up and expulsion.," For older, low mass stars found in elliptical galaxies, the AGB star envelopes have a high oxygen content as opposed to the high carbon content present in high mass AGB envelopes, which occurs through inner layer carbon dredge up and expulsion."
 The theoretical models of dust condensation show amorphous silicates (Si-O) to be the major constituent of the dust in these oxvgen-rich environments. where the stretching and bending of the Si-O bond is responsible for the 10 and 20mm absorpton/emission. (," The theoretical models of dust condensation show amorphous silicates (Si-O) to be the major constituent of the dust in these oxygen-rich environments, where the stretching and bending of the Si-O bond is responsible for the 10 and m absorption/emission. ("
Ilereafter AGB. refers to the low mass. oxveen-rich subelass of AGB stars.),"Hereafter AGB, refers to the low mass, oxygen-rich subclass of AGB stars.)"
 One of the surprises revealed by ISO are that in addition to amorphous silicates. ervstalline silicates are observed through many narrow band features from 15-45;num (Tielensetal.1993).," One of the surprises revealed by ISO are that in addition to amorphous silicates, crystalline silicates are observed through many narrow band features from m \citep{Ref A}."
. These observations have sparked renewed interest in dust particle formation in AGD environments (Sogawa&Ixozasa1999:GailSedlmayvr1999).," These observations have sparked renewed interest in dust particle formation in AGB environments \citep{Ref B,Ref C}."
. The dust composition. size distribution and other properties (albedo. dipole strength. etc) deline the interaction with the gas and ulümatelv control the observational consequences.," The dust composition, size distribution and other properties (albedo, dipole strength, etc) define the interaction with the gas and ultimately control the observational consequences."
" Consequently. much laboratory work is being devoted to ""ervowing silicate dust particles and attempting to extract similar particles from interplanetary dust particles brought to Earth through comets and meteorites 2000;ATarGin 1995).."," Consequently, much laboratory work is being devoted to “growing” silicate dust particles and attempting to extract similar particles from interplanetary dust particles brought to Earth through comets and meteorites \citep{Ref D,Ref E}."
 These theoretical and lab studies have led to a working moclel of AGB star dust formation and emission., These theoretical and lab studies have led to a working model of AGB star dust formation and emission.
 Although much work has been done on the study of individual AGB stars. population studies of these AGB dust features in earlv-tvpe galaxies is relatively unexplored due to the [aint nature of the emission involved.," Although much work has been done on the study of individual AGB stars, population studies of these AGB dust features in early-type galaxies is relatively unexplored due to the faint nature of the emission involved."
 In an elliptical galaxy. it is the sum of many individual oxvgen-rich. AGB stars that will produce in aggregate a similar feature to individual AGB stars.," In an elliptical galaxy, it is the sum of many individual oxygen-rich AGB stars that will produce in aggregate a similar feature to individual AGB stars."
 This feature can be used to confirm (his general stellar evolution picture aud determine {he mass loss rate into the entire galaxy., This feature can be used to confirm this general stellar evolution picture and determine the mass loss rate into the entire galaxy.
 An effort to detect this mid-Ilt excess in early-type ealaxies was first. carried out by Knappetal.(1992). (IARGW hereafter) who used the IRAS All Sky Survey coupled with ground based 10j/muim data to search for the dusty component of the ISM in nearby elliptical galaxies., An effort to detect this mid-IR excess in early-type galaxies was first carried out by \citet{KGWW} (KGW hereafter) who used the IRAS All Sky Survey coupled with ground based m data to search for the dusty component of the ISM in nearby elliptical galaxies.
 Because (his emission is quite faint. (hese galaxies were detected. at low signal-to-noise ratio.," Because this emission is quite faint, these galaxies were detected at low signal-to-noise ratio."
 Nevertheless. {μον observecl excess emission al 12;mnm relative to a derived stellar continuum. indicating emission from the eircumstellar dust envelope.," Nevertheless, they observed excess emission at m relative to a derived stellar continuum, indicating emission from the circumstellar dust envelope."
 KGW scaled the emission from Galactic AGB stars to their signal and they determined galaxv-wide mass loss rates of —0.7|.., KGW scaled the emission from Galactic AGB stars to their signal and they determined galaxy-wide mass loss rates of $\sim$.
. Also. thev show that the dust enission is extended on the scale of the galaxy. ruling out the hypothesis that the dust emission originates only in the nuclear regions.," Also, they show that the dust emission is extended on the scale of the galaxy, ruling out the hypothesis that the dust emission originates only in the nuclear regions."
 ILowever. many of the galaxy detections were," However, many of the galaxy detections were"
Fig.,Fig.
 [8] shows the maximum amount of memory allocated during a variety of SHTs carried out using aand the Fortran HEALPix facility., \ref{memsize} shows the maximum amount of memory allocated during a variety of SHTs carried out using and the Fortran HEALPix facility.
" Since the memory required for forward and backward transforms is essentially equal, only one direction has been plotted."," Since the memory required for forward and backward transforms is essentially equal, only one direction has been plotted."
 The sizes of SHTs with Imax«512 could not be measured reliably due to their short running time and are therefore not shown., The sizes of SHTs with $\lmax<512$ could not be measured reliably due to their short running time and are therefore not shown.
" As can be seen, nneeds slightly less memory than for equivalent operations; overall, the scaling is very close to the expected P... if no precomputed aare used."," As can be seen, needs slightly less memory than for equivalent operations; overall, the scaling is very close to the expected $\lmax^2$ if no precomputed are used."
" For comparison purposes, a few data points for runs using precomputed scalar and tensor wwere also plotted; they clearly indicate the extremely high memory usage of these transforms, which scales with /2,.."," For comparison purposes, a few data points for runs using precomputed scalar and tensor were also plotted; they clearly indicate the extremely high memory usage of these transforms, which scales with $\lmax^3$."
" During the tests with multiple simultaneous threads it was observed that increasing the number of cores does not have a significant influence on the total required memory; for the concrete test discussed in refopenmp;caling, , memoryusage f ortherunwithl6thread swaso. ibpsht/. nlyaSolargerthan forthesingle— threadedrun."," During the tests with multiple simultaneous threads it was observed that increasing the number of cores does not have a significant influence on the total required memory; for the concrete test discussed in \\ref{openmp_scaling}, memory usage for the run with 16 threads was only larger than for the single-threaded run."
" Looking back at the goals outlined in refgoals and the results of the benchmark computations above, it can be concluded that the current version of the ppackage meets almost all specified requirements."," Looking back at the goals outlined in \\ref{goals} and the results of the benchmark computations above, it can be concluded that the current version of the package meets almost all specified requirements."
" It implements SHTs that have the same or a higher degree of accuracy as the other available implementations, can work on all spherical grids relevant to CMB studies, and is written in standard C, which is very widely supported."," It implements SHTs that have the same or a higher degree of accuracy as the other available implementations, can work on all spherical grids relevant to CMB studies, and is written in standard C, which is very widely supported."
 The employed algorithms are significantly more efficient than those published and used in this field of research before; their memory usage is economic and allows transforms whose input and output data occupy almost all available space., The employed algorithms are significantly more efficient than those published and used in this field of research before; their memory usage is economic and allows transforms whose input and output data occupy almost all available space.
" When appropriate, vector capabilities of the CPUs, as well as shared memory parallelism are exploited, resulting in further performance gains."," When appropriate, vector capabilities of the CPUs, as well as shared memory parallelism are exploited, resulting in further performance gains."
" Despite this range of capabilities, the package only consists of approximately 7000 lines of code, including the FFT implementation and in-line documentation for developers; except for a C compiler, it has no external dependencies."," Despite this range of capabilities, the package only consists of approximately 7000 lines of code, including the FFT implementation and in-line documentation for developers; except for a C compiler, it has no external dependencies."
" hhas been integrated into the simulation package of the mission(Level-S,, ?)), where it is interfaced with a development version of the C++ HEALPix and a Fortran90 code performing SHTs on detector beam patterns; this demonstrates the feasibility of interfacing the library with C++ and Fortran code."," has been integrated into the simulation package of the mission, \citealt{reinecke-etal-2006}) ), where it is interfaced with a development version of the C++ HEALPix and a Fortran90 code performing SHTs on detector beam patterns; this demonstrates the feasibility of interfacing the library with C++ and Fortran code."
" There are two areas in which libpsht’ss functionality should probably be extended: it currently does not provide transforms for spins larger than 2, and there is no active support for distributing SHTs over several independent computing nodes."," There are two areas in which s functionality should probably be extended: it currently does not provide transforms for spins larger than 2, and there is no active support for distributing SHTs over several independent computing nodes."
" As was mentioned in reftensorSHT,, addressing the first point is probably not too difficult, and enhancing the library with ,Y;, generators based on Wigner d matrix elements is planned for one of the next releases."," As was mentioned in \\ref{tensorSHT}, addressing the first point is probably not too difficult, and enhancing the library with ${}_sY_{lm}$ generators based on Wigner $d$ matrix elements is planned for one of the next releases."
" Regarding the second point, a combination of libpsht’ss central SHT functionality with the parallelisation strategy implemented by Radek Stompor’s library appears very promising and is currently being investigated."," Regarding the second point, a combination of s central SHT functionality with the parallelisation strategy implemented by Radek Stompor's library appears very promising and is currently being investigated."
" iis distributed under the terms of the GNU General Public License (GPL) version 2 (or later versions at the user’s choice); the most recent version, alongside online technical documentation can be obtained at the URL sourceforge.net/projects/libpsht/.."," is distributed under the terms of the GNU General Public License (GPL) version 2 (or later versions at the user's choice); the most recent version, alongside online technical documentation can be obtained at the URL ."
The TTauri class of objects contains highly luminous stars showing large-amplitude photometric variations with alternating deep and shallow minima.,The Tauri class of objects contains highly luminous stars showing large-amplitude photometric variations with alternating deep and shallow minima.
 The members are located in the high Juminosity end of the HI instability strip and the photometric variations are interpreted as being due to radial pulsations., The members are located in the high luminosity end of the II instability strip and the photometric variations are interpreted as being due to radial pulsations.
 There are two different photometric classes: the RVa stars are objects with a constant mean magnitude while the RVb objects display a long-term variation in their mean magnitude., There are two different photometric classes: the RVa stars are objects with a constant mean magnitude while the RVb objects display a long-term variation in their mean magnitude.
 ? introduced a spectroscopic classification of the TTaurt stars. using unfortunately the same alphabetic letters for the naming: RVA objects show strong absorption lines. RVB objects are of a somewhat hotter spectral type but are weak lined with enhanced CN and CH bands.," \citet{preston63} introduced a spectroscopic classification of the Tauri stars, using unfortunately the same alphabetic letters for the naming: RVA objects show strong absorption lines, RVB objects are of a somewhat hotter spectral type but are weak lined with enhanced CN and CH bands."
 The RVC objects are also weak lined but show no enhanced CN and CH molecular band strength., The RVC objects are also weak lined but show no enhanced CN and CH molecular band strength.
 A significant fraction of the TTauri stars show a large IR excess due to circumstellar dust and ? identified them as post-AGB objects on the basis of this IR excess. their luminosities and mass-loss history.," A significant fraction of the Tauri stars show a large IR excess due to circumstellar dust and \citet{jura86} identified them as post-AGB objects on the basis of this IR excess, their luminosities and mass-loss history."
 The photospheric content of TTaurt stars 15. however. very different from what could be expected in post-3rd dredge-up objects: they do not show high C-abundances or s-process overabundances but instead often show a depletion pattern in their photospheres (2222222)..," The photospheric content of Tauri stars is, however, very different from what could be expected in post-3rd dredge-up objects: they do not show high C-abundances or s-process overabundances but instead often show a depletion pattern in their photospheres \citep{giridhar94,giridhar98,giridhar00,gonzalez97b,gonzalez97a, vanwinckel98, maas05}."
 This abundance pattern is the result of gas-dust separation followed by reaccretion of the gas. which ts poor in refractory elements.," This abundance pattern is the result of gas-dust separation followed by reaccretion of the gas, which is poor in refractory elements."
 ? proposed that the most likely circumstance for this process to occur is when the dust is trapped 1n a circumstellar disc., \citet{waters92} proposed that the most likely circumstance for this process to occur is when the dust is trapped in a circumstellar disc.
 This depletion phenomenon is also observed in binary post-AGB stars with a disc (?).., This depletion phenomenon is also observed in binary post-AGB stars with a disc \citep{vanwinckel95}.
 This led ? to suggest that the depleted TTaurt stars must also be binaries with a disc., This led \citet{vanwinckel99} to suggest that the depleted Tauri stars must also be binaries with a disc.
 The likely presence of a Keplerian circumstellar dise was further shown by the systematic study by ?? of a large sample of binary post-AGB and TTaurt stars.," The likely presence of a Keplerian circumstellar disc was further shown by the systematic study by \citet{deruyter05,deruyter06} of a large sample of binary post-AGB and Tauri stars."
 ? list the post-AGB evolutionary tracks ofsingle stars of different initial mass., \citet{blocker95} list the post-AGB evolutionary tracks ofsingle stars of different initial mass.
 Evolutionary tracks for binary post-AGB stars have not yet been determined. however. since we find evidence that these stars have been shortcut on their AGB evolution (Sect. 8)).," Evolutionary tracks for binary post-AGB stars have not yet been determined, however, since we find evidence that these stars have been shortcut on their AGB evolution (Sect. \ref{conclusion}) ),"
 we expect longer lifetime scales because of the expected lower core masses., we expect longer lifetime scales because of the expected lower core masses.
 Typical post-AGB lifetimes of both stars are estimated to be of the order of ~107yr.," Typical post-AGB lifetimes of both stars are estimated to be of the order of $\sim 10^4\,$ yr."
 To investigate the special evolutionary status of TTauri stars and to research in detail the interplay between the photospheric and circumstellar environment in these systems. we focus in this paper on two well known TTaurt stars: HHer and CCen.," To investigate the special evolutionary status of Tauri stars and to research in detail the interplay between the photospheric and circumstellar environment in these systems, we focus in this paper on two well known Tauri stars: Her and Cen."
 HHer has been extensively studied in the literature., Her has been extensively studied in the literature.
 It is a binary TTauri star of photometric class a. with an orbital period of 1200 days (?)..," It is a binary Tauri star of photometric class a, with an orbital period of 1200 days \citep{vanwinckel98}."
 It is a very regular pulsator with a formal pulsation period (timespan between two successive deep photometric minima) of 75.5 days (?).., It is a very regular pulsator with a formal pulsation period (timespan between two successive deep photometric minima) of 75.5 days \citep{zsoldos93}.
 The mean magnitude of HHer is zy=7.69 mag and the amplitude Amy=2.3] mag (?).., The mean magnitude of Her is $m_V=7.69$ mag and the amplitude $\bigtriangleup m_V=2.31$ mag \citep{lloydevans85}.
 The presence of two strong shock waves in every formal pulsation cycle. causing line-profile deformations in the spectra of HHer and RSScuti has been discussed by ??..," The presence of two strong shock waves in every formal pulsation cycle, causing line-profile deformations in the spectra of Her and Scuti has been discussed by \citet{gillet89,gillet90}."
 HHer shows a chemical depletion pattern (22) that is attributed to the presence of a stable Keplerian dusty disce.," Her shows a chemical depletion pattern \citep{vanwinckel98,giridhar00} that is attributed to the presence of a stable Keplerian dusty disc."
 The presence of such a dise has also been proposed by ? who interpret the detection of weak CO rotational emission lines with a small velocity width (??) as a signature of such a long-lived dust reservoir.," The presence of such a disc has also been proposed by \citet{jura99} who interpret the detection of weak CO rotational emission lines with a small velocity width \citep{bujarrabal88,jura95} as a signature of such a long-lived dust reservoir."
 The presence of highly crystallme silicates in the infrared spectrum (?) and the strong millimeter continuum flux from large dust grains (?2) further corroborate this conclusion.," The presence of highly crystalline silicates in the infrared spectrum \citep{molster99} and the strong millimeter continuum flux from large dust grains \citep{shenton95,jura99} further corroborate this conclusion."
 ? claimed to have resolved the circumstellar material around HHer using N and Q-band imaging., \citet{jura00} claimed to have resolved the circumstellar material around Her using N and Q-band imaging.
 ? however detect no significant extended structure around HHer. using adaptive optics at mid-infrared wavelengths with a higher spatial resolution.," \citet{close03} however detect no significant extended structure around Her, using adaptive optics at mid-infrared wavelengths with a higher spatial resolution."
 CCen also is an TTaurt star of spectroscopic class B and photometric class a. It is a regular pulsator with a pulsation period of 64.6 days. a mean magnitude of wy=9.05 mag and an amplitude of Amy=1.28 mag (?)..," Cen also is an Tauri star of spectroscopic class B and photometric class a. It is a regular pulsator with a pulsation period of 64.6 days, a mean magnitude of $m_V=9.05$ mag and an amplitude of $\bigtriangleup m_V=1.28$ mag \citep{pollard96}. ."
 During every formal, During every formal
"Of these, the residual relative intensities of CIIA1334.5, ΟΙ A1302.2, and SilI A 1260.4 (the strongest metal lines tracing the neutral ISM) are very small in MS 1512-cB58 (Pettini et al.","Of these, the residual relative intensities of$\lambda$ 1334.5, OI $\lambda$ 1302.2, and SiII $\lambda$ 1260.4 (the strongest metal lines tracing the neutral ISM) are very small in MS 1512-cB58 (Pettini et al."
" 2002), The Cosmic Eye (Quider et al."," 2002), The Cosmic Eye (Quider et al."
" 2010), and the 8 O'Clock Arc (Dessauges-Zavadsky et al."," 2010), and the 8 O'Clock Arc (Dessauges-Zavadsky et al."
" but large in The Cosmic Horseshoe (Quider 2010),et al."," 2010), but large $\sim$ ) in The Cosmic Horseshoe (Quider et al."
 2009)., 2009).
" In fact, (~40%))the last object appears to a good example of the type of picket fence situation seen in the DCOs in our sample."," In fact, the last object appears to a good example of the type of picket fence situation seen in the DCOs in our sample."
" As discussed above, dust can also be a significant source of opacity to ionizing radiation in galaxies."," As discussed above, dust can also be a significant source of opacity to ionizing radiation in galaxies."
 Following Shapley et al. (, Following Shapley et al. (
"2006) and G09, we can distinguish between the absolute and relative escape fractions.","2006) and G09, we can distinguish between the absolute and relative escape fractions."
" The former is the actual fraction of ionizing photons that escape the galaxy, including the effect of dust."," The former is the actual fraction of ionizing photons that escape the galaxy, including the effect of dust."
 The latter neglects the effect of dust and is defined as the ratio of the fraction of escaping ionizing photons to escaping non-ionizing far-UV photons (e.g. it measures only the effect of the photoelectric opacity of the gas)., The latter neglects the effect of dust and is defined as the ratio of the fraction of escaping ionizing photons to escaping non-ionizing far-UV photons (e.g. it measures only the effect of the photoelectric opacity of the gas).
 Our measurements discussed above provide information about the relative escape fraction., Our measurements discussed above provide information about the relative escape fraction.
" Given that the ratio of far-IR to far-UV fluxes for the galaxies with DCOs are of-order ten, the implied absolute escape fractions will be proportionately smaller."," Given that the ratio of far-IR to far-UV fluxes for the galaxies with DCOs are of-order ten, the implied absolute escape fractions will be proportionately smaller."
 We list our estimates for both the relative and absolute escape fractions in Tables 1 (FUSE) and 2 (COS)., We list our estimates for both the relative and absolute escape fractions in Tables 1 (FUSE) and 2 (COS).
" As can be seen in Figure 4, the CIIA1334.5 absorption-line profiles in all the LBAs are blueshifted with respect to the galaxy systemic velocity, implying a large-scale outflow of gas."," As can be seen in Figure 4, the $\lambda$ 1334.5 absorption-line profiles in all the LBAs are blueshifted with respect to the galaxy systemic velocity, implying a large-scale outflow of gas."
 Further evidence for outflows is provided by the broad blue-asymmetric wings seen on the optical emission-line profiles (O09)., Further evidence for outflows is provided by the broad blue-asymmetric wings seen on the optical emission-line profiles (O09).
 The absorption-line profiles in Figure 4 show evidence for exceptionally high outflow speeds in the DCOs., The absorption-line profiles in Figure 4 show evidence for exceptionally high outflow speeds in the DCOs.
" To further investigate this, we turn to the Si III 1206.5 feature."," To further investigate this, we turn to the Si III $\lambda$ 1206.5 feature."
" This is the strongest metal line that is accessed in all four DCO spectra, and arises in the ionized gas."," This is the strongest metal line that is accessed in all four DCO spectra, and arises in the ionized gas."
" As shown in Figure 6, outflowing gas is detected at extraordinarily high velocities in all four galaxies with DCOs."," As shown in Figure 6, outflowing gas is detected at extraordinarily high velocities in all four galaxies with DCOs."
 The flux-weighted line centroid is blueshifted by about 700 km s! in these objects (Table 2)., The flux-weighted line centroid is blueshifted by about 700 km $^{-1}$ in these objects (Table 2).
" It is difficult to determine the maximum outflow velocity because of blending with the NIA1200 triplet, which lies ~1600 km s-! blueward of the SiIII line."," It is difficult to determine the maximum outflow velocity because of blending with the $\lambda$ 1200 triplet, which lies $\sim$ 1600 km $^{-1}$ blueward of the SiIII line."
" Conservatively, the maximum outflow speed seen in the DCOs reaches at least 1500 km s-!."," Conservatively, the maximum outflow speed seen in the DCOs reaches at least 1500 km $^{-1}$."
 These velocities are much higher than in the local starburst sample and in the LBAs without a DCO., These velocities are much higher than in the local starburst sample and in the LBAs without a DCO.
 This is shown in Figure 7 where we plot the outflow speeds in the ionized gas in local starburst and LBA samples the star formation rate and galaxy mass (see Tables 1 and 2)., This is shown in Figure 7 where we plot the outflow speeds in the ionized gas in local starburst and LBA samples the star formation rate and galaxy mass (see Tables 1 and 2).
 For galaxies with FUSE data we use the CIIIA977.0 and/or NIIA1084.0 and/or CIIA1036.3 lines., For galaxies with FUSE data we use the $\lambda$ 977.0 and/or $\lambda$ 1084.0 and/or $\lambda$ 1036.3 lines.
 The outflow velocities in the DCOs are also significantly larger than those typically seen in high-z galaxies., The outflow velocities in the DCOs are also significantly larger than those typically seen in high-z galaxies.
 Steidel et al. (, Steidel et al. (
2010) find line centroids that are blueshifed on-average by 164 km s! and maximum outflow speeds that are typically ~ 800 km s-! for galaxies with star formation rates of ~10! to 10?Mo year-!.,2010) find line centroids that are blueshifed on-average by 164 km $^{-1}$ and maximum outflow speeds that are typically $\sim$ 800 km $^{-1}$ for galaxies with star formation rates of $\sim 10^1$ to $10^2 M_{\odot}$ $^{-1}$.
 Outflows from intensely star-forming galaxies are believed to be produced as gas clouds are accelerated by the ram pressure of a hot and fast wind driven by the collective thermal/kinetic energy supplied by supernovae and massive stellar winds (e.g. Heckman et al., Outflows from intensely star-forming galaxies are believed to be produced as gas clouds are accelerated by the ram pressure of a hot and fast wind driven by the collective thermal/kinetic energy supplied by supernovae and massive stellar winds (e.g. Heckman et al.
 2000; Veilleux et al., 2000; Veilleux et al.
 2005; Marcolini et al., 2005; Marcolini et al.
" 2005; Strickland Heckman 2010), and/or by radiation pressure acting on dust (Murray et al."," 2005; Strickland Heckman 2010), and/or by radiation pressure acting on dust (Murray et al."
 2005; 2010)., 2005; 2010).
 Can the unusually high velocities seen in the DCOs be explained in this way?, Can the unusually high velocities seen in the DCOs be explained in this way?
 We consider a simple idealized model of a cloud with a column density N accelerated outward from an initial radius rg by the ram pressure of a wind that carries momentum at a rate p into a solid angle 2., We consider a simple idealized model of a cloud with a column density $N$ accelerated outward from an initial radius $r_0$ by the ram pressure of a wind that carries momentum at a rate $\dot{p}$ into a solid angle $\Omega$.
" Then the terminal velocity of this cloud will 2) Vterm=1800937(0/42)όροΝο km su! Here the momentum flux is in units of 1055 dynes, the initial radius is in units of 100 pc and the cloud column density is in units of 1031 cm-?."," Then the terminal velocity of this cloud will 2) $v_{term} = 1800 \dot{p}_{35}^{1/2} (\Omega/ 4\pi)^{-1/2} r_{0,100}^{-1/2}
N_{21}^{-1/2}$ km $^{-1}$ Here the momentum flux is in units of $10^{35}$ dynes, the initial radius is in units of 100 pc and the cloud column density is in units of $10^{21}$ $^{-2}$ ."
" The DCO star formation rates imply momentum fluxes of this order (Leitherer Heckman 1995), while the HST images yield typical DCO radii of 100 pc (O9)."," The DCO star formation rates imply momentum fluxes of this order (Leitherer Heckman 1995), while the HST images yield typical DCO radii of 100 pc (O9)."
" As discussed below in section 4.6,"," As discussed below in section 4.6,"
"Errors in. M, were calculated by Monte Carlo sampling of uncorrelated errors in SSPP metallicity. &o magnitude and ro magnitude. drawn randomly from a Gaussian with a width equal to the reported errors on those quantities.","Errors in $M_{\mathrm{r}}$ were calculated by Monte Carlo sampling of uncorrelated errors in SSPP metallicity, $g_{\mathrm{0}}$ magnitude and $r_{\mathrm{0}}$ magnitude, drawn randomly from a Gaussian with a width equal to the reported errors on those quantities."
" This addition of error was done 107 times for each star. and we take the standard deviation in those 10 determinations of M, as the error on M,."," This addition of error was done $10^{4}$ times for each star, and we take the standard deviation in those $10^{4}$ determinations of $M_{\mathrm{r}}$ as the error on $M_{\mathrm{r}}$."
" This error has a typical value of 0.3 magnitudes. with the largest values (20.5 magnitudes) on the 3% of stars with the largest errors in apparent ry and go.The age of the isochrones used had minimal effects on the derived M, values. with a change of only £0.06 magnitudes for a shift of +1 Gyr."," This error has a typical value of $0.3$ magnitudes, with the largest values $\ga 0.5$ magnitudes) on the $3\%$ of stars with the largest errors in apparent $r_{\mathrm{0}}$ and $g_{\mathrm{0}}$.The age of the isochrones used had minimal effects on the derived $M_{\mathrm{r}}$ values, with a change of only $\pm 0.06$ magnitudes for a shift of $\mp 1$ Gyr."
 The limited metallicity range of the isochrones i5 sufficient for the purposes of the analysis in Section3. but lower-metallicity isochrones would allow us to study the distance distribution of the roughly 1/3 of our final dataset at lower metallicity. to compare the spatial distribution of our final data set to the inner and outer halos identified in).," The limited metallicity range of the isochrones is sufficient for the purposes of the analysis in Section 3, but lower-metallicity isochrones would allow us to study the distance distribution of the roughly 1/3 of our final dataset at lower metallicity, to compare the spatial distribution of our final data set to the inner and outer halos identified in."
. We measured $(3839)1981)... a bandstrength index for the CN band at 38834. for all halo giant spectra.," We measured $S(3839)$, a bandstrength index for the CN band at $3883 \hbox{\AA}$, for all halo giant spectra."
 $(3839 measures the magnitude difference between the integrated flux in the CN feature and the integrated flux in a nearby continuum band. with more absorption in the feature resulting m larger bandstrength.," $S(3839)$ measures the magnitude difference between the integrated flux in the CN feature and the integrated flux in a nearby continuum band, with more absorption in the feature resulting in larger bandstrength."
 As can be seen in Fig. 2..," As can be seen in Fig. \ref{ff2},"
 there is a strong concentration at low $(3839) in our data set., there is a strong concentration at low $S(3839)$ in our data set.
 This is to be expected. since the halo is primarily composed of CN-weak stars with typical Pop.," This is to be expected, since the halo is primarily composed of CN-weak stars with typical Pop."
 II abundances., II abundances.
 There is also a clear trend with temperature. in the sense that cooler stars have larger CN bandstrengths.," There is also a clear trend with temperature, in the sense that cooler stars have larger CN bandstrengths."
 The cooling of stars as they ascend the giant branch reddens the spectra and permits more CN molecule formation. both of which increase the flux difference between the feature and continuum bands of $(3839).," The cooling of stars as they ascend the giant branch reddens the spectra and permits more CN molecule formation, both of which increase the flux difference between the feature and continuum bands of $S(3839)$."
 There are. however. interesting outliers in Fig. 2: ," There are, however, interesting outliers in Fig. \ref{ff2}: :"
many of the stars with dramatically large $(3839) wre carbon stars (shown as open triangles). with correspondingly large CH and C» bandstrengths.," many of the stars with dramatically large $S(3839)$ are carbon stars (shown as open triangles), with correspondingly large CH and $_{2}$ bandstrengths."
 Figure 3. shows four sample spectra from the halo giant data set: the uppermost spectrum (of SDSS J035123.904-092451.3) is à low-metallicity carbon star. the next (of SDSS 1115934.87-002748.0) is a possible CEMP star (Carbon-Enhanced Metal-Poor. having [Fe/H]<—2.0 and [C/Fe]2+1.0. and deseribed in Lucatello et al.," Figure \ref{ff3} shows four sample spectra from the halo giant data set: the uppermost spectrum (of SDSS J035123.90+092451.3) is a low-metallicity carbon star, the next (of SDSS J115934.87+002748.0) is a possible CEMP star (Carbon-Enhanced Metal-Poor, having $\la -2.0$ and $\ga +1.0$, and described in Lucatello et al."
 and references therein). with strong CH absorption redward of the G band and a low metallicity. the next spectrum (of SDSS JO64411.96+275351.6) is not a carbon star. but is CN-strong. and the lowest spectrum (of SDSS J145301.24-001954.1) isa typical CN-weak halo star.," and references therein), with strong CH absorption redward of the G band and a low metallicity, the next spectrum (of SDSS J064411.96+275351.6) is not a carbon star, but is CN-strong, and the lowest spectrum (of SDSS J145301.24-001954.1) is a typical CN-weak halo star."
 The broad molecular features in the example carbon-star spectrum are quite clear., The broad molecular features in the example carbon-star spectrum are quite clear.
 We use the strength of the CH feature around 4350A and the Swan (1.0) C» band at 4737A to identify carbon and CEMP stars and remove them from the final data set.," We use the strength of the CH feature around $4350\hbox{\AA}$ and the Swan (1,0) $_{2}$ band at $4737\hbox{\AA}$ to identify carbon and CEMP stars and remove them from the final data set."
" Specifically. we measure these indices: and consider all stars with [Fe/H]x-1.8 and s(e0)> —0.093: 1.8 <[Fe/H]<—1.4. s(00)>—0.05. and s(cl)>O.1S: and [Fe/H]>—1.4. 0)>—0.02. and s(el)>0.158 to have ""strong carbon features""."," Specifically, we measure these indices: and consider all stars with $\le -1.8$ and $s(c0)\ge -0.093$ ; $-1.8 \le $ $\le -1.4$, $s(c0) \ge -0.05$, and $s(c1)\ge 0.15$; and $\ge -1.4$, $s(c0)\ge -0.02$, and $s(c1)\ge 0.158$ to have “strong carbon features”."
 Although these carbon and CEMP stars tend to have strong UV CN bands. as is demonstrated in. e.g... (2002)... that is a result of the unusually large ratio of carbon to oxygen in their atmospheres. and is not a result of the anticorrelated C-Nabundance pattern," Although these carbon and CEMP stars tend to have strong UV CN bands, as is demonstrated in, e.g., , that is a result of the unusually large ratio of carbon to oxygen in their atmospheres, and is not a result of the anticorrelated C-Nabundance pattern"
mechanism is at work iu the case of a binary.,mechanism is at work in the case of a binary.
 However. the trajectories of both perturbing masses are now closed orbits.," However, the trajectories of both perturbing masses are now closed orbits."
 Further complicating the analysis is the fact that each star draws tu matter [rom the wake of its companion., Further complicating the analysis is the fact that each star draws in matter from the wake of its companion.
 Naive application of the staudarcd dyuamical friction formulae would grossly misestiniate the torque., Naive application of the standard dynamical friction formulae would grossly misestimate the torque.
 The new approach introduced bere concentrates on the fact that the binary as a whole must shed angular momenuttun to the external medium., The new approach introduced here concentrates on the fact that the binary as a whole must shed angular momentum to the external medium.
 The actual mechanism is that the orbiting stars create au oscillatingMm Oogravitational potential that torques uearby eas., The actual mechanism is that the orbiting stars create an oscillating gravitational potential that torques nearby gas.
 This fluctuating torque eenerates outgolug acoustic waves. which transport angular momentum.," This fluctuating torque generates outgoing acoustic waves, which transport angular momentum."
 By calculating the total augular momentum efflux [rou the binary. the braking torque may be obtained without cousicering the complex star-gas interaction close to the stars themselves.," By calculating the total angular momentum efflux from the binary, the braking torque may be obtained without considering the complex star-gas interaction close to the stars themselves."
 Section 2 below formulates the problem mathematically aud derives the governing wave equation., Section 2 below formulates the problem mathematically and derives the governing wave equation.
 The quadrupolar driving potential created by the rotating stars is established in Section 3. while Section [ discusses the physical character of the generated) waves.," The quadrupolar driving potential created by the rotating stars is established in Section 3, while Section 4 discusses the physical character of the generated waves."
 Section 3 presents the centra result of this investigation., Section 5 presents the central result of this investigation.
 Here the augular momentum trausported by the spiral wave is obtained. auc thereby the torque (see eq. (," Here the angular momentum transported by the spiral wave is obtained, and thereby the torque (see eq. ("
39)).,39)).
 It is also shown that the wave carries off all tle mechauica euergvy released by the sluiuking binary., It is also shown that the wave carries off all the mechanical energy released by the shrinking binary.
 The resulting temporal evolution of the syste is cousidere in Section 6. along with the efficacy of braking in infrared dark clouds.," The resulting temporal evolution of the system is considered in Section 6, along with the efficacy of braking in infrared dark clouds."
 Finally. Sectiou 7 compares the results obtained here with the traditional analysis aid cliscusses future extensious of this study.," Finally, Section 7 compares the results obtained here with the traditional analysis and discusses future extensions of this study."
 We wish to treat the binary as a perturbiug mass embeccled in otherwise static. uniform gas.," We wish to treat the binary as a perturbing mass embedded in otherwise static, uniform gas."
 If the binary itself recently formed. then gas surrouncing it would not be perfectly quiescent.," If the binary itself recently formed, then gas surrounding it would not be perfectly quiescent."
 Hence. our assiiiied background is a highly idealized represeutation of a real cloud. or at least that portion oL a cloud iu which we cau accurately follow the propagation of acoustic waves from the stars.," Hence, our assumed background is a highly idealized representation of a real cloud, or at least that portion of a cloud in which we can accurately follow the propagation of acoustic waves from the stars."
 Oue obvious stipulation is that this region cannot be so close to the binary that cloud gas is infalline onto the stars., One obvious stipulation is that this region cannot be so close to the binary that cloud gas is infalling onto the stars.
 This requiremeut sets ai inuer radius of validity for the analysis. which we take to be the sonic point in the Bondi accretion problem: Here. Mj is the total binary mass and c; the sound speed of the surrounding gas. assumed to be isothermal.," This requirement sets an inner radius of validity for the analysis, which we take to be the sonic point in the Bondi accretion problem: Here, $M_{\rm tot}$ is the total binary mass and $c_s$ the sound speed of the surrounding gas, assumed to be isothermal."
 Iu the numerical evaluation of rij. we have used for e; the typical observed velocity," In the numerical evaluation of $r_{\rm in}$, we have used for $c_s$ the typical observed velocity"
low velocity disversion. star streams du ealactic halos are sensitive to the deeree to which the dark matter in the halo is sub-structured imto thousancs of orbiting sub-halos.,low velocity dispersion star streams in galactic halos are sensitive to the degree to which the dark matter in the halo is sub-structured into thousands of orbiting sub-halos.
" The sub-halos fold aid chop the star-strealus and eraduallv increase the veocity disper""on around the mean motio1 to about x1! of the indo circular velocity. typically over a Hubble time (hataYoon.Johuston&Ποσο 2010)."," The sub-halos fold and chop the star-streams and gradually increase the velocity dispersion around the mean motion to about $\simeq$ of the halo circular velocity, typically, over a Hubble time \citep{Ibata:02,SGV:08,StarStreams,YJH:10}."
. C'onsequenIv. cool star Sreams are sensitive 1idicators of the presence of the predicted dark matter substrucure.," Consequently, cool star streams are sensitive indicators of the presence of the predicted dark matter substructure."
 About half à dozen of the currentv know1i Milkv Wav streams qualify as cool (iiost couficently. Pal 5. GD-1. Orphan. Archerou aud Styx) tlat ls. having loca velocity dispersions beow about or. width less than ayout 0.1 radiau as seen from the ceuter of the host ealaxy.," About half a dozen of the currently known Milky Way streams qualify as cool (most confidently, Pal 5, GD-1, Orphan, Archeron and Styx) that is, having local velocity dispersions below about, or, width less than about 0.1 radian as seen from the center of the host galaxy."
 A dark matter sub-halo crossing sucli a coo stream will lead to visible distiwhaices., A dark matter sub-halo crossing such a cool stream will lead to visible disturbances.
 Uifortunately. the star count «ata often do not vet have sfhcieut loca numbers to allow staistically significant measurements of desity variations] relative to the ealactic foregrounc and backeroux (Odeusirchenctal.2001:Fergusonetmaiz2009:Odenkircheuetal.2 009).," Unfortunately, the star count data often do not yet have sufficient local numbers to allow statistically significant measurements of density variations relative to the galactic foreground and background \citep{Odenkirchen:02, Ferguson:02,Majewski:04,Chapman:08, Grillmair:09,Odenkirchen:09}. ."
. The Pan Aucromea Archeological Survey (PAudAS) (MeConunachieetal.20109) in one fell swoop provides deep and uniform daa to a distance of about 150 kpc from M31's center., The Pan Andromeda Archeological Survey (PAndAS) \citep{Pandas} in one fell swoop provides deep and uniform data to a distance of about 150 kpc from M31's center.
 The spectacular star stream north-west of M31. nore than 100 Ipc lone. was first displaved in its cutirety in Richardsonetal.(2011).," The spectacular star stream north-west of M31, more than 100 kpc long, was first displayed in its entirety in \citet{Richardson:11}."
.. The ereat leugth of the stream. as well as beiug fairly distant frou the disturbing effects of the main body of M31. make the stream an exceptionally interesting case for analysisof density variations.," The great length of the stream, as well as being fairly distant from the disturbing effects of the main body of M31, make the stream an exceptionally interesting case for analysisof density variations."
 Aud. the star couuts have sufficieut," And, the star counts have sufficient"
Each GRB was prepared in a manuer similar o that described above for BATSE trigger 2193.,Each GRB was prepared in a manner similar to that described above for BATSE trigger 2193.
 Here. however. it was tested to see if the pulse slape in energy chaunels 1 aud 2 could be scaled in time and fIuence to match the pulse shape [οιud in energy channel 3.," Here, however, it was tested to see if the pulse shape in energy channels 1 and 2 could be scaled in time and fluence to match the pulse shape found in energy channel 3."
 The uuuber of couuts was so few in enerey chauuel | for each burst that any test involving data [rom this chanuel was essentially meatingless., The number of counts was so few in energy channel 4 for each burst that any test involving data from this channel was essentially meaningless.
H A X> per degree ofH fH'eecom. (d) statistic was computed betwee the ↕∖∖↽∩↥∩∖∖↽≺↵↕⋅≺↵∐≺↵↥⋅∑≟⊽∖⊽∢∙∐⋜↕∐∐≺↵↥⊳∖⋜↕, A $\chi^2$ per degree of freedom $d$ ) statistic was computed between the two lower energy channels and channel 3.
"⋯⇂∢∙∐⋜↕∐∐≺↵↥∙⋝⋅∌∩↥⋅⊟⇀−∖⊺⊱↕↢↵⋃⋅↓∑≟∑≟≺↵↕⋅−≻∙⋝↖∖↙⋅↕∐↩∖↕⋮⋡⋝↥↽≻≺↵↕⋅≺⇂≺↵∑⇁⋟↓⋅≺↵≺↲∩⊔↕⋅≺↲≺↵≺⇂∩⋯ . E SÓLO . QT 2 ⋅⋅ between channels 1aid 3 was 5.06. while between cinnels 2 and 3 \3,/d=2.20."," For BATSE trigger 2387, the $\chi_{13}^2$ per degree of freedom between channels 1 and 3 was 5.06, while between channels 2 and 3 $\chi_{23}^2/d = 2.20$."
" For BATSE triggerMD Via/d=125 and \3,/d= 1.07."," For BATSE trigger 3003, $\chi_{13}^2/d = 1.25$ and $\chi_{23}^2/d = 1.07$ ."
" For BATSE trigger 3267. ντα=0.98 and \3,/d=1.09."," For BATSE trigger 3267, $\chi_{13}^2/d = 0.98$ and $\chi_{23}^2/d = 1.09$."
" Lastly. for BATSE lggerMD 63106. Via/d=0.72 and \3,/d=0.71."," Lastly, for BATSE trigger 6346, $\chi_{13}^2/d = 0.72$ and $\chi_{23}^2/d = 0.74$."
 These tests were all carrie: oul uxing 61-nms time-binec data., These tests were all carried out using 64-ms time-binned data.
 Plots where the data is binned o 1.02l-secoucd bins are shown iu. Figures 6-9., Plots where the data is binned to 1.024-second bins are shown in Figures 6-9.
 Trends i the data are easier for the human eve to cliscer Lou this large time-bin size., Trends in the data are easier for the human eye to discern on this large time-bin size.
 The best-found scale factors between the first two energy. channelJd. all Channel 3 has been applied., The best-found scale factors between the first two energy channels and channel 3 has been applied.
" Chanuel 3 data is plotted with a coutinuous line. while channel 1 atd channel 2 data are plotted with the sviubols 717 aud 72"" respectively."," Channel 3 data is plotted with a continuous line, while channel 1 and channel 2 data are plotted with the symbols “1"" and “2"" respectively."
 From these results. it appears that tle combination Pulse Start aud Pulse Scale Conjectures hold in all cases with the possible exceptio 1of BATSE trigger 2387.," From these results, it appears that the combination Pulse Start and Pulse Scale Conjectures hold in all cases with the possible exception of BATSE trigger 2387."
 Even for BATSE trigger 2387. the Pulse Start aud Pulse Scale Conjectures were mareiually cousisteut between enerey chanuels 2 aud 3.," Even for BATSE trigger 2387, the Pulse Start and Pulse Scale Conjectures were marginally consistent between energy channels 2 and 3."
) It is speculated that the conjecture tests mieht have been compromised between chauuels 1 and 3 because of a second dim soft pulse that occurred well after the peak of the main pulse., It is speculated that the conjecture tests might have been compromised between channels 1 and 3 because of a second dim soft pulse that occurred well after the peak of the main pulse.
" From prelimitary luspectiou of the most [Iuent S""ele-pulse GRB on the Norris.(1999) ist. BATSE triggerMD 2193. Iwo conjectwes have been suggested."," From preliminary inspection of the most fluent single-pulse GRB on the \citet{Nor99} list, BATSE trigger 2193, two conjectures have been suggested."
 They are: 1., They are: 1.
 The Pulse Start Coujecure: GRB pulses each have a unique starting time that is indepeucdent of energy., The Pulse Start Conjecture: GRB pulses each have a unique starting time that is independent of energy.
 2., 2.
 The Pulse Scale Conjecture: GRB pulses have a uiicte shape that is independent of energy., The Pulse Scale Conjecture: GRB pulses have a unique shape that is independent of energy.
 The relation between tle shape of a GRB pulse at any euerey aud the shape of the same GRB pulse at auy other energy a'e simple scale factors in time ane aiplitucle., The relation between the shape of a GRB pulse at any energy and the shape of the same GRB pulse at any other energy are simple scale factors in time and amplitude.
 These two conjectures were then tested on BATSE viewer 2193 and found statistically valid iu sieuificantly [Iteut euergy chanuels., These two conjectures were then tested on BATSE trigger 2193 and found statistically valid in significantly fluent energy channels.
 The two conjecttJes were also indicated as true for three of the next four most fluent siugle-pulse GRBs. with the ¢Inc'paut case possibly beiug affected by the presence of a suall secondary. pulse.," The two conjectures were also indicated as true for three of the next four most fluent single-pulse GRBs, with the discrepant case possibly being affected by the presence of a small secondary pulse."
 When the Pulse Scale coujecure οί». a unique tiue-scaliug factor is revealed between GRB energy bands.," When the Pulse Scale conjecture holds, a unique time-scaling factor is revealed between GRB energy bands."
 For BATSE trigger 2193. this scale [actor was found to be a power aw lucreasing monotonically from about LOO IxeV o 1000 IxeV. One might be surprised that the Pulse Scale Conjecture can be tested at all with current BATSE data. since the energy channels available all have finite energy width.," For BATSE trigger 2193, this scale factor was found to be a power law increasing monotonically from about 100 KeV to 1000 KeV. One might be surprised that the Pulse Scale Conjecture can be tested at all with current BATSE data, since the energy channels available all have finite energy width."
 In its purest form. the Pulse Scale Coyecttwe predicts a scaling relation between pulseslapes at mouochromatic energies," In its purest form, the Pulse Scale Conjecture predicts a scaling relation between pulseshapes at monochromatic energies"
"color-color diagrams, luminosity functions (LFs) and MFs.","color–color diagrams, luminosity functions (LFs) and MFs."
 Such studies are not possible using heterogenous datasets where unknown biases may be present., Such studies are not possible using heterogenous datasets where unknown biases may be present.
 The initial mass function (IMF) is the distribution of stars of varying masses from the original parent cloud., The initial mass function (IMF) is the distribution of stars of varying masses from the original parent cloud.
 The universality of the IMF and the influence of environment on star formation is still a matter of debate., The universality of the IMF and the influence of environment on star formation is still a matter of debate.
" As these clusters are very young (age < 10 Myr), their MF may be approximated as the IMF."," As these clusters are very young (age $\leq$ 10 Myr), their MF may be approximated as the IMF."
" The sample of clusters have been made from differing initial conditions and subject to varied influences of external interactions with the galactic field, thus leading to observable differences, which we explore."," The sample of clusters have been made from differing initial conditions and subject to varied influences of external interactions with the galactic field, thus leading to observable differences, which we explore."
" However, from a recent study of ?,, even in the case of very young clusters, there is a change in the MF due to the dynamics of young clusters which loose a significant fraction of their stars at an early age."," However, from a recent study of \cite{kroupa07}, even in the case of very young clusters, there is a change in the MF due to the dynamics of young clusters which loose a significant fraction of their stars at an early age."
" Mass segregation is the redistribution of stars according to their masses, thus leading to the concentration of high mass stars near the centre and the low mass ones away from the centre."," Mass segregation is the redistribution of stars according to their masses, thus leading to the concentration of high mass stars near the centre and the low mass ones away from the centre."
" This has been observed in a variety of clusters, both young and old."," This has been observed in a variety of clusters, both young and old."
 The variation of the MF of these clusters is determined in different regions of the clusters and their values are compared., The variation of the MF of these clusters is determined in different regions of the clusters and their values are compared.
" Further, we estimate from the value of 7=/,,/!,,;,,. the degree of mass segregation expected due to dynamical effects and compare it with our observations."," Further, we estimate from the value of $\tau = t_{age}/t_{relax}$, the degree of mass segregation expected due to dynamical effects and compare it with our observations."
 The relaxation time {ως is a characteristic time in which there is an equipartition of energy and the high mass stars with lesser kinetic energy sink to the core and the low mass stars move to the outer regions of the cluster (?).., The relaxation time $t_{relax}$ is a characteristic time in which there is an equipartition of energy and the high mass stars with lesser kinetic energy sink to the core and the low mass stars move to the outer regions of the cluster \citep{binney08}.
 The value of 7 indicates whether an excess of high mass stars in the cores of clusters is a result of dynamical evolution or the imprint of the star formation process itself., The value of $\tau$ indicates whether an excess of high mass stars in the cores of clusters is a result of dynamical evolution or the imprint of the star formation process itself.
 The parameter 7 has been described as an evolutionary parameter (?) , The parameter $\tau$ has been described as an evolutionary parameter \citep{bonatto05} 
The solar magnetic 11-vear evele is thought to be sustained by the dvnamo motion of (he internal ionized plasma (Parker1955)..,The solar magnetic $11$ -year cycle is thought to be sustained by the dynamo motion of the internal ionized plasma \citep{1955ApJ...122..293P}.
 Dased on (he internal structure of the velocity field. i.e.. the meridional flow and the differential rotation revealed by. helioseismology (seeThompsonetal. 2003).. the flux-transport dynamo was suggested (Choudluri1995:Dikpati&Charbonneau1999;INükeretal.2001:HottaYokovama 2010).. as a model to successfully explain some features of solar activity such as the equatorware migration of sunspots and (he poleward migration of the surface," Based on the internal structure of the velocity field, i.e., the meridional flow and the differential rotation revealed by helioseismology \citep[see review by][]{2003ARA&A..41..599T}, the flux-transport dynamo was suggested \citep{1995A&A...303L..29C,1999ApJ...518..508D,2001A&A...374..301K,2010ApJ...709.1009H}, as a model to successfully explain some features of solar activity such as the equatorward migration of sunspots and the poleward migration of the surface"
corresponding to a mass of 0.93 would produce a 0.51 wwhite dwarf remnant.,corresponding to a mass of 0.93 – would produce a 0.51 white dwarf remnant.
 The above empirical initialfinal mass relation has been derived using while dwarfs from the thin disk. while the relation lor halo white dwarls which is even more poorly known is probably similar to that of elobular clusters.," The above empirical initial-final mass relation has been derived using white dwarfs from the thin disk, while the relation for halo white dwarfs — which is even more poorly known — is probably similar to that of globular clusters."
 As discussed by (1996).. the mass of the white dwarls currently being formed in globular clusters can be constrained by the liminosities of the red eiut branch tip. the horizontal branch. the AGB termination. and the post-AGB stars. all of which are sensitive to the mass of the hycrogen exhausted core.," As discussed by \citet{renzini96}, the mass of the white dwarfs currently being formed in globular clusters can be constrained by the luminosities of the red giant branch tip, the horizontal branch, the AGB termination, and the post-AGB stars, all of which are sensitive to the mass of the hydrogen exhausted core."
 All observations point to values between Mw=0.51 and 0.55AL... virtually independent of metallicity (Renzini&FusiPecci1988).," All observations point to values between $M_{\rm
WD}=0.51$ and 0.55, virtually independent of metallicity \citep{renzini88}."
. Hence it is reasonable to assume that white chwarls currently being formed in the halo should have masses in the same range., Hence it is reasonable to assume that white dwarfs currently being formed in the halo should have masses in the same range.
 The empirical mitial-final mass relation we derived above is certainly consistent. wilh these results. although it should be considered a good approximation at best. and white clwarls curently being formed in the halo could still be as massive as 0.55M.," The empirical initial-final mass relation we derived above is certainly consistent with these results, although it should be considered a good approximation at best, and white dwarfs currently being formed in the halo could still be as massive as 0.55."
.. The isochrones representing the white dwarl cooling ages plus main sequence ages using (he initial-Dinal mass relation described above are reproduced in Figure 9.., The isochrones representing the white dwarf cooling ages plus main sequence ages using the initial-final mass relation described above are reproduced in Figure \ref{fg:f9}.
 It is clear that the total age of a white dwarl is strongly mass-dependent. a result which stresses the importance of determining reliable masses through precise (rigonometric parallax measurements.," It is clear that the total age of a white dwarf is strongly mass-dependent, a result which stresses the importance of determining reliable masses through precise trigonometric parallax measurements."
 For instance. all white dwarls with masses below A/S0.5 cecannot have been formed within the lifetime of the Galaxy. and thev must be the result ol common envelope evolution.," For instance, all white dwarfs with masses below $M\lesssim0.5$ cannot have been formed within the lifetime of the Galaxy, and they must be the result of common envelope evolution."
 Alternatively. these could be unresolved degenerate binaries. and their overluminosity would be wrongly interpreted here as single white dwarls wilh large radii and low masses (see BRL and BLR for further discussion).," Alternatively, these could be unresolved degenerate binaries, and their overluminosity would be wrongly interpreted here as single white dwarfs with large radii and low masses (see BRL and BLR for further discussion)."
 Also. the results of Figure 9 illustrate how a 12 Gyr old white dwarL. sav. could be found at any effective temperature. as long as ils mass is precisely on the horizontal part of the isochrones near ~0.5M... implving that is has recently (a lew Gyr) evolved [rom a main sequence star slightly below el1 ((see discussion above).," Also, the results of Figure \ref{fg:f9}
 illustrate how a 12 Gyr old white dwarf, say, could be found at any effective temperature, as long as its mass is precisely on the horizontal part of the isochrones near $\sim0.5$, implying that is has recently (a few Gyr) evolved from a main sequence star slightly below $\sim 1$ (see discussion above)."
 Only three white dwarls Irom the OILDIDS sample have trigonometric parallax measurements., Only three white dwarfs from the OHDHS sample have trigonometric parallax measurements.
 One of them is the extremely massive white dwarf LIIS 4033 (Dahnetal.2004) seen in the upper left corner of Figure 9.., One of them is the extremely massive white dwarf LHS 4033 \citep{dahn04} seen in the upper left corner of Figure \ref{fg:f9}.
 50 not only this star does not have (he proper kinematics to be associated with the halo population. but it is also much too voung (7<2 Gyr).," So not only this star does not have the proper kinematics to be associated with the halo population, but it is also much too young $\tau<2$ Gyr)."
 The other two objects. LIIS 147 and LIIS 542. have more normal masses of Af=0.64 and 0.67AL... respectively.," The other two objects, LHS 147 and LHS 542, have more normal masses of $M=0.64$ and 0.67, respectively."
 Taken at face value. thev both appear too voung to be associated with the halo population. despite their halo kinematics.," Taken at face value, they both appear too young to be associated with the halo population, despite their halo kinematics."
 However. when the mass uncertainties are taken," However, when the mass uncertainties are taken"
of stellar mass.,of stellar mass.
 The two panels are for galaxies residing in haloes of different mass., The two panels are for galaxies residing in haloes of different mass.
 In the following we will refer to galaxies in haloes with mass >8x1014 as ‘cluster’ ellipticals and as ‘field’ ellipticals to all modelΜω ellipticals residing in less massive haloes.," In the following we will refer to galaxies in haloes with mass $\gtrsim 8\times10^{14}\,{\rm M}_{\sun}$ as `cluster' ellipticals and as `field' ellipticals to all model ellipticals residing in less massive haloes."
" In both panels, the solid line shows the average star formation history for all the elliptical galaxies in the sample under investigation."," In both panels, the solid line shows the average star formation history for all the elliptical galaxies in the sample under investigation."
" The long dashed, dash-dotted, dashed, and dotted lines refer to galaxies with stellar mass ~1x 10, 1x 10'', 1x10'°, and 1x10°Mo,respectively?.."," The long dashed, dash-dotted, dashed, and dotted lines refer to galaxies with stellar mass $\simeq 1\times10^{12}$ , $1\times10^{11}$ , $1\times10^{10}$, and $1\times10^{9}\,{\rm M}_{\sun}$,."
 The vertical lines are included to guide the eye and mark the peak of the 10/3 Mo-ellipticals in the top panel.," The vertical lines are included to guide the eye and mark the peak of the $10^{12}\,{\rm M}_\odot$ -ellipticals in the top panel."
 Fig., Fig.
" 1 shows the most important result of this paper: more massive elliptical galaxies have star formation histories that peak at higher redshifts (~ 5) than lower mass systems, and can reach star formation rates up to several thousands of solar masses per year for galaxies ending up in overdense regions."," \ref{fig:sfr} shows the most important result of this paper: more massive elliptical galaxies have star formation histories that peak at higher redshifts $\simeq 5$ ) than lower mass systems, and can reach star formation rates up to several thousands of solar masses per year for galaxies ending up in overdense regions."
 Less massive elliptical galaxies have star formation histories that peak at progressively lower redshifts and are extended over a longer time interval., Less massive elliptical galaxies have star formation histories that peak at progressively lower redshifts and are extended over a longer time interval.
 A comparison of the top and bottom panels of Fig., A comparison of the top and bottom panels of Fig.
" 1 shows that the qualitative behaviour for ‘field’ and ‘cluster’ ellipticals is the same, but that for fixed mass, the star formation histories of field ellipticals are predicted to be more extended than those of ellipticals in clusters."," \ref{fig:sfr} shows that the qualitative behaviour for `field' and `cluster' ellipticals is the same, but that for fixed mass, the star formation histories of field ellipticals are predicted to be more extended than those of ellipticals in clusters."
" This is a natural outcome of the hierarchical scenario, where haloes in regions of the Universe that are destined to form a cluster collapse earlier and merge more rapidly."," This is a natural outcome of the hierarchical scenario, where haloes in regions of the Universe that are destined to form a cluster collapse earlier and merge more rapidly."
 The star formation histories shown in Fig., The star formation histories shown in Fig.
" 1 represent averages computed over all the elliptical galaxies in the simulation box, but the trends remain true also when a much smaller volume of the simulation, and hence a much smaller sample size, is analysed."," \ref{fig:sfr} represent averages computed over all the elliptical galaxies in the simulation box, but the trends remain true also when a much smaller volume of the simulation, and hence a much smaller sample size, is analysed."
 Fig., Fig.
 2 shows the star formation histories of randomly selected elliptical galaxies in different mass bins and in different environments.," \ref{fig:sfronefile}
 shows the star formation histories of randomly selected elliptical galaxies in different mass bins and in different environments."
 The figure shows that individual star formation histories display a much more ‘bursty’ behaviour than those shown in Fig. 1.., The figure shows that individual star formation histories display a much more `bursty' behaviour than those shown in Fig. \ref{fig:sfr}.
" This reflects our assumption that bulge formation takes place during merger-induced bursts, which naturally gives the star formation histories of individual systems a bursty nature quite different from the smooth history seen for the population average."," This reflects our assumption that bulge formation takes place during merger-induced bursts, which naturally gives the star formation histories of individual systems a bursty nature quite different from the smooth history seen for the population average."
 We will comment more on the implications of this for the scatter of the ages for the model elliptical galaxies in the following section., We will comment more on the implications of this for the scatter of the ages for the model elliptical galaxies in the following section.
" In Fig. 3,,"," In Fig. \ref{fig:sfrhalo},"
" we show the star formation histories again, but this time split into bins of different parentmass."," we show the star formation histories again, but this time split into bins of different parent."
" 'The long dashed, dash-dotted, dashed, and dotted lines are for elliptical galaxies in haloes of mass c1x101°, 1x10!4, 1x101%, and 1x10?Merespectively?."," The long dashed, dash-dotted, dashed, and dotted lines are for elliptical galaxies in haloes of mass $\simeq 1\times10^{15}$, $1\times10^{14}$, $1\times10^{13}$, and $1\times10^{12}\,{\rm M}_{\sun}$."
. Only galaxies with stellar mass larger than 4x10°Mo are used here.," Only galaxies with stellar mass larger than $4\times10^{9}\,{\rm M}_{\odot}$ are used here."
 The solid line shows the average mass—weighted star formation history for all the galaxies in the sample., The solid line shows the average mass–weighted star formation history for all the galaxies in the sample.
 The faster evolution of proto-cluster regions produces star formation histories that peak at higher redshifts for galaxies in more massive haloes., The faster evolution of proto–cluster regions produces star formation histories that peak at higher redshifts for galaxies in more massive haloes.
" Given that galaxies of a fixed stellar mass occur in haloes covering a wide range of masses, it is not surprising that the dependence of the star formation history on halo mass is much weaker than that on galaxy stellar mass."," Given that galaxies of a fixed stellar mass occur in haloes covering a wide range of masses, it is not surprising that the dependence of the star formation history on halo mass is much weaker than that on galaxy stellar mass."
 We now turn to an analysis of the distribution of ages and metallicities of model elliptical galaxies as a function of stellar mass and environment., We now turn to an analysis of the distribution of ages and metallicities of model elliptical galaxies as a function of stellar mass and environment.
" In Fig. 4,,"," In Fig. \ref{fig:formtime},"
 we show the distribution ofthe formation redshifts for model elliptical galaxies., we show the distribution ofthe formation redshifts for model elliptical galaxies.
 We define the formation redshift as the redshift when 50 per cent (or 80 per cent) of the stars that make up the final elliptical galaxy at redshift zero are already formed., We define the formation redshift as the redshift when $50$ per cent (or $80$ per cent) of the stars that make up the final elliptical galaxy at redshift zero are already formed.
" The shaded histograms are for model elliptical galaxies with stellar mass larger than 101! while the open histogram is for all galaxies with secure morphologyMo, (stellar mass larger than 4x10? "," The shaded histograms are for model elliptical galaxies with stellar mass larger than $10^{11}\,{\rm M}_{\sun}$, while the open histogram is for all galaxies with secure morphology (stellar mass larger than $4\times10^{9}\,{\rm M}_{\sun}$ )."
"The figure clearly demonstrates that the stars in more Mo).massive ellipticals are on average older than those in their less massive counterparts, but the scatter of the distribution is rather large and there is a non-negligible fraction of model galaxies whose stars are formed relativelylate."," The figure clearly demonstrates that the stars in more massive ellipticals are on average older than those in their less massive counterparts, but the scatter of the distribution is rather large and there is a non-negligible fraction of model galaxies whose stars are formed relativelylate."
"It is important, however, to distinguish the early times of the stars that make up the elliptical galaxy","It is important, however, to distinguish the early of the stars that make up the elliptical galaxy"
aas published in various LAU. Circulars (nos.,as published in various IAU Circulars (nos.
 7176. 7177. 7179. TISL. 7193. 7203. 7200. 7236. aud 7238). since they provide the best overall coverage throughout the first 3) mouths of the uova.," 7176, 7177, 7179, 7184, 7193, 7203, 7209, 7236, and 7238), since they provide the best overall coverage throughout the first 3 months of the nova."
 We lave supplemented these with magnitude estimates from pre-discovery photographic plates. aud with photometry at Mt. Jolin University observatory (silmartin(1999) aud Gilmore(1999))) vetween Jun 26 aud Jul 14 (a period for which uo visual maguitude estimates are available in IAUCS). although. there could be an offset between visual maguitudes and. photographic or »hotoelectrie measurements.," We have supplemented these with magnitude estimates from pre-discovery photographic plates, and with photometry at Mt. John University observatory \citet{k99} and \citet{g99}) ) between Jun 26 and Jul 14 (a period for which no visual magnitude estimates are available in IAUCs), although there could be an offset between visual magnitudes and photographic or photoelectric measurements."
 lis a very [ast nova: DellaValleetal(1999) have measured the rate of decline of the uova to be οξύ days and (4210 days. aud hence estimated a peak absolute visual magnitude My of —8.7250.2: hiis implies a distauce to the nova of about 2 kpe.," is a very fast nova: \citet{d99} have measured the rate of decline of the nova to be $_2$ =6 days and $_3$ =10 days, and hence estimated a peak absolute visual magnitude $_{\rm V}$ of $-8.7 \pm 0.2$; this implies a distance to the nova of about 2 kpc."
 It is also a neon nova (Woodwardetal1999).. as evidenced by the detection of strong [NelI] 12.515: line.," It is also a neon nova \citep{w99}, as evidenced by the detection of strong [NeII] $\mu$ line."
 The rare briguuess of the nova (the brightest since the advent of imaging X-ray astronomy) las Inacle aa prime target [ο: X-ray observations., The rare brightness of the nova (the brightest since the advent of imaging X-ray astronomy) has made a prime target for X-ray observations.
 Accordingly. by the eud of 1999. thas been observed with (5 times). ((twice).( (once) aud ((ouce.( with three more polntiugs duriug 2000: Starrfieldetal (2000))).," Accordingly, by the end of 1999, has been observed with (5 times), (twice), (once) and (once, with three more pointings during 2000; \citet{s2000}) )."
 Here we coucentrate ou iheΙ aad cddata (sumanarized in Table L: see also Fig., Here we concentrate on the and data (summarized in Table 1; see also Fig.
 1)., 1).
 We also cite the preliminary results of the oobservatious (Orioetal.(1999a) and Orioetal (1999b)))., We also cite the preliminary results of the observations \citet{o99a} and \citet{o99b}) ).
 The oobservatiou (see also the preliminary report by Mukai&Ishida (1999))) was performed between, The observation (see also the preliminary report by \citet{m99a}) ) was performed between
setting the ¢np nass function.,setting the clump mass function.
 However. it raises the prospect that it may be very difficult inceed to measure a chump mass function which is lognormal.," However, it raises the prospect that it may be very difficult indeed to measure a clump mass function which is lognormal."
 If all chip lass functiois appear lognormal. it may be difficult to distinguish trose Whose width and normalization were set bv the plvsics of star formation and those whose characteristic‘ss were set by our data acquisition and analysis techiques.," If all clump mass functions appear lognormal, it may be difficult to distinguish those whose width and normalization were set by the physics of star formation and those whose characteristics were set by our data acquisition and analysis techniques."
 We have investigated the effects on the derived chump mass function of iuase noise. image angular resolution. and two kinds of spatia filtering.," 	We have investigated the effects on the derived clump mass function of image noise, image angular resolution, and two kinds of spatial filtering."
 We have fouud twat adding noise to an image| aud coarseniug its resolutiou to the poiut where the ojects in the image are clearly no longer the precursors of individual stars frequeulv does not cause its lass fiction to become incompatible with the Salpeter stelar IMF., We have found that adding noise to an image and coarsening its resolution to the point where the objects in the image are clearly no longer the precursors of individual stars frequently does not cause its mass function to become incompatible with the Salpeter stellar IMF.
 When the siuulated nass functions are fit with power laws. the cistilition of the power law ex))nents caused by OIC alid degraded resolution mirrors the distribution of measured exponenuts in the stellar IAF.," When the simulated mass functions are fit with power laws, the distribution of the power law exponents caused by noise and degraded resolution mirrors the distribution of measured exponents in the stellar IMF."
 The clamp mass functioi only deviates conclusively from the Salpeter, The clump mass function only deviates conclusively from the Salpeter
and find halos whose locations lie on the light cone emanating from this position.,and find halos whose locations lie on the light cone emanating from this position.
" To do this we use saved checkpoint files at z= 0.0, 0.25, 0.5, 0.75, and 1.0."," To do this we use saved checkpoint files at $z=0.0$ , $0.25$ , $0.5$ , $0.75$, and $1.0$ ."
" Moving outwards in small (Az= 0.05) redshift slices covering the same range, we locate the nearest output (in to each slice and use the center-of-mass peculiar redshift)velocities of the clusters found in that output to estimate new positions in the slice under consideration."," Moving outwards in small $\Delta z = 0.05$ ) redshift slices covering the same range, we locate the nearest output (in redshift) to each slice and use the center-of-mass peculiar velocities of the clusters found in that output to estimate new positions in the slice under consideration."
" We then compute the flux as Ρ/(Απά1), where dz is the luminosity distance of the cluster."," We then compute the flux as $P/(4 \pi d_L^2)$, where $d_L$ is the luminosity distance of the cluster."
 We do this at both 1.4GHz and 150MHz assuming a spectral index of 1.2., We do this at both $1.4 \ghz$ and $150 \mhz$ assuming a spectral index of $1.2$.
" We begin with 10,, where we show the total counts of radio halos at 1.4GHz and 150MHz in the observable universe as a function of flux limit in mJy."," We begin with , where we show the total counts of radio halos at $1.4 \ghz$ and $150 \mhz$ in the observable universe as a function of flux limit in mJy."
" To generate error bars we propagate the lo uncertainties in the derived M, - My,pm and I,—M, fits to generate a minimum andmaximum radiopowerforeach cluster."," To generate error bars we propagate the $1\sigma$ uncertainties in the derived $M_v$ - $M_{v,{\rm DM}}$ and $\Gamma_v-M_v$ fits to generate a minimum andmaximum radiopowerforeach cluster."
 Our number counts arebound by the resolvability limit, Our number counts arebound by the resolvability limit
indicative of a recent novae. so we examine other possible Xray signatures of nuclear reprocessing in BY Cana.,"indicative of a recent novae, so we examine other possible X–ray signatures of nuclear reprocessing in BY Cam."
 The core composition of the accreting white dwarf can either be CO or ONcALe., The core composition of the accreting white dwarf can either be CO or ONeMg.
 For a CO white dwarf. calculations bv Wovetz Prialnik (1997) show that while Nis almos always enhanced in a nova explosion. the O abundance can be almost unallected or strongly depleted. while the heavier elements are not substantially enriched.," For a CO white dwarf, calculations by Kovetz Prialnik (1997) show that while N is almost always enhanced in a nova explosion, the O abundance can be almost unaffected or strongly depleted, while the heavier elements are not substantially enriched."
 This is consisten with the ASCA observations., This is consistent with the ASCA observations.
 Conversely. for an ONeMg white ciwarf core. N/Osx1I requires that the abundances of P and S should also be strongly enhanced. (Politano et a 1995).," Conversely, for an ONeMg white dwarf core, $\le 1$ requires that the abundances of P and S should also be strongly enhanced (Politano et al 1995)."
 Phis tvpe of nova explosion can then be clearly rule out since S is not overabundant (see Section 3.1)., This type of nova explosion can then be clearly ruled out since S is not overabundant (see Section 3.1).
 Thus if a recent novae has enriched the N abundance. it would have to be [rom accretion onto a CO core. although a more likely explanation for anomalous C/N ratios is that the secondary is evolved (Alouchet et al 1997).," Thus if a recent novae has enriched the N abundance, it would have to be from accretion onto a CO core, although a more likely explanation for anomalous C/N ratios is that the secondary is evolved (Mouchet et al 1997)."
 hesically/-— motivated (as. opposed. to phenomolgical) descriptions of the spectrum of BY Cam are given. by Ixallman et al (1996). using the same ASC'A data as studied rere.," Physically motivated (as opposed to phenomolgical) descriptions of the spectrum of BY Cam are given by Kallman et al (1996), using the same ASCA data as studied here."
 However. their conclusions are rather dillerent in that hey propose a strong component at 6.7 keV rom the preshock column. rather than absorption.," However, their conclusions are rather different in that they propose a strong component at 6.7 keV from the pre–shock column, rather than absorption."
 They require this because their assumed continuum form. (a 30 keV. bremsstrahlung) produces copious 6.9 keV line from ike iron. but has a low Llelike iron ion fraction. so canno ooduce much of the observed. 6.7 keV line.," They require this because their assumed continuum form (a 30 keV bremsstrahlung) produces copious 6.9 keV line from H--like iron, but has a low He–like iron ion fraction, so cannot produce much of the observed 6.7 keV line."
 We show tha his conclusion is obviated by using a continuum given by the (physically expected) multitemperature. plasma. emission rom a cooling shock. as this can produce both the 6.7 au 6.9 keV lines.," We show that this conclusion is obviated by using a continuum given by the (physically expected) multi–temperature plasma emission from a cooling shock, as this can produce both the 6.7 and 6.9 keV lines."
 Ixallman et al (1996) discount this possibility due to the weakness of the observed 6.9 keV line compare o that expected from lower temperature plasma with solar abundances., Kallman et al (1996) discount this possibility due to the weakness of the observed 6.9 keV line compared to that expected from lower temperature plasma with solar abundances.
 Εις objection is removed by allowing iron (ane rw other element abundances) to be subsolar., This objection is removed by allowing iron (and the other element abundances) to be sub–solar.
 Kallman e I (1996) also comment that with lower temperature plasma 1ο observed. continuum is too Lat., Kallman et al (1996) also comment that with lower temperature plasma the observed continuum is too flat.
 In our analysis this is resolved. by the inclusion. of complex absorption from the Ολα., In our analysis this is resolved by the inclusion of complex absorption from the column.
 We note that Ixallman ct al (1996) require a pregajiock column of order 107* overlaving the Xray mission region., We note that Kallman et al (1996) require a pre--shock column of order $\sim 10^{23}$ $^{-2}$ overlaying the X–ray emission region.
 Thus their model is not self.consistent as un does not take into account the strong absorption that this predicts., Thus their model is not self–consistent as it does not take into account the strong absorption that this predicts.
 Our model is also inconsistent in not including re [rom the X.ray illuminated. column. but the subsolar iron abundances means that any predicted: iron lx line emission is reduced. by a factor 3. below that. of Ixallman et al (1996) ie. <25 eV equivalent width. which is smaller than the error bars.," Our model is also inconsistent in not including the from the X–ray illuminated column, but the sub–solar iron abundances means that any predicted iron K line emission is reduced by a factor 3 below that of Kallman et al (1996) i.e. $\le 25$ eV equivalent width, which is smaller than the error bars."
 The photoionised column is instead expected to produce the bull of its. emission from continuum/recombination continuumlines at energies below 1 keV. as noted by Ixallman et al (1993) where a large neutral absorption column is required to suppress the strong predicted low energy. emission.," The photo–ionised column is instead expected to produce the bulk of its emission from continuum/recombination continuum/lines at energies below $\sim 1$ keV, as noted by Kallman et al (1993) where a large neutral absorption column is required to suppress the strong predicted low energy emission."
 Ixallman ct al (1996) also speculate on the presence of a rellection component to produce the observed. 6.4. keV line., Kallman et al (1996) also speculate on the presence of a reflection component to produce the observed 6.4 keV line.
 Fheir derived. limit of 0/236 is again assuming solar abundances. without including the inclination ellects. and using the results of early. reflection calculations. which eave EW ~2100/2; eV for a 20 keV. bremsstrahlung illumination.," Their derived limit of $\Omega/2\pi\le 0.6$ is again assuming solar abundances, without including the inclination effects, and using the results of early reflection calculations, which gave EW $\sim
210\Omega/2\pi$ eV for a 20 keV bremsstrahlung illumination."
 More recent Monte.Carlo results (c.g. Cieorge Fabian 1991. Matt. Perola Piro 1991. Van Teesling. lxaastra Heise 1996). scaled to a 2) keV. bremsstrahlung continuum (e.g. Beardmore et al 1995) indicate that even for a lace on. solar abundance slab the line emission is no more than ~1500/2: eV. while for a mean (phase averaged) viewing angle of 60 this reduces to ~1100/2z eV. while abundances of ~0.4. solar reduces it further to 900/25 eV (George Fabian 1991).," More recent Monte–Carlo results (e.g. George Fabian 1991, Matt, Perola Piro 1991, Van Teesling, Kaastra Heise 1996), scaled to a 20 keV bremsstrahlung continuum (e.g. Beardmore et al 1995) indicate that even for a face on, solar abundance slab the line emission is no more than $\sim
150 \Omega/2\pi$ eV, while for a mean (phase averaged) viewing angle of $60^\circ$ this reduces to $\sim 110 \Omega/2\pi$ eV, while abundances of $\sim
0.4\times$ solar reduces it further to $\sim 90
\Omega/2\pi$ eV (George Fabian 1991)."
 Thus the 6.4 keV line strength is easily. consistent with rellection from a shock just. above the white ανα surface ie. with 0/23—I., Thus the 6.4 keV line strength is easily consistent with reflection from a shock just above the white dwarf surface i.e. with $\Omega/2\pi\sim 1$.
 The BY Cam data from ASCA and GINGA give a physically self consistent picture of an accretion column shock. cooling w both evelotron and bremsstrahlung emission.," The BY Cam data from ASCA and GINGA give a physically self consistent picture of an accretion column shock, cooling by both cyclotron and bremsstrahlung emission."
 This multitemperature X.ray continuum illuminates the white dwarf surface. producing a rellection. continuum anc iron luorescence Dine.," This multi–temperature X–ray continuum illuminates the white dwarf surface, producing a reflection continuum and iron fluorescence line."
 Absorption from the preshock column stronely mocdifies the observed spectral form. and is a significant source of uncertainty since it depends on whether he preshock column is uniform. or blobby (ancl hence ionised or nearly neutral). circular. or arelike in crosssection. and. whether there is. radial density structure.," Absorption from the pre–shock column strongly modifies the observed spectral form, and is a significant source of uncertainty since it depends on whether the pre–shock column is uniform, or blobby (and hence ionised or nearly neutral), circular or arc–like in cross--section and whether there is radial density structure."
 We model this absorption by neutral material where the covering fraction is a power law function of the column., We model this absorption by neutral material where the covering fraction is a power law function of the column.
 This could represent a physical situation where the column is blobhy (unionised). with more blobs accreting towards the centre of a circular column.," This could represent a physical situation where the column is blobby (unionised), with more blobs accreting towards the centre of a circular column."
 However. it is more likely that this form merely gives a suitable approximation to a more complex situation. and we urge further theoretical modelling of the preshock How.," However, it is more likely that this form merely gives a suitable approximation to a more complex situation, and we urge further theoretical modelling of the pre–shock flow."
 The multitemperature emission ancl reflection. moclel. together with the power law neutral absorption moclel will be mace publically available in the next release of NSPEC.," The multi–temperature emission and reflection model, together with the power law neutral absorption model will be made publically available in the next release of XSPEC."
 We thank Dave Smith for his help with the GINGA data extraction. and Mark Cropper for useful conversations and the use of his bremsstrahlungevclotron cooling code.," We thank Dave Smith for his help with the GINGA data extraction, and Mark Cropper for useful conversations and the use of his bremsstrahlung–cyclotron cooling code."
 CD acknowledges support from a PPARC Advanced Fellowship. and PM acknowledges support. from. Polish Academy of Sciences and The Roval Society.," CD acknowledges support from a PPARC Advanced Fellowship, and PM acknowledges support from Polish Academy of Sciences and The Royal Society."
 This research. has been supported in part by the Polish KBN erant. 2P03DO01008 and has mace use of data obtained through the High ]5nergy Astrophysics Science Archive research Center Online Service. provided. by the NASA/CGoddard. Space blight Center. and from the Leicester Database and Archive Service," This research has been supported in part by the Polish KBN grant 2P03D01008 and has made use of data obtained through the High Energy Astrophysics Science Archive research Center Online Service, provided by the NASA/Goddard Space Flight Center, and from the Leicester Database and Archive Service"
"requirements for all the GPU stages, the number of samples which can be processed can then be calculated.","requirements for all the GPU stages, the number of samples which can be processed can then be calculated."
" The size of the input and output buffers can be computed by taking the size of the respective largest buffer from the processing stages, which can then be used to compute the number of samples which will fit in memory."," The size of the input and output buffers can be computed by taking the size of the respective largest buffer from the processing stages, which can then be used to compute the number of samples which will fit in memory."
" 'The subband de-dispersion kernel is very similar to the brute force one, the only major change being that not all the channels are summed up to generate the series, and more than one value is generated per input sample."," The subband de-dispersion kernel is very similar to the brute force one, the only major change being that not all the channels are summed up to generate the series, and more than one value is generated per input sample."
" This makes the algorithm less compute intensive and more memory limited (same number of input requests, more output requests)."," This makes the algorithm less compute intensive and more memory limited (same number of input requests, more output requests)."
" However the number of nominal DM values is only a fraction of the total number of DM values, and this greatly reduces the number of calculations which need to be performed by the brute force algorithm."," However the number of nominal DM values is only a fraction of the total number of DM values, and this greatly reduces the number of calculations which need to be performed by the brute force algorithm."
" To test the code, a file containing a pulsed signal was generated using the fake pulsar generator within (Lorimer, http://sigproc.sourceforge.net)."," To test the code, a file containing a pulsed signal was generated using the fake pulsar generator within (Lorimer, http://sigproc.sourceforge.net)."
 The parameters which were used to generate this fake file are listed in table 2.., The parameters which were used to generate this fake file are listed in table \ref{surveyPlanTable}.
" The fake filterbank data are generated as 1024 time-series, one for each frequency channel."," The fake filterbank data are generated as 1024 time-series, one for each frequency channel."
 Each one is made up of a square pulse of height ὃν1024=0.25 and Gaussian noise with mean 0 and standard deviation 1., Each one is made up of a square pulse of height $8\sqrt{1024} = 0.25$ and Gaussian noise with mean 0 and standard deviation 1.
" The S/N of the average simulated pulse, integrated over frequency has a mean value of 8."," The S/N of the average simulated pulse, integrated over frequency has a mean value of 8."
 Brute-force de-dispersion using 1000 DM values with a DM step of 0.1 pccm? was performed., Brute-force de-dispersion using 1000 DM values with a DM step of 0.1 $\text{pc cm}^{-3}$ was performed.
" Figure 5 shows the output of the de-dispersion code, which captures all pulses with S/N greater than 5."," Figure \ref{brute2Figure}
 shows the output of the de-dispersion code, which captures all pulses with S/N greater than 5."
 The performance of the CUDA implementations has been measured., The performance of the CUDA implementations has been measured.
" Fake data is generated in the testing runs themselves, with all the elements initialized to the same value."," Fake data is generated in the testing runs themselves, with all the elements initialized to the same value."
 The time taken to generate and copy the data to and from GPU memory is not included in the timings., The time taken to generate and copy the data to and from GPU memory is not included in the timings.
" Figure 6 shows the performance achieved when de-dispersing with different number of channels, samples and DM values."," Figure \ref{timingGpuFigure} shows the performance achieved when de-dispersing with different number of channels, samples and DM values."
" Different parameter configurations will result in some different optimal combinations, for example, in cases where the number of data partitions to process is exactly divisible by the number of processors and thread blocks being used."," Different parameter configurations will result in some different optimal combinations, for example, in cases where the number of data partitions to process is exactly divisible by the number of processors and thread blocks being used."
" The general tendency is for performance to increase linearly as the number of channels, samples and DM values increases until the maximum GPU occupancy level is reached, after which the behaviour becomes asymptotic."," The general tendency is for performance to increase linearly as the number of channels, samples and DM values increases until the maximum GPU occupancy level is reached, after which the behaviour becomes asymptotic."
" The optimal block size is 128, since fewer threads will result in less latency hiding and more threads will increase scheduling latency without performance benefits."," The optimal block size is 128, since fewer threads will result in less latency hiding and more threads will increase scheduling latency without performance benefits."
" The grid size does not affect performance too much, except for the case where there are too many threads in each block."," The grid size does not affect performance too much, except for the case where there are too many threads in each block."
" As was already stated, the de-dispersion algorithm is memory-bound, and both the GPU and CPU will spend most of their time waiting for data."," As was already stated, the de-dispersion algorithm is memory-bound, and both the GPU and CPU will spend most of their time waiting for data."
" For this reason, the flop rate achieved on the GPU is a small percentage of the theoretical peak for the C-1060, between 80 and 120 Gflops, which comes to about15-20%."," For this reason, the flop rate achieved on the GPU is a small percentage of the theoretical peak for the C-1060, between 80 and 120 Gflops, which comes to about."
". The memory bandwidth achieved within the GPU is about 55 GB/s, which is about of the theoretical peak."," The memory bandwidth achieved within the GPU is about 55 GB/s, which is about of the theoretical peak."
" The same tests were performed on a CPU, specifically on one core of a QuadCore Intel Xeon 2.7 GHz."," The same tests were performed on a CPU, specifically on one core of a QuadCore Intel Xeon 2.7 GHz."
 CPU-performance decreases quasi-linearly as the number of samples or channels increases due to cache misses., CPU-performance decreases quasi-linearly as the number of samples or channels increases due to cache misses.
 This performance is then compared with the appropriate GPU performance to produce the comparison plots in figure 7.., This performance is then compared with the appropriate GPU performance to produce the comparison plots in figure \ref{bruteSpeedupFigure}.
" This shows the speedup gained in performing brute force de-dispersion when using GPUs, for different parameter values."," This shows the speedup gained in performing brute force de-dispersion when using GPUs, for different parameter values."
" From these plots it follows that on average we get a speed of aboutmagnitude, between 50x - 200x depending on the parameters used, with the speedup increasing as the number of input samples/channels increases."," From these plots it follows that on average we get a speed of about, between $\times$ - $\times$ depending on the parameters used, with the speedup increasing as the number of input samples/channels increases."
 The CUDA implementation was then compared to the, The CUDA implementation was then compared to the
Cold Dark Matter (CDAL) cosmogonies foretell the existence of a significant amount of substructure in galaxy halos (IxXIvpin. ct al.,Cold Dark Matter (CDM) cosmogonies foretell the existence of a significant amount of substructure in galaxy halos (Klypin et al.
 1999. Moore. et. al.," 1999, Moore et al."
 1999)., 1999).
 Llowever. the requency of low-mass clark-matter halos around the Milky. Way predicted. in numerical simulations of structure ormation. is an order of magnitude higher than the known number of dwarf galaxies in the Local Group. (Mateo 1998).," However, the frequency of low-mass dark-matter halos around the Milky Way predicted in numerical simulations of structure formation, is an order of magnitude higher than the known number of dwarf galaxies in the Local Group (Mateo 1998)."
" Theoretical solutions of this CDAL ""crisis"" include opositions to change the properties of the clark matter xwiicle (c.g. Sperecl Steinhardt 2000: Colinn. Avila-Loose Valenzucla 2000). as well as proposals which seek o change with cosmic time the astrophysical conclitions hat make gas conducive to star formation in small idos (eg. Bullock. Ixravstov Weinberg. 2001)."," Theoretical solutions of this CDM “crisis"" include propositions to change the properties of the dark matter particle (e.g., Spergel Steinhardt 2000; Colínn, Avila-Reese Valenzuela 2000), as well as proposals which seek to change with cosmic time the astrophysical conditions that make gas conducive to star formation in small halos (e.g. Bullock, Kravstov Weinberg 2001)."
 An alternative solution is the hypothesis that the subhalos wedicted by CDAL simulations have been overlooked so far observationallv., An alternative solution is the hypothesis that the subhalos predicted by CDM simulations have been overlooked so far observationally.
" A potential source for these missing"" dwarf systems is the large population of Compact High-Velocity Clouds (ΗΧΟ)."," A potential source for these “missing"" dwarf systems is the large population of Compact High-Velocity Clouds (CHVC)."
 lligh-Velocitv. Clouds are concentrations. of neutral hvdrogen with extremely high. radial velocities that are inconsistent with Galactic rotation models., High-Velocity Clouds are concentrations of neutral hydrogen with extremely high radial velocities that are inconsistent with Galactic rotation models.
 In. spite of decades. of intense investigation. the nature of νο.," In spite of decades of intense investigation, the nature of HVCs"
Figure 1.,Figure \ref{pzzeta}.
" Thus, for the dataset in which both z and ¢ are available, both the convolution and deconvolution approaches are valid, whether or not the means (or, for that matter, the most probable values) of p(z|¢) and p(¢|z) are unbiased, and however complicated (skewed, multimodal) the shape of these two distributions."," Thus, for the dataset in which both $z$ and $\zeta$ are available, both the convolution and deconvolution approaches are valid, whether or not the means (or, for that matter, the most probable values) of $p(z|\zeta)$ and $p(\zeta|z)$ are unbiased, and however complicated (skewed, multimodal) the shape of these two distributions."
 This remains true in the larger dataset where only ¢ is known., This remains true in the larger dataset where only $\zeta$ is known.
" However, whereas the convolution approach assumes that p(z|¢) is the same in the calibration subset as in the full one, the deconvolution approach assumes that p(C|z) is the same."," However, whereas the convolution approach assumes that $p(z|\zeta)$ is the same in the calibration subset as in the full one, the deconvolution approach assumes that $p(\zeta|z)$ is the same."
" The integral in equation (3)) is really a sum over all the objects in the photometric dataset, where each object with estimated ¢ contributes to dN/dz with weight p(z|¢): Now, recall that ¢ was the mean (or most probable) value of a distribution returned by a photometric redshift code."," The integral in equation \ref{Nz}) ) is really a sum over all the objects in the photometric dataset, where each object with estimated $\zeta$ contributes to ${\rm d}N/{\rm d}z$ with weight $p(z|\zeta)$: Now, recall that $\zeta$ was the mean (or most probable) value of a distribution returned by a photometric redshift code."
" In cases where the observed colours c map to a unique value of 6, then this sum over C is really a sum over c, and the expression above is really Equation (5)) is one of the key results of this paper."," In cases where the observed colours ${\bm c}$ map to a unique value of $\zeta$, then this sum over $\zeta$ is really a sum over ${\bm c}$, and the expression above is really Equation \ref{Npz|c}) ) is one of the key results of this paper."
" Although we arrived at equation (5)) by requiring the mapping c—C be one-to-one (as may be the case for, e.g., LRGs), it is actually more general."," Although we arrived at equation \ref{Npz|c}) ) by requiring the mapping ${\bm c}\to \zeta$ be one-to-one (as may be the case for, e.g., LRGs), it is actually more general."
" This is because one can simply measure p(z|c) in the sample for which spectra are in hand, for the same reason that one could measure p(z|C)."," This is because one can simply measure $p(z|{\bm c})$ in the sample for which spectra are in hand, for the same reason that one could measure $p(z|\zeta)$."
" In fact, p(z|c) is an easier measurement, since it does not depend on the output of a photo-z code!"," In fact, $p(z|{\bm c})$ is an easier measurement, since it does not depend on the output of a $z$ code!"
 The constraint on the mapping between c and ¢ in the discussion above was simply to motivate the connection between photo-z codes and the convolution method., The constraint on the mapping between ${\bm c}$ and $\zeta$ in the discussion above was simply to motivate the connection between $z$ codes and the convolution method.
" Once the connection has been made, however, there is no real reason to go through the intermediate step of estimating C, since all photo-z codes use the observed colors c anyway."," Once the connection has been made, however, there is no real reason to go through the intermediate step of estimating $\zeta$, since all $z$ codes use the observed colors ${\bm c}$ anyway."
" In this respect, equation (5)) is the more direct and natural expression to work with than is equation (4))."," In this respect, equation \ref{Npz|c}) ) is the more direct and natural expression to work with than is equation \ref{Npz|zeta}) )."
" In particular, because p(z|c) is an observable, the convolution approach of equation (5)) is independent of"," In particular, because $p(z|{\bm c})$ is an observable, the convolution approach of equation \ref{Npz|c}) ) is independent of"
masses of the AcBe primarics were not known and we averaged the restuts obtained using a flat distribution of masses ranging from L5AL.. to LOAL...,masses of the AeBe primaries were not known and we averaged the results obtained using a flat distribution of masses ranging from $1.5 M_\odot$ to $10 M_\odot$.
 When the distance to an AcBe star was unknown. we accdlitionally averaged the results obtained by placing the star at distances ranging frou LOOpe and 1500pc.," When the distance to an AeBe star was unknown, we additionally averaged the results obtained by placing the star at distances ranging from 100pc and 1500pc."
 Table 1 shows the thus predicted binary discovery rate of cach of the stars in our sample., Table 1 shows the thus predicted binary discovery rate of each of the stars in our sample.
 By assuming an intrinsic binary frequency ofLOOM... we computed that +1.6 of the companious would be detected or 6.2 40.6 stars for a sample of 39 binary stars.," By assuming an intrinsic binary frequency of, we computed that $\pm$ of the companions would be detected or 6.2 $\pm$ 0.6 stars for a sample of 39 binary stars."
 We also calculated that 0.2 field stars brighter than 10.5 I< maenitude would be observed within 8 (based «on colputations using the model)., We also calculated that 0.2 field stars brighter than 10.5 K magnitude would be observed within 8” (based on computations using the \cite{soniera} model).
 Adjusting for fortuitous field stars. we therefore expected to detect 6.1 £0.6 multiple svsteis if all of the 39 stars m our sample were binaries.," Adjusting for fortuitous field stars, we therefore expected to detect 6.4 $\pm$ 0.6 multiple systems if all of the 39 stars in our sample were binaries."
" The successful detection of 9 multiple systems out of 39 AcBe stars corresponds to a mean ος of coumpanious of 1.1,", The successful detection of 9 multiple systems out of 39 AeBe stars corresponds to a mean number of companions of 1.4.
 Taking iuto account tle nucertaimty in our monte-carlo prediction aud Poisson statistics we calcadated that there is a likelihood that the AcBe star binarity is exeater thanC., Taking into account the uncertainty in our monte-carlo prediction and Poisson statistics we calculated that there is a likelihood that the AeBe star binarity is greater than.
 This is comparable to that of TTS ( (1993).. Chezetal. (1993))) and exceeds that of near solar type AS stars(s (19913)).," This is comparable to that of TTS ( $\pm$ \cite{leinert}, \cite{ghez}) ) and exceeds that of near solar type MS stars, \cite{duq}) )."
" Note that sources were previously detected around some of t1ος AeDe stars (Z CMa (19933)). VW Ser and V376 Cas (Lietal. (199139) for exiuuple). but hat these were not combined with our observations because they were not within the completeness regioi of this study <sep.<<c8"" aud A «10.5)."," Note that sources were previously detected around some of these AeBe stars (Z CMa \cite{haas}) ), VV Ser and V376 Cas \cite{li}) ) for example), but that these were not combined with our observations because they were not within the completeness region of this study $<$ $<$ 8” and $K<$ 10.5)."
 It should be noted that to obtain a significant increase in fhe coufidence level would require a saluple at least four times as huge., It should be noted that to obtain a significant increase in the confidence level would require a sample at least four times as large.
 Such a large sample is not availabe in the northerm hemisphere. but this statistical study can be improved by observing additional AcBe stars in the southern sky.," Such a large sample is not available in the northern hemisphere, but this statistical study can be improved by observing additional AeBe stars in the southern sky."
 Tt is also important to note that we found companions (which are assumed to be the fainter stars) around of the tvpe I AcBe stars iu our sample (IID 200775. BD|16 3171. MWC 1080. KIS Oph. and ΠΟ 150193).," It is also important to note that we found companions (which are assumed to be the fainter stars) around of the type I AeBe stars in our sample (HD 200775, BD+46 3471, MWC 1080, KK Oph, and HD 150193)."
 Iu the case of WD 200775 and BD|16 3171 the companions were found to be significautly £ünter ( AK = L9 and L7 respectively) than the primary star and can oulv account for a small portion of the observed mfrared lieht excess (Eilleubraudetal.(199 1)))., In the case of HD 200775 and BD+46 3471 the companions were found to be significantly fainter ( $\Delta$ K $=$ 4.9 and 4.7 respectively) than the primary star and can only account for a small portion of the observed infrared light excess \cite{hillenbrand}) ).
 The nearer of the two companions around MW 1080 contributes about of the 2.2 pau radiation and we estimate twat the companions around Ik Oph aud ΠΟ 150193 coutribute about and of the total 2.2 sau radiation., The nearer of the two companions around MWC 1080 contributes about of the 2.2 $\mu$ m radiation and we estimate that the companions around KK Oph and HD 150193 contribute about and of the total 2.2 $\mu$ m radiation.
 The companion near V892 Tau. the ouly type II star with a companion in this study. is secu to coitvibute less than of the total 2.2 jan light.," The companion near V892 Tau, the only type II star with a companion in this study, is seen to contribute less than of the total 2.2 $\mu$ m light."
 The three stars that have not previously been classified by 1Iillenbrandetal.(1991). have been found to all have very bright compauions coutributing between aid of the total infraved. light from the svstem., The three stars that have not previously been classified by \cite{hillenbrand} have been found to all have very bright companions contributing between and of the total infrared light from the system.
present.,present.
 In the case of NS system. the presence of a solid surface (reflection boundary condition) plays a erucial role. given that the radiation pressure gradient gets increasing with the accretion rate until bulk flow is stopped.," In the case of NS system, the presence of a solid surface (reflection boundary condition) plays a crucial role, given that the radiation pressure gradient gets increasing with the accretion rate until bulk flow is stopped."
 Thus no index saturation can be observed in NS sources., Thus no index saturation can be observed in NS sources.
 The innermost part of the system evolves towards thermal equilibrium when the emergent spectrum consists of two blackbody-like components which are related to NS surface and disk emissions (?).., The innermost part of the system evolves towards thermal equilibrium when the emergent spectrum consists of two blackbody-like components which are related to NS surface and disk emissions \citep{ts05}.
 On the basis of these considerations we may tentatively discuss the parameters behaviour reported in Tab. 1::, On the basis of these considerations we may tentatively discuss the parameters behaviour reported in Tab. \ref{tab_fit}:
 along the horizontal branch (HB). at lower AI level. BC effect can efficiently pronounced as it is indicated by the high é-value.," along the horizontal branch (HB), at lower $\dot{M}$ level, BC effect can efficiently pronounced as it is indicated by the high $\delta$ -value."
 A slight decrease of the BC effect is observed in the spectra in the low part of HB (LHB). where both 6 and log(.A) get lower. even though this effect seems to be marginal given the high error bars for both these parameters.," A slight decrease of the BC effect is observed in the spectra in the low part of HB (LHB), where both $\delta$ and $\log(A)$ get lower, even though this effect seems to be marginal given the high error bars for both these parameters."
 The total keV) source luminosity at LHB increases about with respect to that at UHB., The total (0.1-200 keV) source luminosity at LHB increases about with respect to that at UHB.
 When the sources is at the UNB (upper normal branch). a further increase (~ 50 %)) of the luminosity with respect to that at LHB occurs. accompanied by significant decrease of the hard X-ray emission (see Fig. 3)).," When the sources is at the UNB (upper normal branch), a further increase $\sim$ 50 ) of the luminosity with respect to that at LHB occurs, accompanied by significant decrease of the hard X-ray emission (see Fig. \ref{residuals}) )."
 We may interpret this behaviour as an increase of Maias which in turn leads to higher AM and consequently leads to the increase of the radiation pressure in the TL., We may interpret this behaviour as an increase of $\dot{M}_{\rm disk}$ which in turn leads to higher $\dot{M}_{\rm tl}$ and consequently leads to the increase of the radiation pressure in the TL.
 This pressure works against an inward matter motion and the bulk effect decreases (though it may be still present but below the instrument threshold)., This pressure works against an inward matter motion and the bulk effect decreases (though it may be still present but below the instrument threshold).
 Figure + illustrates this spectral transition. due to A change. from the intermediate (UHB) to soft (NB) state. following the terminology defined in P006.," Figure \ref{spectral_ev} illustrates this spectral transition, due to $\dot{M}$ change, from the intermediate (UHB) to soft (NB) state, following the terminology defined in P06."
 In the spectrum the power-law extension reduces as a result of the BC efficiency (0. parameter) reduction., In the spectrum the power-law extension reduces as a result of the BC efficiency $\delta-$ parameter) reduction.
 Looking at the parameters of Tab., Looking at the parameters of Tab.
 1. 3 (2). , \ref{tab_fit} \ref{residuals} \citep{is99} 
This is a blurred version of FD defined in Eq. (30)),This is a blurred version of FD defined in Eq. \ref{equ_fd}) )
 and will decrease the different performance of the models but more robustly than FD., and will decrease the different performance of the models but more robustly than FD.
 Numerical comparison results can be found in Table [I]., Numerical comparison results can be found in Table \ref{tab:numbericalcoompare}.
" From the first part of the table (the results of the uniform weighting test), we can see that IUWT-based CS provides the best reconstruction results for both FD and CFD. figureB]."," From the first part of the table (the results of the uniform weighting test), we can see that IUWT-based CS provides the best reconstruction results for both FD and CFD. \ref{fig_none}."
surface reactions become important.,surface reactions become important.
 As a result. a large fraction of oxvgen is locked iu water ice that is hard to desorb (Law~5500 I). while half of the clemeutal carbo eoes to volatile micthane ice (E~1300 I).," As a result, a large fraction of oxygen is locked in water ice that is hard to desorb $E_{\rm des} \sim 5500$ K), while half of the elemental carbon goes to volatile methane ice $E_{\rm des} \sim 1300$ K)."
 Upon CRP heating of dust eraius. this leads to πιο higher gas-pliase abundance of CSIT in he cloud core or the cold model compared to the wari. uodoel.," Upon CRP heating of dust grains, this leads to much higher gas-phase abundance of $_2$ H in the cloud core for the cold model compared to the warm model."
 The effectis not hat stroug for less clense regions a larger radi from the cOuter., The effect is not that strong for less dense regions at larger radii from the center.
" Since the C5IT cussion is auti-correated wih the dus continuun emission iu he ¢case of EE8089-1732 roflsü0s89)). we* do uot have the Il, column deusities to quantitatively conpare the abudace profiles. of IRASIISOSO-1732 wit1 our model."," Since the $_2$ H emission is anti-correlated with the dust continuum emission in the case of 18089-1732 \\ref{18089}) ), we do not have the $_2$ column densities to quantitatively compare the abundance profiles of 18089-1732 with our model."
 However. data ane model allow a qualitaive cOparison of hne spatia structures.," However, data and model allow a qualitative comparison of the spatial structures."
 Estimating an exact CVOutiouary time for IRAS1ISO0S9-1132 Is lhiucdly »ossible. but based on the strong molecular ne CUSSION. its high centra eas tfeniperatires aud fle (observed outfiow-disk. svsteni (Deutheretal.2001b.a.2YOST). an approxinate age of 5«10! yr appears reasonho.," Estimating an exact evolutionary time for 18089-1732 is hardly possible, but based on the strong molecular line emission, its high central gas temperatures and the observed outflow-disk system \citep{beuther2004a,beuther2004b,beuther2005c}, an approximate age of $5\times10^4$ yr appears reasonable."
 Although dyuuuica ancl chemical les are iot necessarily exactly the same. in highauass star formatio itrev should not differ to much: Followiie the models by Melee&Tan(2003) Or Ixiuhlolzeta.(2007).. he Tuninositv rises stronelv right from the ouse of collapse which cau be considered as a starting poiut or the chenical evolution.," Although dynamical and chemical times are not necessarily exactly the same, in high-mass star formation they should not differ to much: Following the models by \citet{mckee2003} or \citet{krumholz2006b}, the luminosity rises strongly right from the onset of collapse which can be considered as a starting point for the chemical evolution."
" At the same time disks aud outfhows evolve. which should lence have similar time-scaOs,"," At the same time disks and outflows evolve, which should hence have similar time-scales."
 Tre diaueter of the shell-like CoIT structure in 115089-1732 21805rif rof] SOUSO3). or ~9000 AAT in radius at the eiven distance of kkpc.," The diameter of the shell-like $_2$ H structure in 18089-1732 is $\sim 5''$ \\ref{18089}) ), or $\sim$ AU in radius at the given distance of kpc."
 This value is well matched by the modeled region with decreased. Coll abundance refinedel))., This value is well matched by the modeled region with decreased $_2$ H abundance \\ref{model}) ).
 Although iu principle optical depths and/or excitation effects could ninie the Col morphology. we consider this as unlikely because the other observed lutdecules with many different transitions all peak toward the central subi coutiuuuu cussion in L1sos9-lv:33 (Beutheretal.2005b). ," Although in principle optical depths and/or excitation effects could mimic the $_2$ H morphology, we consider this as unlikely because the other observed molecules with many different transitions all peak toward the central submm continuum emission in 18089-1732 \citep{beuther2005c}. ."
Since Col is the ouly exception iu that rich dataset. chemical effects appear the lutwe plausihle explanation.," Since $_2$ H is the only exception in that rich dataset, chemical effects appear the more plausible explanation."
 The fact that we see oll at the earliest aud the later οτεdutionary stages can be explained by the reactive nature of (ΟΠ: it is produced quickly carly on aud ects replenished at the core edges by the UV photodissociation of CO., The fact that we see $_2$ H at the earliest and the later evolutionary stages can be explained by the reactive nature of $_2$ H: it is produced quickly early on and gets replenished at the core edges by the UV photodissociation of CO.
 The inner ~chemical” hole observed toward LI8089-1732 can be explained by CI eine consumed iu the chemical network forming CO and more complex molecules like larecr carbon-hivdrogeu complexes and/or depletion., The inner “chemical” hole observed toward 18089-1732 can be explained by $_2$ H being consumed in the chemical network forming CO and more complex molecules like larger carbon-hydrogen complexes and/or depletion.
" T1ο data show that CoIT is uot suited to investigate the central gas cores in more evolved sources. however. our analysis medicates that Col may be a suitable tracer of the earliest stages of (luassive) star formation. like N2 II! or NIL, (e.g. Bereinetal2002:Tafallaal.200Beutheretal.2005a:Pillai 20061)."," The data show that $_2$ H is not suited to investigate the central gas cores in more evolved sources, however, our analysis indicates that $_2$ H may be a suitable tracer of the earliest stages of (massive) star formation, like $_2$ $^+$ or $_3$ (e.g., \citealt{bergin2002,tafalla2004,beuther2005a,pillai2006}) )."
 While a spatial analysis of the line cussion will eive insights iuto he οποιαιός of the eas and also the evolutionary stage roni chemical models. multiple Col lines will even allow ate uperature characterization.," While a spatial analysis of the line emission will give insights into the kinematics of the gas and also the evolutionary stage from chemical models, multiple $_2$ H lines will even allow a temperature characterization."
" With its lowest Jὐ ransitious around GCGIIz. CoM has easily accessible spectral lines in several bauds between the 21m and yan. Furthermore. even the GGIIz lines preseuted iere have still relatively low upper level excitation cucreics GE,ko 12«I). hence allowing to studycold cores even at sub-anilhineter waveleugths."," With its lowest $J=1-0$ transitions around GHz, $_2$ H has easily accessible spectral lines in several bands between the mm and $\mu$ m. Furthermore, even the GHz lines presented here have still relatively low upper level excitation energies $E_u/k\sim42$ K), hence allowing to studycold cores even at sub-millimeter wavelengths."
 This prediction can urther be proved via Ligh spectral aud sxitial resolution observatious of different ΟΠ lines towiux voung IRDC., This prediction can further be proved via high spectral and spatial resolution observations of different $_2$ H lines toward young IRDCs.
Mass outflows play an essential role in the process of star formation and have been found to be ubiquitous in various stages of star formation(Arceetal. 2001).,"Mass outflows play an essential role in the process of star formation and have been found to be ubiquitous in various stages of star formation\citep{arc07,bal07}."
. The specilic angular momentum of a star-forming molecular core is about four orders of magnitude (mes higher (han that of a (vpical T Tawi star (Boclenheimer 1995).., The specific angular momentum of a star-forming molecular core is about four orders of magnitude times higher than that of a typical T Tauri star \citep{bod95}. .
 One, One
Over the last decade. novel discoveries of transiting extrasolar plaucts have often been fraught by nuexplaimed radius anomalies.,"Over the last decade, novel discoveries of transiting extra-solar planets have often been fraught by unexplained radius anomalies."
" Iu particular. many eas eiut planets. residing on orbits in extreme proximuty to their host stars. have been found to have much larecr radi than what was previously thought possible iu the contest of ""standard gas elaut theory (Stevenson1982:Caullot 1999).."," In particular, many gas giant planets, residing on orbits in extreme proximity to their host stars, have been found to have much larger radii than what was previously thought possible in the context of “standard"" gas giant theory \citep{1982AREPS..10..257S,1999Sci...286...72G}."
 A nuuber of possible explanations to this problem have been proposed. most notably tidal heating (Bodenheimeretal.2001.2003).. kinetic heating due to breakingC» C»eravitv waves (Caullot&Showman2002). chhanced opacity: (Burrowsetal.0n2007).. se1ü1-:convection⋅ (Chabricr:&Daraffeto2007).MOT and turbulent burial: of: heat by a mechauical ereenhouse. effect.: (Youdinr:&Mitchell 2010)..," A number of possible explanations to this problem have been proposed, most notably tidal heating \citep{2001ApJ...548..466B,2003ApJ...592..555B}, kinetic heating due to breaking gravity waves \citep{2002A&A...385..156G}, enhanced opacity \citep{2007ApJ...661..502B}, semi-convection \citep{2007ApJ...661L..81C} and turbulent burial of heat by a mechanical greenhouse effect \citep{2010ApJ...721.1113Y}."
 IHTowever. it appears unlikely that auy of the above solutions can simultaneously explain all of the observed auonmales (Fortuev&Nettelmaun2010)..," However, it appears unlikely that any of the above solutions can simultaneously explain all of the observed anomalies \citep{2010SSRv..152..423F}."
 Receuth. a new Olunic heating mechanisin. that relies ou the electro-magnetic interactions between atmospheric flows aud the planetary magnetic field. has been suggested» ax a pronusing∙∙↜ explanation∙ to the inflation⋅ problem (Batvein)&V|Stevensonson2010)2," Recently, a new Ohmic heating mechanism, that relies on the electro-magnetic interactions between atmospheric flows and the planetary magnetic field, has been suggested as a promising explanation to the inflation problem \citep{2010ApJ...714L.238B}."
" "" A characteristic that is unique to hot Jupiters is the enornous amount of incident enerev that they receive from their host stars.", A characteristic that is unique to hot Jupiters is the enormous amount of incident energy that they receive from their host stars.
 The extreme irradiation naturally results iu high atinospheric temperatures (sometimes m excess200018) (Spiegeletal.2009).. and the latitudinal dependenceof of this heating leads to fast. super-rotatiug equatorial jets with characteristic wiud-speeds of order ~ 1 dm/s (Showmanctal.2009)...," The extreme irradiation naturally results in high atmospheric temperatures (sometimes in excess of 2000K) \citep{2009ApJ...699.1487S}, and the latitudinal dependence of this heating leads to fast, super-rotating equatorial jets with characteristic wind-speeds of order $\sim$ 1 km/s \citep{2009ApJ...699..564S}."
" Consequently, the atmospheres of hot Jupiters are hot enough to thermally ionize alkali metals. that are preseut in trace abuidauces iand eeive tise: fo ia finite electricali conductivitv of order ~1031 ↽⊳↜↽≼∣πι, depending⋅ primarily⋅ on the effective⋅⋅ temperature."," Consequently, the atmospheres of hot Jupiters are hot enough to thermally ionize alkali metals, that are present in trace abundances, and give rise to a finite electrical conductivity of order $\sim 10^{-3} - 1$ S/m, depending primarily on the effective temperature."
 Rapid advection of the ious bv the jets. while immersed: in the planetaryian magneticae Beld.ο leadslead: to an induction: ofd an electroauotive: forceJg that sets up electrical current loops through the deep interior audthe atinosplere of the plauct.," Rapid advection of the ions by the jets, while immersed in the planetary magnetic field, leads to an induction of an electro-motive force that sets up electrical current loops through the deep interior and the atmosphere of the planet."
 These electrical cimveuts. eive rise to Oluuic heating as they flow throughout the planet., These electrical currents give rise to Ohmic heating as they flow throughout the planet.
 It has been shown. iu the context of a static structural model. that the level of resulting dissipation in the interior is approximately that required to maintain the radi of hot Jupiters (Batvein&Stevenson2010).," It has been shown, in the context of a static structural model, that the level of resulting dissipation in the interior is approximately that required to maintain the radii of hot Jupiters \citep{2010ApJ...714L.238B}."
 llowever. time-evolution of Olinically heated gas-eiauts relmains au open question. since the degree of dissipation exlibits strong dependence on the structure of the planet. and particularly on the planct’s ἆρεε mass. aud the location of the radiative/convective bomudary.," However, time-evolution of Ohmically heated gas-giants remains an open question, since the degree of dissipation exhibits strong dependence on the structure of the planet, and particularly on the planet's $T_{eff}$, mass, and the location of the radiative/convective boundary."
 More explicith. the thermal structure can be affected bv the heating and its radial variation as well as the ↸∖⋪↴∖↴↸∖≼ap tuuBe.," More explicitly, the thermal structure can be affected by the heating and its radial variation as well as the elapsed time."
" d thο saluean nnBe. ""nthe ‘conductivityemtyBM ""arises. fron)Ὁ 4ΤΟΥOrlnal1 AtTOMZATION,olüzatiolnn"" CAUSAcase THCe mtOrrorπου structurestruetus ο toflback feedbackοou the heatingheating."," At the same time, the conductivity arises from thermal ionization, causing the interior structure to feedback on the heating."
 ThinsThus. ththe «πιο pi'ofileil of the planet dictates the Olunic heatingstructural distribution.," Thus, the structural profile of the planet dictates the Ohmic heating distribution."
 As a result. selfconsistent evolutionary calculations are required to assess whether Olunic dissipation can indeed resolve the inflation problem.," As a result, self-consistent evolutionary calculations are required to assess whether Ohmic dissipation can indeed resolve the inflation problem."
 To perform such calculations is the primary purpose of this study., To perform such calculations is the primary purpose of this study.
" Tere. we construct an approximate model with specific assumptions. that couples calculations of Olunic heating⋅ and structural evolution.⋅ and show how this⋅ . clement . . ⋅⋅⋅ ∐↸∖⋜↧↑↕∐∶↴∙⊾↑∪∶↴∙⊾↸∖↑∐↸∖↥⋅↖↖↽↕↑↕↑∐↸∖↕∐↘↽↸∖↕⋅↖↽↖⇁⋜∐⋅↕⋜∏⋝↕∐↑⋅↖⇁∪↕↕∐∖⋜↧↖↽⋅↖↽ abuudances cau provide an explanation for all ob northe cumHUY1 ορ] radius7 ANOMANCS,|"," Here, we construct an approximate model with specific assumptions, that couples calculations of Ohmic heating and structural evolution, and show how this heating together with the likely variability of heavy element abundances can provide an explanation for all of the currently observed radius anomalies."
 E‘ compute the ANANTH νοαattainable bx evolved Tot Jupiters as functions of radiplanetars aud effective aud fine tendency u TassDN radii x PMtemperatures 1 COPFESPONK to the onset © 1sieuificautect coniity that. from metals in the atmosphere., We compute the maximum radii attainable by evolved Hot Jupiters as functions of planetary mass and effective temperature and find a tendency for larger radii for effective temperatures that correspond to the onset of significant conductivity arising from alkali metals in the atmosphere.
 This uctisIchavior is arisingconatiile alkaliwith the trend of the data., This behavior is compatible with the trend of the data.
" nrtherukme. our suggest that in absence of a relatively massive. caleulatiousligh-ietallicity core. oN“aeDuas hot. ∙:wu Ga) expandow): INTOcvond tlvh: RoxRochehe-lobes aud spill their Tupiterscuvelopescan onto their parent stars ou billion year time-scales. possibly leaving behind small 1""OCIXV. cotCO"," Furthermore, our calculations suggest that in absence of a relatively massive, high-metallicity core, low-mass hot Jupiters can expand beyond their Roche-lobes and spill their envelopes onto their parent stars on billion year time-scales, possibly leaving behind small rocky cores."
YCS. paper Ts --struetüpec.] as follTOLOWS., The paper is structured as follows.
 1Hp SCCTIOMü 52. WX consideri the euergeties of the Oluuic mechanisin. aud show that in steady state. the degree of dissipation is lanited by the thermocdvuamic efficiency of the," In section 2, we consider the energetics of the Ohmic mechanism, and show that in steady state, the degree of dissipation is limited by the thermodynamic efficiency of the"
siineq. Ecgc110-152007).,"0, $E_{\rm cut}$$\sim$ 11-15."
 The absorption column density was —0.1 < 10°? 7. compatible with the interstellar value in. the direction of the source (Dickey&Lockman.1990).," The absorption column density was $\sim$ $\times$ $^{22}$ $^{-2}$, compatible with the interstellar value in the direction of the source \citep{dickey90}."
 The 0O.1- keV luminosity of the two outbursts was « 10°? and 3.1079 erg/s. wwas also caught in outburst by 44 times between 2006 and 2009 (Romanoetal..2008;Sidolietal..2009b:Barthelmy 2009).," The 0.1-100 keV luminosity of the two outbursts was $\times$ $^{35}$ and $\times$ $^{36}$ erg/s. was also caught in outburst by 4 times between 2006 and 2009 \citep{romano08,sidoli09b,barthelmy09}."
. During the first outburst. which occurred on 2006 October 4. JIXRT slewed to the source ~2000 s after the BAT trigger.," During the first outburst, which occurred on 2006 October 4, /XRT slewed to the source $\sim$ 2000 s after the BAT trigger."
 The combined XRT+BAT spectrum could be reproduced well by using à cut-off power-law model (T=0.31. ΓΗ keV).," The combined XRT+BAT spectrum could be reproduced well by using a cut-off power-law model $\Gamma$ =0.31, $E_{\rm c}$$\simeq$ 11 keV)."
 The column density was 70.3 « 1077 7., The column density was $\sim$ $\times$ $^{22}$ $^{-2}$.
 The second outburst. which took place on 2008 July 5. was characterized. by complex behavior and comprised two different flares separated by ~10° s. The time-resolved spectral analysis found that the soft X-ray emission of the source could be reasonably well described by using an absorbed power-law model with a constant photon index of ~1 and a variable absorptior column density ranging from ~0.5 to —11 4 10?? 7.," The second outburst, which took place on 2008 July 5, was characterized by complex behavior and comprised two different flares separated by $\sim$ $^{5}$ s. The time-resolved spectral analysis found that the soft X-ray emission of the source could be reasonably well described by using an absorbed power-law model with a constant photon index of $\sim$ 1 and a variable absorption column density ranging from $\sim$ 0.5 to $\sim$ $\times$ $^{22}$ $^{-2}$."
 The simultaneous XRT+BAT broad band (0.3-80 keV) spectrum of the source could be reproduced well using an absorbed cut-off power-law model. with ΓΙ keV. D-—l.4. and Aq26. < 107? 7.," The simultaneous XRT+BAT broad band (0.3-80 keV) spectrum of the source could be reproduced well using an absorbed cut-off power-law model, with $E_{\rm c}$$\gtrsim$ 14 keV, $\Gamma$$\sim$ 1.4, and $N_{\rm H}$ $\times$ $^{22}$ $^{-2}$."
 The third outburst occurred on 2008 September 21., The third outburst occurred on 2008 September 21.
 On this occasion. JIXRT slewed to the source  150 s after the BAT trigger. a- time-resolved spectroscopy did not reveal any variation in the absorption column density.," On this occasion, /XRT slewed to the source $\sim$ 150 s after the BAT trigger, and time-resolved spectroscopy did not reveal any variation in the absorption column density."
 However. the total XRT spectrum of the observation (exposure time 1160 s) clearly. revealec a previously undetected soft component below ~2 keV. A reasonably good fit to these data could be obtained by using an absorbed CVjj20.4 < 107? 2) power-law (D2.2) component and a BB at the softer energies (with à temperature of AkT-1.95 keV. and a radius of ~1.2 km).," However, the total XRT spectrum of the observation (exposure time 1160 s) clearly revealed a previously undetected soft component below $\sim$ 2 keV. A reasonably good fit to these data could be obtained by using an absorbed $N_{\rm H}$ $\times$ $^{22}$ $^{-2}$ ) power-law $\Gamma$$\simeq$ 2.2) component and a BB at the softer energies (with a temperature of $kT$$\sim$ 1.95 keV, and a radius of $\sim$ 1.2 km)."
 This BB component was also required to fit the combined XRT-BAT broad band spectrum of the outburst (a single comptonization model. BMC. only provided a poor fit to the data).," This BB component was also required to fit the combined XRT+BAT broad band spectrum of the outburst (a single comptonization model, BMC, only provided a poor fit to the data)."
 Sidolietal.(2009b) found that the seed photon temperature of the BMC component and the BB temperature could not be linked to the same value in the fit., \citet{sidoli09b} found that the seed photon temperature of the BMC component and the BB temperature could not be linked to the same value in the fit.
 This was interpreted as being caused by the presence of two distinct photon populations. a colder one with a temperature of about 0.3 keV and a radius of ~10 km. and a hotter one with a temperature of 1.4-1.8 keV and a radius," This was interpreted as being caused by the presence of two distinct photon populations, a colder one with a temperature of about 0.3 keV and a radius of $\sim$ 10 km, and a hotter one with a temperature of 1.4-1.8 keV and a radius"
uncertainties in the SDSS ABcalibration.,uncertainties in the SDSS AB.
". In the near-infrared, there is a tendency for the model to slightly overestimate the observed K-band luminosity (by an amount corresponding to about 0.10 mag)."," In the near-infrared, there is a tendency for the model to slightly overestimate the observed $K_s$ -band luminosity (by an amount corresponding to about $0.10$ mag)."
 This offset is well within the uncertainties of current stellar population synthesis models for this spectral region., This offset is well within the uncertainties of current stellar population synthesis models for this spectral region.
 We have checked that these offsets have a negligible influence on our results., We have checked that these offsets have a negligible influence on our results.
" To illustrate the quality of these fits, we show in Fig."," To illustrate the quality of these fits, we show in Fig."
" 4 three examples of the best-fit spectral energy distributions of galaxies spanning wide ranges in star formation and dust properties in our sample, from quiescent, moderately dusty (IRAS F15105+5959, top panel) to actively star-forming, highly dusty (IRAS F09253+1724, bottom panel)."," \ref{fig:dust4}
 three examples of the best-fit spectral energy distributions of galaxies spanning wide ranges in star formation and dust properties in our sample, from quiescent, moderately dusty (IRAS F15105+5959, top panel) to actively star-forming, highly dusty (IRAS F09253+1724, bottom panel)."
 The middle panel shows the spectral energy distribution of a galaxy corresponding roughly to the median star formation and dust properties of the sample (IRAS F15028+0820)., The middle panel shows the spectral energy distribution of a galaxy corresponding roughly to the median star formation and dust properties of the sample (IRAS F15028+0820).
" The method described in Section 3 above allows us to derive statistical constraints on the star formation rate averaged over the last 10° yr,w,, the stellar mass,M.,, the total dust luminosity,Líi?*. the fraction of this contributed by dust in the diffuse interstellar medium,f,,., and the total dust mass,Ma., for the 3258 galaxies in the matched ssampledescribed in Section 2.."," The method described in Section \ref{dust:model} above allows us to derive statistical constraints on the star formation rate averaged over the last $10^8$ yr, the stellar mass, the total dust luminosity, the fraction of this contributed by dust in the diffuse interstellar medium, and the total dust mass, for the 3258 galaxies in the matched sampledescribed in Section \ref{dust:sample}."
" In Fig. 5,,"," In Fig. \ref{fig:dust5},"
 we show the resulting relation between dust mass aand star formation rate~., we show the resulting relation between dust mass and star formation rate.
". The grey contour shows the distribution for the full sample, while grey points show the distribution of a sub-sample of 1658 galaxies with highest-S/N photometry at ultraviolet, optical and infrared wavelengths."," The grey contour shows the distribution for the full sample, while grey points show the distribution of a sub-sample of 1658 galaxies with highest-S/N photometry at ultraviolet, optical and infrared wavelengths."
 This high-S/N subsample includes only galaxies with relative photometric errors less than 2c larger than the sample mean in each band., This high-S/N subsample includes only galaxies with relative photometric errors less than $\sigma$ larger than the sample mean in each band.
 Also shown in the lower-right corner of Fig., Also shown in the lower-right corner of Fig.
" 5 are the median errors in aandMa,, derived from the likelihood distributions of these parameters for all the galaxies in the sample."," \ref{fig:dust5} are the median errors in and, derived from the likelihood distributions of these parameters for all the galaxies in the sample."
" The correlation between dust mass and star formation rate is remarkably tight in Fig. 5,,"," The correlation between dust mass and star formation rate is remarkably tight in Fig. \ref{fig:dust5},"
 spanning 4 orders of magnitude in both aandMa., spanning 4 orders of magnitude in both and.
". We perform a linear fit to the grey points (i.e., the highest-S/N subsample) by computing the bisector of two square regression lines (?):: one in aas a function ofw,, and one in aas a function of Ma."," We perform a linear fit to the grey points (i.e., the highest-S/N subsample) by computing the bisector of two least-square regression lines \citep{Isobe1990}: one in as a function of, and one in as a function of ."
. This gives (black line in Fig. 5)), This gives (black line in Fig. \ref{fig:dust5}) )
1n each of these integrals. the factor can be taken out.,"In each of these integrals, the factor can be taken out."
 We are then left with the following three building blocks: The binary LOSVD. averaged. over all orbital parameters. finally takes the following form: As can be expected. the LOSVD goes to zero for or equivalently for the highest line-ol-sieht velocity any binary can reach.," We are then left with the following three building blocks: The binary LOSVD, averaged over all orbital parameters, finally takes the following form: As can be expected, the LOSVD goes to zero for $m_3 \ge m_2$ or equivalently for the highest line-of-sight velocity any binary can reach."
 As long as my«mj. or equivalently for the LOSVD is velocity independent.," As long as $m_0 < m_1$, or equivalently for the LOSVD is velocity independent."
 This is the highest velocity that the slowest binaries. i.e. the ones with the owest secondary mass and the largest orbital radius. can reach.," This is the highest velocity that the slowest binaries, i.e. the ones with the lowest secondary mass and the largest orbital radius, can reach."
 Usually. rias is very large ancl consequently. ey will rc à very small velocity. of the order ofa few m/s (see Figure 1).," Usually, $r_{\rm max}$ is very large and consequently, $v_1$ will be a very small velocity, of the order of a few m/s (see Figure \ref{losvdps}) )."
 lo be exact. the binary LOSVD (39)) still. needs o be multiplied. with the probability that the primary star is actually observed. and. then be averaged over the ηαν mass.," To be exact, the binary LOSVD \ref{losvd}) ) still needs to be multiplied with the probability that the primary star is actually observed and then be averaged over the primary mass."
 Lt is this final function that quantifies the extra broadening of the absorption lines due to the binary x»pulation., It is this final function that quantifies the extra broadening of the absorption lines due to the binary population.
 Llowever. the probability of observing a star of a particular spectral type depends both on its intrinsic xieghtness and on the number of such stars that are around.," However, the probability of observing a star of a particular spectral type depends both on its intrinsic brightness and on the number of such stars that are around."
 As the brightest stars are also the most scarce ones while he faintest stars are extremely abundant. this probability will be sharply peaked.," As the brightest stars are also the most scarce ones while the faintest stars are extremely abundant, this probability will be sharply peaked."
 For old to intermediate stellar populations. this peak corresponds to early Ix. giants on the ip of the giant branch.," For old to intermediate stellar populations, this peak corresponds to early K giants on the tip of the giant branch."
 We will therefore not do this last integration ancl instead usc an average value for the primary mass., We will therefore not do this last integration and instead use an average value for the primary mass.
 Most authors use Ad=0.8M. and this is the value we will adopt here.," Most authors use $M=0.8\,{\rm M}_\odot$ and this is the value we will adopt here."
 The binary LOSVD for elliptic orbits is not analytically tractable. but its velocity moments. which we denote by fi. are.," The binary LOSVD for elliptic orbits is not analytically tractable, but its velocity moments, which we denote by $\tilde{\mu}_n$, are."
" We set out with the expression (15)) for the moments of the specific binary LOSVD and average it over the argument of the periastron z. the inclination 7. the secondary. mass m. the semi-major axis e, and the eccentricitv 6."," We set out with the expression \ref{mome}) ) for the moments of the specific binary LOSVD and average it over the argument of the periastron $\omega$ , the inclination $i$, the secondary mass $m$, the semi-major axis $a_p$ and the eccentricity $e$."
 For { oand m. one can use the distributions (19)) and. (24)).," For $i$ and $m$, one can use the distributions \ref{disi}) ) and \ref{dism}) )."
 The distribution (23)) for the semi-major axis is also still applicable but care has to be taken with the integration limits: the bounds for αν are set by the conditions which makes the integration limits of αν dependent of , The distribution \ref{disap}) ) for the semi-major axis is also still applicable but care has to be taken with the integration limits: the bounds for $a_p$ are set by the conditions which makes the integration limits of $a_p$ dependent of $e$.
"As can beseen in Figure 5.. the distribution of the eccentricity. of a sample of 200. visual. binarics can be approximated fairly accurately by with a=05. 3=| and maximum eccentricitvy Cua,=mlo"," As can beseen in Figure \ref{soee}, , the distribution of the eccentricity of a sample of 200 visual binaries can be approximated fairly accurately by with $\alpha=0.5$, $\beta=1$ and maximum eccentricity $e_{\rm
 max}=1$."
 The eccentricity distribution found by Duquennov Mayor (1991). shows that tight. binarics have more circular orbits than wide ones but. overall. their eccentricity. distribution is similar to the one derived. by Soederhjelm.," The eccentricity distribution found by Duquennoy Mayor \cite{dm} shows that tight binaries have more circular orbits than wide ones but, overall, their eccentricity distribution is similar to the one derived by Soederhjelm."
 In theremainder. we consider the parameters a. 3 andὅμως bo be fixed by the observations.," In theremainder, we consider the parameters $\alpha$ , $\beta$ and$e_{\rm max}$ to be fixed by the observations."
 This distribution reaches a maximum for, This distribution reaches a maximum for
"where n4 is the number density of dust grains. o4 is their mean geometric cross-section. Py, is ihe mean velocity of the hydrogen5 atoms striking5 the 5erains. 9 is the sticking5 coefficient (the xobabilitv that a hvdrogen5 atom striking5 the 5grain. will stick to the surface) and Jcf, is the fraction of adsorbed hydrogen atoms that actually form I». rather (han simply escaping back into the eas phase.","where $n_{\rm d}$ is the number density of dust grains, $\sigma_{\rm d}$ is their mean geometric cross-section, $\bar{v}_{\mH}$ is the mean velocity of the hydrogen atoms striking the grains, $S$ is the sticking coefficient (the probability that a hydrogen atom striking the grain will stick to the surface) and $f_{a}$ is the fraction of adsorbed hydrogen atoms that actually form $\mHt$, rather than simply escaping back into the gas phase."
" Thev argued (hat lor gas ancl grain temperatures (vpical of molecular clouds. both S and f, should be of order unit."," They argued that for gas and grain temperatures typical of molecular clouds, both $S$ and $f_{a}$ should be of order unity."
" llollenbach&AIckee(1979). later used this prescription to derive an IH» formation rate for the local ISM that continues to be widely cited: where and where 75 and 75,» are the gas and grain temperatures in units of 100Ix."," \citet{hm} later used this prescription to derive an $\mHt$ formation rate for the local ISM that continues to be widely cited: where and where $T_{2}$ and $T_{{\rm gr},2}$ are the gas and grain temperatures in units of $100 \: \rm{K}$."
" They argued that f, should be approximately constant. ancl of order unity for erain temperatures below some critical value τω, but that [or Ta>Tj. 1t should fall off exponentially. with most of the hydrogen atoms evaporating from (he erain surlace before (μον have time to form Ils."," They argued that $f_{a}$ should be approximately constant and of order unity for grain temperatures below some critical value $T_{\rm cr}$, but that for $T_{\rm gr} > T_{\rm cr}$, it should fall off exponentially, with most of the hydrogen atoms evaporating from the grain surface before they have time to form $\mHt$."
" The value of Z;, has proved hard to determine precisely. but is of the order of 100Ix."," The value of $T_{\rm cr}$ has proved hard to determine precisely, but is of the order of $100 \: \rm{K}$."
 Although this rate has been widely adopted in the literature. recent experiments have cast doubt on its validity at hieh temperatures. aud suggest that the II» formation rate may be smaller (han previously assumed (Pirronello 1997a.b. 1999: INatz 1999: Bihan 2001).," Although this rate has been widely adopted in the literature, recent experiments have cast doubt on its validity at high temperatures, and suggest that the $\mHt$ formation rate may be smaller than previously assumed (Pirronello 1997a,b, 1999; Katz 1999; Biham 2001)."
" ILowever. since (his conclusion is not entirely clear and their work is still ongoing. I have tentatively adopted the Hollenbach&AIckee rate below. with the proviso that the values of D, (that I derive may prove to be lower limits if the results of Pirronello are borne out by [uture work."," However, since this conclusion is not entirely clear \citep{tiel} and their work is still ongoing, I have tentatively adopted the \citeauthor{hm} rate below, with the proviso that the values of ${\cal D}_{\rm cr}$ that I derive may prove to be lower limits if the results of Pirronello are borne out by future work."
 We can combine equations 19. and 23.eBi (o write the total gas-phase IH» formation rate as while the grain-catalvzed rate can be written as, We can combine equations \ref{final_hm} and \ref{final_h2p} to write the total gas-phase $\mHt$ formation rate as while the grain-catalyzed rate can be written as
disks and growing spheroids.,disks and growing spheroids.
" Many would be moderately luminous, detectable byChandra deep surveys at z=2 if un-obscured, but the high expected obscuration makes their detectability dubious, especially for galaxies below "," Many would be moderately luminous, detectable by deep surveys at $z\!=\!2$ if un-obscured, but the high expected obscuration makes their detectability dubious, especially for galaxies below $10^{11}\,\msun$."
These BHs can grow inside classical bulges 10!!Mg.(Elmegreenetal.2008b)., These BHs can grow inside classical bulges \citep{EBE08b}.
". Observationally, the fraction of X-ray-detectable AGN indeed declines with decreasing stellar mass (Xueetal."," Observationally, the fraction of X-ray-detectable AGN indeed declines with decreasing stellar mass \citep{xue10}."
" Many AGN are highly absorbed, individually undetectable2010).. in X-rays (e.g.,Daddietal.2007),, and many of these obscured AGN have moderate intrinsic luminosities and lie in star-forming galaxies (Juneauetal. 2011)."," Many AGN are highly absorbed, individually undetectable in X-rays \citep[e.g.,][]{daddi07}, and many of these obscured AGN have moderate intrinsic luminosities and lie in star-forming galaxies \citep{J11}."
. Host galaxies of moderate-luminosity AGNs do not show an excess of major-merger signatures (Groginetal.2005;Gabor2009) and often have disky morphologies (Schawinskietal.," Host galaxies of moderate-luminosity AGNs do not show an excess of major-merger signatures \citep{grogin05,gabor09} and often have disky morphologies \citep{schawinski11}."
" Small mergers may also help feeding the BH while 2011)..leaving the disk intact (Shankar2010),, but they are a component of the cosmological streams, and can be considered to be part of the inflow within the disk."," Small mergers may also help feeding the BH while leaving the disk intact \citep{shankar10}, but they are a component of the cosmological streams, and can be considered to be part of the inflow within the disk."
 The violent instability phase should end when the cosmological accretion rate declines and the system becomes stellar dominated., The violent instability phase should end when the cosmological accretion rate declines and the system becomes stellar dominated.
" We expect less massive galaxies to remain unstable for longer times, because they retain higher gas fractions — one of the aspects of the downsizing of star formation (e.g.,Juneauet"," We expect less massive galaxies to remain unstable for longer times, because they retain higher gas fractions — one of the aspects of the downsizing of star formation \citep[e.g.,][]{J05}."
 This can result from the regulation of gas consumption in smaller galaxies (Dekel&Silk1986;Krumholz&Dekel and from the continuation of cold accretion to2011) lower redshift for lower-mass halos (Dekel&Birnboim2006).," This can result from the regulation of gas consumption in smaller galaxies \citep{DekelSilk86, KD11} and from the continuation of cold accretion to lower redshift for lower-mass halos \citep{DB06}."
". It induces an inverse gradient of gas fraction with galaxy mass (as observed, Kannappan and a downsizing in gravitational instability which2004)) could result in downsizing of BH growth."," It induces an inverse gradient of gas fraction with galaxy mass (as observed, \citealt{kannappan04}) ) and a downsizing in gravitational instability which could result in downsizing of BH growth."
" This is a longer growth phase into later redshifts in lower-mass galaxies, with »10!M galaxies growing BHs mostly at z>1 in their unstable disk phase, while clumpy disks of ~10!°M. may show AGN activity even after z~1."," This is a longer growth phase into later redshifts in lower-mass galaxies, with $\sim \! 10^{11}\msun$ galaxies growing BHs mostly at $z > 1$ in their unstable disk phase, while clumpy disks of $\sim \! 10^{10}\msun$ may show AGN activity even after $z \! \sim \! 1$."
 Simulations were performed at CCRT and TGCC under GENCI allocation 2011-GEN2192., Simulations were performed at CCRT and TGCC under GENCI allocation 2011-GEN2192.
" We acknowledge discussions with Frangooise Combes,| Bruce Elmegreen, Tiziana Di Matteo, James Mullaney, a constructive referee report, and support from grants ERC-StG-257720, CosmoComp ITN, ISF 6/08, GIF G-1052-104.7/2009, NSF AST-1010033 and a DIP grant."," We acknowledge discussions with Françooise Combes, Bruce Elmegreen, Tiziana Di Matteo, James Mullaney, a constructive referee report, and support from grants ERC-StG-257720, CosmoComp ITN, ISF 6/08, GIF G-1052-104.7/2009, NSF AST-1010033 and a DIP grant."
Field contamination is likely to be small in the anti-centre direction. but separation of field and cluster stars simply on the basis of photometry is neither easy nor unambiguous. since the cluster extends quite far (see Fig.,"Field contamination is likely to be small in the anti-centre direction, but separation of field and cluster stars simply on the basis of photometry is neither easy nor unambiguous, since the cluster extends quite far (see Fig."
 5 where panel a) corresponds to the cluster centre. ancl panels bh) to e) to increasing distances).," \ref{fig-rad} where panel a) corresponds to the cluster centre, and panels b) to e) to increasing distances)."
 Even at about 30 aremin from the centre there might be a slight cluster component mixed with the field stars., Even at about 30 arcmin from the centre there might be a slight cluster component mixed with the field stars.
 We will not discuss the cluster. dimension and racial distribution. since these will be amply treated by Testa Hamilton (1997).," We will not discuss the cluster dimension and radial distribution, since these will be amply treated by Testa Hamilton (1997)."
 Chiu van Altena (1981). published. a membership study of this cluster for the stars in AICP: their table 2 contains SOL stars down to V about. 18., Chiu van Altena (1981) published a membership study of this cluster for the stars in MCTF; their table 2 contains 801 stars down to V about 18.
 Of the S63 stars in our catalogue. 462 were cross-identified with MCTE on the basis of position and magnitude. while 401. mostly faint. are present only in our photometry.," Of the 863 stars in our catalogue, 462 were cross-identified with MCTF on the basis of position and magnitude, while 401, mostly faint, are present only in our photometry."
 Phe svnthetic CALDs deseribed in section 4 will be compared to our whole sample. since we lack membership information for V—120 and there is no way to assess the incompleteness of the membership study at the faint limit.," The synthetic CMDs described in section 4 will be compared to our whole sample, since we lack membership information for V=18--20 and there is no way to assess the incompleteness of the membership study at the faint limit."
 Phe CX cata are insteac very useful at the bright end. since we can be sure of how the MS behaves near the TO. of the position of the subgiant and red giant branches. and of the location of the chump (see Fig. 2((," The CvA data are instead very useful at the bright end, since we can be sure of how the MS behaves near the TO, of the position of the subgiant and red giant branches, and of the location of the clump (see Fig. \ref{fig-cmd4}( ("
0). where we show only those stars with membership probability z 0.70).,"c), where we show only those stars with membership probability $\ge$ 0.70)."
 Alost open clusters show indications of a sizeable binary population. detected. either speetroscopically (see e.g. Alermilliod Alavor 1989. 1990) or photometrically (see e.g. M 67. Fan et al.," Most open clusters show indications of a sizeable binary population, detected either spectroscopically (see e.g. Mermilliod Mayor 1989, 1990) or photometrically (see e.g. M 67, Fan et al."
 1996. or two clusters quite similar to NGC2506: NGC 2243. Bonilazi et al.," 1996, or two clusters quite similar to NGC2506: NGC 2243, Bonifazi et al."
 1990. and NGC 2420. Anthony-Pwaroe et al.," 1990, and NGC 2420, Anthony-Twarog et al."
 1990)., 1990).
 NGC2506 has been unsuccessfully surveveck for closemoderately close binaries in the past: Cameron Reid (1987) found no candidate binary among 26 subgiant/lower red giant cluster members., NGC2506 has been unsuccessfully surveyed for close/moderately close binaries in the past: Cameron Reid (1987) found no candidate binary among 26 subgiant/lower red giant cluster members.
 Their technique was based on the etection of chromospheric emission in the Call HHLIN lines due to enhanced axial rotation produced by orbital locking uxd was claimed to be sensible to periods of about 1 to 50 davs., Their technique was based on the detection of chromospheric emission in the CaII H+K lines due to enhanced axial rotation produced by orbital locking and was claimed to be sensible to periods of about 1 to 50 days.
 Waluzny Shara (1988) found no contact binaries of rw W UMa type. with periods in the range a few hours 1.5 days. in their CCD survey of 6 old open clusters. among which was NGC2506.," Kaluzny Shara (1988) found no contact binaries of the W UMa type, with periods in the range a few hours – 1.5 days, in their CCD survey of 6 old open clusters, among which was NGC2506."
 MC'TE. on the basis of the scatter of re MS. estimated a crude 50% of binaries in NOGC2506.," MCTF, on the basis of the scatter of the MS, estimated a crude 50 of binaries in NGC2506."
 The VibV diagram is not the most suitable to separate the secondary sequence due to binary systems rom the single-star MS: much better resolution in colour is obtained in the CMD's involving the U band (see Fig. 3)).," The $V,B-V$ diagram is not the most suitable to separate the secondary sequence due to binary systems from the single-star MS: much better resolution in colour is obtained in the CMD's involving the U band (see Fig. \ref{fig-cmd5}) )."
 ‘To quantify the visual impression of well populated. binary star sequences. we have performed a simple experiment on 1CΑΕΕ. our best field. both for photometric accuracy ancl crowcding conditions. using the C.C1 CAID.," To quantify the visual impression of well populated binary star sequences, we have performed a simple experiment on RCA-F1, our best field both for photometric accuracy and crowding conditions, using the $U,U-I$ CMD."
 After delining a MS ridge linc. we have measured the distance in colour of every star from this MS and plotted the distance histograms in [our separa magnitude bins (see Fig.," After defining a MS ridge line, we have measured the distance in colour of every star from this MS and plotted the distance histograms in four separate magnitude bins (see Fig."
 ο left panel for a definition ofthe MS and magnitude intervals. and Fig.," \ref{fig-bin} – left panel – for a definition of the MS and magnitude intervals, and Fig."
 6 micelle panel — for the histograms)., \ref{fig-bin} – middle panel – for the histograms).
 A secondary. peak in the colour distribution is clearly present: counting stars in it and in the “AIS peak. we derive a binary Frequency of about ..τν.," A secondary peak in the colour distribution is clearly present: counting stars in it and in the “MS peak”, we derive a binary frequency of about 17."
 Finally. the right panel of Fie.," Finally, the right panel of Fig."
 6 shows the histogram of the distance of cach star in RCA-F1 from the AIS ridge line: here too the secondary. peak is clearly prominent., \ref{fig-bin} shows the histogram of the distance of each star in RCA-F1 from the MS ridge line: here too the secondary peak is clearly prominent.
 This value is referred to the central part of the cluster.," This value is referred to the central part of the cluster,"
 Inserting (14) together with (18). and making the further definitions v=(y—yi/c and m=(y;—b)fc. we transform this integral as follows. The approximation in the first equality refers to replacing the lower boundary of the integral —y;/c by —eo. which is valid if dosevilco»lobe. if GQ)) is compact (c«I) and peaks not too to zero (y;>0).," Inserting ) together with ), and making the further definitions $y \equiv (\chi-\chi_i)/\sigma$ and $m \equiv (\chi_i- b)/\sigma$, we transform this integral as follows, The approximation in the first equality refers to replacing the lower boundary of the integral $-\chi_i/\sigma$ by $-\infty$, which is valid if $\chi_i/\sigma \gg 1$ , i.e. if $G^{(i)}(\chi)$ is compact $\sigma \ll 1$ ) and peaks not too close to zero $\chi_i \gg 0$ )."
 The root of the term in curly brackets can be found numerically. resulting i We have solved (19) directly and plot the resulting 6 1.," The root of the term in curly brackets can be found numerically, resulting in We have solved ) directly and plot the resulting $b$ in $\,$."
 We find excellent agreement with the approximate solution (21) as long as the assumption discussed above ts fulfilled., We find excellent agreement with the approximate solution ) as long as the assumption discussed above is fulfilled.
 Significant deviations from the linear behaviour of b as a function of c are only found for y;jocx2., Significant deviations from the linear behaviour of $b$ as a function of $\sigma$ are only found for $\chi_i/\sigma \lesssim 2$.
 For reasons of simplicity we will restrict ourselves to cases where the approximation (21) holds., For reasons of simplicity we will restrict ourselves to cases where the approximation ) holds.
 This means in particular that we will not consider signals at very small redshifts. where y; Is necessarily small.," This means in particular that we will not consider signals at very small redshifts, where $\chi_i$ is necessarily small."
 In. practice we use the condition GP(Wai)xO0 with μη the minimum redshift used in the survey as a simple cross-check to ensure that this approximation is sufficiently accurate., In practice we use the condition $G^{(i)}(\chi(z_{\rm min})) \approx 0$ with $z_{\rm min}$ the minimum redshift used in the survey as a simple cross-check to ensure that this approximation is sufficiently accurate.
 The conditions specified above are strictly fulfilled only for continuous y or z., The conditions specified above are strictly fulfilled only for continuous $\chi$ or $z$.
 However. we will in practice use the discretised transformation (10) and thus have to make sure that GI boosting and GG suppression work accurately also in this case.," However, we will in practice use the discretised transformation ) and thus have to make sure that GI boosting and GG suppression work accurately also in this case."
 Via a procedure outlined in the following. we optimise the remaining free parameter c to guarantee a good sampling of αγ) by the discrete set of weights B(y(z;) with j=1....N.. thereby fulfilling dG= Oand (19) to good accuracy.," Via a procedure outlined in the following, we optimise the remaining free parameter $\sigma$ to guarantee a good sampling of $G^{(i)}(\chi)$ by the discrete set of weights $B^{(i)}(\chi(z_j))$ with $j=1,\,..\,,N_z$, thereby fulfilling $\partial G^{(i)}/ \partial \chi\, |_{\chi_i} = 0$ and ) to good accuracy."
 As the sampling points of (10) we choose the medians of the redshift distributions of the galaxy samples employed., As the sampling points of ) we choose the medians of the redshift distributions of the galaxy samples employed.
 It is expected that the optimal choice of the parameter ος. denoted by cry in the following. will depend intricately on the positions of these sampling points and hence on the redshift distributions of the different galaxy samples in the cosmic shear data. in particular if the number of sampling points 15 small. e.g. if the distributions have a large scatter.," It is expected that the optimal choice of the parameter $\sigma$, denoted by $\sigma_{\rm opt}$ in the following, will depend intricately on the positions of these sampling points and hence on the redshift distributions of the different galaxy samples in the cosmic shear data, in particular if the number of sampling points is small, e.g. if the distributions have a large scatter."
 Since the binning is done in terms of redshift. it is convenient to work with the quantity σιΞdpac instead of c.," Since the binning is done in terms of redshift, it is convenient to work with the quantity $\sigma_z \equiv \bc{ \chi'(z_i) }^{-1} \sigma$ instead of $\sigma$."
 We will also give our choices of cry in terms of co- throughout., We will also give our choices of $\sigma_{\rm opt}$ in terms of $\sigma_z$ throughout.
 We introduce the discrete version of the function Gy) Then we consider the root mean square deviation of all function values Gt(yG4) used. as a criterion for how well G'(y) is MN by the discrete set of function values B(y(z;)) dependenceentering (22).," We introduce the discrete version of the function $G^{(i)}(\chi)$, Then we consider the root mean square deviation of all function values ${G'}^{(i)} (\chi(z_k))$ used, as a criterion for how well $G^{(i)}(\chi)$ is sampled by the discrete set of function values $B^{(i)}(\chi(z_j))$ entering )."
 In the equation above we have made the on c- explicit in the arguments., In the equation above we have made the dependence on $\sigma_z$ explicit in the arguments.
" We emphasise that the determination of σ- via the diagnostic Z is optimalonly in the sense that it allows us to find a representative sampling of G''(y) such that dG""/dy|,=0 and (19) hold to good accuracy."," We emphasise that the determination of $\sigma_z$ via the diagnostic $\zeta$ is optimalonly in the sense that it allows us to find a representative sampling of $G^{(i)}(\chi)$ such that $\partial G^{(i)}/ \partial \chi\, |_{\chi_i} = 0$ and ) hold to good accuracy."
 It willin general not yield an, It willin general not yield an
velocity dispersion). is the Fundamental Plane thereafter FP: Djorgovski Davis 1987: Faber et al.,"velocity dispersion), is the Fundamental Plane (hereafter FP; Djorgovski Davis 1987; Faber et al."
 1987. Dressler et al.," 1987, Dressler et al."
 1987: Bernardi et al., 1987; Bernardi et al.
 2003)., 2003).
" In theory. the FP is derived from the virial theorem as Hoxστ1(MÉL)+.where 4. is the effective surface brightness in flux units. calculated within the half-light radius. /7,.. of the galaxy. σ is the galaxy internal velocity dispersion and AZ/L is its mass-to-light ratio."," In theory, the FP is derived from the virial theorem as $R_e\propto \sigma^2 I_e^{-1} (M/L)^{-1}$,where $I_e$ is the effective surface brightness in flux units, calculated within the half-light radius, $R_e$, of the galaxy, $\sigma$ is the galaxy internal velocity dispersion and $M/L$ is its mass-to-light ratio."
 Assuming that the AZ/2£ is expressed by a power-law function of σ and/or the effective surface brightness. /... the FP relation is simplitied as where i5. is the mean surface brightness in megaresec? unit and is defined as fo.Ξ2.5log(.) cle.," Assuming that the $M/L$ is expressed by a power-law function of $\sigma$ and/or the effective surface brightness, $I_e$ , the FP relation is simplified as where $\langle \mu \rangle_e$ is the mean surface brightness in $mag/arcsec^2$ unit and is defined as $\langle \mu \rangle_e=-2.5log(I_e)+cte$ ."
 Although the shape of the FP and its coefficients differs for different gravitationally bound systems from globular clusters (Burstein et al., Although the shape of the FP and its coefficients differs for different gravitationally bound systems from globular clusters (Burstein et al.
 1997) to galaxy clusters (Schaeffer et al., 1997) to galaxy clusters (Schaeffer et al.
 1993: Fritsch Buchert 1999: Zaritsky et al., 1993; Fritsch Buchert 1999; Zaritsky et al.
 2006a: ZGZ06). there is no doubt about its existence (Lucey. Bower Ellis 1991 and its references).," 2006a: ZGZ06), there is no doubt about its existence (Lucey, Bower Ellis 1991 and its references)."
 In reality. the coefticients of the FP relation differ from the prediction of the virial theorem.," In reality, the coefficients of the FP relation differ from the prediction of the virial theorem."
 The observed coefficients are AZ1.24 and B=0.33 (Jorgensen. Franks and Kpergaard 1996: JFK96) while the virial theorem predicts A-2.0 and B=0.4.," The observed coefficients are A=1.24 and B=0.33 rgensen, Franks and rgaard 1996: JFK96) while the virial theorem predicts A=2.0 and B=0.4."
" This difference. often referred to as the ""tilt"" of the FP. is mainly attributed to different formation histories and evolutionary processes."," This difference, often referred to as the “tilt"" of the FP, is mainly attributed to different formation histories and evolutionary processes."
 The difference in FP coefficients of different spheroidal systems with different mass. size and luminosities. can be explained by evolution of the A7/L ratio as a function of stellar population [age. metalicity or initial mass function (IMF)] and/or dark matter. content (Tortora et al.," The difference in FP coefficients of different spheroidal systems with different mass, size and luminosities, can be explained by evolution of the $M/L$ ratio as a function of stellar population [age, metalicity or initial mass function (IMF)] and/or dark matter content (Tortora et al."
 2009: Grillo Gobat 2010: Graves Faber 2010)., 2009; Grillo Gobat 2010; Graves Faber 2010).
 In addition. the absence of homology. ie. the fact that. the structure of spheroids is scalable regardless of their size. can be the source of the FP tilt (D'Onofrio et al.," In addition, the absence of homology, i.e, the fact that, the structure of spheroids is scalable regardless of their size, can be the source of the FP tilt (D'Onofrio et al."
 2008: Trujillo. Burkert Bell 2004).," 2008; Trujillo, Burkert Bell 2004)."
 On the other hand. some authors studied the role of dissipation in explaining the nature of the FP (Ribeiro Dantas 2010: Hopkins. Thoms Hernquist 2008: HCHO8).," On the other hand, some authors studied the role of dissipation in explaining the nature of the FP (Ribeiro Dantas 2010; Hopkins, Thoms Hernquist 2008: HCH08)."
 HCHOS claimed that the non-homology or change in the dark matter distribution are not the main drivers of FP tilt., HCH08 claimed that the non-homology or change in the dark matter distribution are not the main drivers of FP tilt.
 Studying the early-type galaxies in 59 nearby galaxy clusters. D'Onofrio et al. (," Studying the early-type galaxies in 59 nearby galaxy clusters, D'Onofrio et al. ("
2008) have found a strong correlation between the FP coefficients and the ocal cluster environment and no strong correlations with internal galaxy properties (e.g. Sversic index and galaxy colour).,2008) have found a strong correlation between the FP coefficients and the local cluster environment and no strong correlations with internal galaxy properties (e.g. S'ersic index and galaxy colour).
 Moreover. FP coefficients are independent of global oroperties of clusters such as radius. X-ray luminosity and cluster velocity dispersion (D'Onofrio et al..," Moreover, FP coefficients are independent of global properties of clusters such as radius, X-ray luminosity and cluster velocity dispersion (D'Onofrio et al.,"
 2008)., 2008).
 On the other hand. while Reda. Forbes Hau (2005) have shown hat isolated early-type galaxies lie on the same fundamental jane as galaxies in high-density environments. cluster galaxies have also less intrinsic scatter in their properties compared to field galaxies (de Carvalho Dyjorgovski. 993).," On the other hand, while Reda, Forbes Hau (2005) have shown that isolated early-type galaxies lie on the same fundamental plane as galaxies in high-density environments, cluster galaxies have also less intrinsic scatter in their properties compared to field galaxies (de Carvalho Djorgovski, 1992)."
 The study of dark matter in dwarf galaxies showed —iat below the critical virial velocity. which is estimated to be ~ 100 km +. interstellar gas removal via supernova explosions become important (Dekel Silk 1986: DS86).," The study of dark matter in dwarf galaxies showed that below the critical virial velocity, which is estimated to be $\sim$ 100 km $^{-1}$, interstellar gas removal via supernova explosions become important (Dekel Silk 1986: DS86)."
 This mechanism has been invoked also to explain the shape of the low mass dwarf galaxies and their mass profiles (Sánnchez-Janssen et al., This mechanism has been invoked also to explain the shape of the low mass dwarf galaxies and their mass profiles (Sánnchez-Janssen et al.
 2010: Governato et al., 2010; Governato et al.
 2010)., 2010).
 Differences between the FP of giant and dwarf galaxies have been Known for some time (Nieto et al., Differences between the FP of giant and dwarf galaxies have been known for some time (Nieto et al.
 1990: Bender et al., 1990; Bender et al.
 1992: Guzman et al., 1992; Guzman et al.
 1993)., 1993).
 Peterson Caldwell (1993). studying a sample of nucleated dwarfs. they found a change in A//£ ratio with luminosity as predicted by the scaling relations of DS86.," Peterson Caldwell (1993), studying a sample of nucleated dwarfs, they found a change in $M/L$ ratio with luminosity as predicted by the scaling relations of DS86."
 The study of 17 Virgo dwarfs (-17.5<My x-15.48 by Geha et al. (, The study of 17 Virgo dwarfs $<M_V<$ -15.48) by Geha et al. (
2002 2003) placed dEs on a plane parallel to. but offset. from that occupied by normal elliptical galaxies.,"2002 2003) placed dEs on a plane parallel to, but offset, from that occupied by normal elliptical galaxies."
 On the other hand. Graham Guzmánn (2003). based on a detailed surface photometry from HST archival images for a sample of «Ες in the Coma cluster. claimed that the E-dE dichotomy is a consequence of a continuously varying profile shape with galaxy luminosity.," On the other hand, Graham Guzmánn (2003), based on a detailed surface photometry from HST archival images for a sample of dEs in the Coma cluster, claimed that the E-dE dichotomy is a consequence of a continuously varying profile shape with galaxy luminosity."
 They attributed the E-dE dichotomy to the fact that the light profile of dwarf galaxies do not follow the de Vaucouleurs law (1948) and therefore. concluded that giant and dwarf ellipticals thereafter dEs) were not formed by different mechanisms.," They attributed the E-dE dichotomy to the fact that the light profile of dwarf galaxies do not follow the de Vaucouleurs law (1948) and therefore, concluded that giant and dwarf ellipticals (hereafter dEs) were not formed by different mechanisms."
 However. as found by Geha et al. (," However, as found by Geha et al. ("
2003). modelling the dEs with Sérrsic functions (1968) does not explain why dEs lie on a different region of the FP.,"2003), modelling the dEs with Sérrsic functions (1968) does not explain why dEs lie on a different region of the FP."
 Although it is elear that dwarf ellipticals do not lie on the fundamental plane as defined by brighter ellipticals (Geha et al., Although it is clear that dwarf ellipticals do not lie on the fundamental plane as defined by brighter ellipticals (Geha et al.
 2003). Zaritskv et al. (," 2003), Zaritsky et al. ("
"2006 a.b) showed that the differences between the ""fundamental planes"". defined by different families of spheroids from dwarf ellipticals to galaxy clusters. is the result of a non-linear relationship between loga and log Al/L.","2006 a,b) showed that the differences between the “fundamental planes”, defined by different families of spheroids from dwarf ellipticals to galaxy clusters, is the result of a non-linear relationship between $\log{\sigma}$ and $\log{M/L}$ ."
 Desroches et al. (, Desroches et al. (
2007) also found thesensitivity of the coefficient A of the FP to the,2007) also found thesensitivity of the coefficient A of the FP to the
las adnaenitude of 10.77 imag in the bandpass and was selected by the Asteroscisunic Science Cousortimm (ISASC.Cillibuidetal.2009) for one-cadence observations.,"has a magnitude of 10.77 mag in the bandpass and was selected by the Asteroseismic Science Consortium \citep[KASC,][]{gilliland2009} for long-cadence observations."
" The time series data lave a 29,[-1in cadence aud show both solar-like oscillations in he frequency regine expected for a red eiut. and an eclipse (κος Fig. 1))."," The time series data have a 29.4-min cadence and show both solar-like oscillations in the frequency regime expected for a red giant, and an eclipse (see Fig. \ref{obs}) )."
 This star has also been part of the TrES-Lyrl exouud-jased. (O'Donovanetal.2006)— from which time series xXotometrv with a time span of dd are at our disposal., This star has also been part of the TrES-Lyr1 ground-based \citep{odonovan2006} from which time series photometry with a time span of d are at our disposal.
 Iu these data no eclipse is visible. while the primary and possibly the secondary eclipse (depth predicted. iu Section 6) would have heen detected. had they occurred caving the time span of the observations.," In these data no eclipse is visible, while the primary and possibly the secondary eclipse (depth predicted in Section 6) would have been detected, had they occurred during the time span of the observations."
 The fact that the primary is not ποσα in the TYES data implies that the period is longer than the duration of the TrES observations., The fact that the primary is not seen in the TrES data implies that the period is longer than the duration of the TrES observations.
 This star was also observed photometrically by the SuperWASP cameras (Pollaccoetal.2006) for 100 davs iu 2007 and for 90 days in 2008., This star was also observed photometrically by the SuperWASP cameras \citep{pollacco2006} for 100 days in 2007 and for 90 days in 2008.
 There is a suspicion of a secondary eclipse. but no primary eclipse has beeu observed during these periods. which have some gaps due to weather conditious.," There is a suspicion of a secondary eclipse, but no primary eclipse has been observed during these periods, which have some gaps due to weather conditions."
 Observations for NICsLLO637 are also available from the All Sky Automated Survey CASAS) (Pojmaiski1997).. iu which there are three points that have lower flux thin the imajoritv.," Observations for KIC8410637 are also available from the All Sky Automated Survey (ASAS) \citep{pojmanski1997}, in which there are three points that have lower flux than the majority."
 These points night be due to the eclipse. which would imply an orbital period of about 380dd or a subimultiple thereof.," These points might be due to the eclipse, which would imply an orbital period of about d or a submultiple thereof."
 Unfortunately. there are too few poiuts. aud their nucertaimty is too laree to draw fiii conclusions.," Unfortunately, there are too few points, and their uncertainty is too large to draw firm conclusions."
 Additionally. we checked whether we could put mere constraints on the period from the time stamps of the data we have at our disposal.," Additionally, we checked whether we could put more constraints on the period from the time stamps of the data we have at our disposal."
 We determüned which periods would have led us to miss primary eclipses because of a lack of observations., We determined which periods would have led us to miss primary eclipses because of a lack of observations.
 We would have iuissed all eclipses for an orbital period of about 135. 220. 253. 261. 270 and 295 davs.," We would have missed all eclipses for an orbital period of about 135, 220, 253, 261, 270 and 295 days."
 All other possible periods are longer than the 380 davs. iieutioned above.," All other possible periods are longer than the 380 days, mentioned above."
 For KICS110637 the (Brownetal.2005) provides a set of stellar paraiicters. as listed in Table 1..," For KIC8410637 the \citep{brown2005} provides a set of stellar parameters, as listed in Table \ref{param}."
 Tweho-2 (Πουetal.2000) and 2MASS (Skrutskieetal.2006) photometry are also available., Tycho-2 \citep{hog2000} and 2MASS \citep{skrutskie2006} photometry are also available.
 Ofek(2008). fitted these colors with svuthetic photometry of stellar templates (Pickles1998)., \citet{ofek2008} fitted these colors with synthetic photometry of stellar templates \citep{pickles1998}.
. Using two methods. Ofek(2008) classified the star as a imietal-weak I21II star with au effective temperature of hy or a IKIV star with a temperature of IKK. We also obtained three spectra with the IIERMES spectrograph (Raskin&VauWinckel2008) mounted ou he Mercator Telescope. La Palma. Spain: two ou 2009 October 13 and one on 2009 October 15.," Using two methods, \citet{ofek2008} classified the star as a metal-weak K2III star with an effective temperature of K or a K4V star with a temperature of K. We also obtained three spectra with the HERMES spectrograph \citep{raskin2008}
 mounted on the Mercator Telescope, La Palma, Spain: two on 2009 October 13 and one on 2009 October 15."
 These spectra rave a resolution of about 850000 and a waveleneth range of 10009000AA., These spectra have a resolution of about 000 and a wavelength range of $4000 - 9000$.
. We analyzed the average spectrum. which has a vpical signal-to-noise ratio of ~ 90 at 6000A.. with nethods described by Tekker&Melendez(2007). aud list he results iu the last column of Table 1..," We analyzed the average spectrum, which has a typical signal-to-noise ratio of $\sim$ 90 at 6000, with methods described by \citet{hekker2007} and list the results in the last column of Table \ref{param}."
" This method does uot directly provide uncertaiuties on the individual xuwanmeters,", This method does not directly provide uncertainties on the individual parameters.
 Tere. we adopted the total of the differcuce and scatter between the values of Hekker&Aleléudez(2007) and the values of Luck&Ueiter(2007) (κουIekker&Moléudez2007.formore details)...," Here, we adopted the total of the difference and scatter between the values of \citet{hekker2007} and the values of \citet{luck2007}
 \citep[see][for more details]{hekker2007}."
 For esin/. we computed the uncertainty from a uncertaiutv in the total FWIIM of the lines. which is combined with the & on the macro turbulence.," For $v\sin i$, we computed the uncertainty from a uncertainty in the total FWHM of the lines, which is combined with the $\sigma$ on the macro turbulence."
 These results are in agreement with those in the KIC aud the literature., These results are in agreement with those in the KIC and the literature.
 They confirm the red-giaut nature of the primary star., They confirm the red-giant nature of the primary star.
 We investigated the spectrun for signatures of the secondary star. in the fori of additional absorption lines and cross-correlation profiles. but no clear signatures could be detected.," We investigated the spectrum for signatures of the secondary star, in the form of additional absorption lines and cross-correlation profiles, but no clear signatures could be detected."
 The three spectra obtained over a time span of 3 davs are not sufficieut to see radial velocity variatious due to the binary orbit., The three spectra obtained over a time span of 3 days are not sufficient to see radial velocity variations due to the binary orbit.
 Solar-like oscillations are visible both durug aud outside the eclipse (Fie. 1))., Solar-like oscillations are visible both during and outside the eclipse (Fig. \ref{obs}) ).
 To analyze these we adopted two approaches., To analyze these we adopted two approaches.
 Iu the first approach. we corrected the light curve inchiding the eclipse bv fitting aud subtracting trends for different parts of the time series using a linear polvuonual for the commissioning data. a second-order polynomial for the part of the QI data before the eclipse. the mean value for the data in full," In the first approach, we corrected the light curve including the eclipse by fitting and subtracting trends for different parts of the time series using a linear polynomial for the commissioning data, a second-order polynomial for the part of the Q1 data before the eclipse, the mean value for the data in full"
"was serendipitously discovered during the first INTErnational Gamma-Ray Astrophysical LaboratoryUNTEGRAL, Winkler et al.","was serendipitously discovered during the first INTErnational Gamma-Ray Astrophysical Laboratory, Winkler et al."
 2003) observation of the Galactic microquasar GRS 1915-105. (Hannikainen. Rodriguez Pottschmidt 2003).," 2003) observation of the Galactic microquasar GRS 1915+105 (Hannikainen, Rodriguez Pottschmidt 2003)."
 Inspection. of the high energy archives showed it to be the most likely hard X-ray counterpart of the poorly studiedEXOSAT source EXO 1912+097 (Lu et al., Inspection of the high energy archives showed it to be the most likely hard X-ray counterpart of the poorly studied source EXO 1912+097 (Lu et al.
 1996)., 1996).
 Soon after its discovery a Target of Opportunity (ToO) was performed on with the(RXTE)., Soon after its discovery a Target of Opportunity (ToO) was performed on with the.
 The preliminary spectral analysis of this ToO showed the source had a rather hard spectrum. fitted with a power law of photon index 1.6. and an absorptior column density Nu26x107 οι (Swank Markwardt 2003).," The preliminary spectral analysis of this ToO showed the source had a rather hard spectrum, fitted with a power law of photon index 1.6, and an absorption column density $_{\mathrm{H}}$ $\times10^{22}$ $^{-2}$ (Swank Markwardt 2003)."
 Recently timing analysis of the RXTE/ASM data revealed ai X-ray period of 13.55 days (Corbet. Hannikainen Remillare 2004).," Recently timing analysis of the /ASM data revealed an X-ray period of 13.55 days (Corbet, Hannikainen Remillard 2004)."
 This analysis showed that the source was detected even during the early days of theRXTE mission. which suggests that is a persistent X-ray source although most of the time in a faint state.," This analysis showed that the source was detected even during the early days of the mission, which suggests that is a persistent X-ray source although most of the time in a faint state."
 In a companior paper (Hannikaimen et al., In a companion paper (Hannikainen et al.
 2004a. hereafter Paper 1) we usec the latest version of the software to refine anc give the most accurate X-ray position of (see also Cabanae et al.," 2004a, hereafter Paper 1) we used the latest version of the software to refine and give the most accurate X-ray position of (see also Cabanac et al."
 2004). which allowed us to obtai the most accurate X-ray/Gamma-ray spectra of the source.," 2004), which allowed us to obtain the most accurate X-ray/Gamma-ray spectra of the source."
 High energy spectral analysis of covering the period of its discovery.ο.," High energy spectral analysis of covering the period of its discovery,."
 during revolutioi 48. was presented for the first time.," during revolution 48, was presented for the first time."
 We have. in particular. shown. that during this observation. the source. although very variable. showed two distinct spectral behaviours.," We have, in particular, shown, that during this observation, the source, although very variable, showed two distinct spectral behaviours."
 The first one manifests a thermal component (black body-like) in the soft X-rays. and a hard X-ray tail. whereas the second one is harder and can be interpreted as originating from thermal Comptonisation (Paper Although it ts very likely that is a Galactic object. the nature of the compact object is unclear.," The first one manifests a thermal component (black body-like) in the soft X-rays, and a hard X-ray tail, whereas the second one is harder and can be interpreted as originating from thermal Comptonisation (Paper Although it is very likely that is a Galactic object, the nature of the compact object is unclear."
 The spectral analysis presented in Paper | would tend to favour a neutron star. but no definite conclusion could be drawn from the data We report here observations of with performed between March and May 2003. during a very well sampled and unprecedented high energy coverage of this source.," The spectral analysis presented in Paper 1 would tend to favour a neutron star, but no definite conclusion could be drawn from the data We report here observations of with performed between March and May 2003, during a very well sampled and unprecedented high energy coverage of this source."
" To our monitoring. we add the analysis of the March 2003RXTE ToO. às well as the analysis of observations performed one year earlier on EXO 19124097, and the first observation of a monitoring campaign we are currently leading on IGR J19140+0951. performed in April 2004."," To our monitoring, we add the analysis of the March 2003 ToO, as well as the analysis of observations performed one year earlier on EXO 1912+097, and the first observation of a monitoring campaign we are currently leading on IGR J19140+0951, performed in April 2004."
 We describe the sequence of observations and the data reduction procedures in Sec., We describe the sequence of observations and the data reduction procedures in Sec.
 2. and then present the results obtained from the different instruments in Sec.," 2, and then present the results obtained from the different instruments in Sec."
 3., 3.
 The results are discussed in the last part of the paper., The results are discussed in the last part of the paper.
 The 1.2212 keV RXTE/ASM. 20-40 keV and 40-80 keV TEGRAL/ISGRI light curves of the source covering the period of interest are shown in Fig. 1..," The 1.2–12 keV ASM, 20–40 keV and 40–80 keV /ISGRI light curves of the source covering the period of interest are shown in Fig. \ref{fig:lc}."
 The main instruments on board are IBIS (Ubertini et al., The main instruments on board are IBIS (Ubertini et al.
 2003). and SPI (Vedrenne et al.," 2003), and SPI (Vedrenne et al."
 2003)., 2003).
 Both instruments use coded masks which allow y-ray imaging over large Fields of View (FOV). ~30x30° up to zero response.," Both instruments use coded masks which allow $\gamma$ -ray imaging over large Fields of View (FOV), $\sim30 \times30^\circ$ up to zero response."
 The Totally Coded FOV (TCFOV) is à smaller part within which the detector has the highest response., The Totally Coded FOV (TCFOV) is a smaller part within which the detector has the highest response.
 IBIS has a TCFOV of 9°x9°. while that of SPI is 16° (corner to corner).," IBIS has a TCFOV of $9^\circ\times9^\circ$, while that of SPI is $^\circ$ (corner to corner)."
 The Soft Gamma-Ray Imager (ISGRI. Lebrun et al.," The Soft Gamma-Ray Imager (ISGRI, Lebrun et al."
 2003) is the top layer of the IBIS detection plane. and covers the energy range from 13 keV to a few hundred The JEM-X monitors (Lund et al.," 2003) is the top layer of the IBIS detection plane, and covers the energy range from 13 keV to a few hundred The JEM-X monitors (Lund et al."
 2003). consist of two identical coded mask instruments designed for X-ray imaging in the range 3-35 keV with an angular resolution of 3 aremin," 2003), consist of two identical coded mask instruments designed for X-ray imaging in the range 3–35 keV with an angular resolution of 3 arcmin"
"The surface molecular gegas uss density Xj, aud the inoleculu gas niass Mg are thus calculated as. where wy, aud dQ are the mass of a hydrogen molecule and the solid augle of the integrated area. respectively.","The surface molecular gas mass density $\Sigma_{\rm H_2}$ and the molecular gas mass $M_{\rm H_2}$ are thus calculated as, where $m_{\rm H_{2}}$ and $\Omega$ are the mass of a hydrogen molecule and the solid angle of the integrated area, respectively."
 Since this conversion factor is used for the !1?CO(J = 1 enission. we assume the simplest case where the ratio of COL = 1)/172CO(J = 10) is unity.," Since this conversion factor is used for the $^{12}$ CO(J = 1--0) emission, we assume the simplest case where the ratio of $^{12}$ CO(J = $^{12}$ CO(J = 1–0) is unity."
 The Ty mass of the clumps measured within oue svuthesized beam are listed in Table 2.., The $_{2}$ mass of the clumps measured within one synthesized beam are listed in Table \ref{t.clumps}.
 We would eet the molecular gas ass CM) by unultiplvine the IT) mass by the mean atomic weight of 1.36 of the Te correction., We would get the molecular gas mass $M_{\rm gas}$ ) by multiplying the $_{2}$ mass by the mean atomic weight of 1.36 of the He correction.
 The ο = 2 1)is usually optically thick aud ouly traces the surface properties of the clouds., The $^{12}$ CO(J = 2–1) is usually optically thick and only traces the surface properties of the clouds.
 Optically thinner PCO = 21) would be a better estimator of the total column deusitv., Optically thinner $^{13}$ CO(J = 2–1) would be a better estimator of the total column density.
" We therefore caleulate the average Ag, ia the uucleus aud the ring to compare the Afg, derived. frou both 11Ο = 21) and PCOU = 21) lines.", We therefore calculate the average $M_{\rm H_{2}}$ in the nucleus and the ring to compare the $M_{\rm H_{2}}$ derived from both $^{12}$ CO(J = 2–1) and $^{13}$ CO(J = 2–1) lines.
" In the case of “COL = 21). the average Af, of the ring is about 1.7410! AL. within one svuthesized beam. aud the corresponding value for the uncleus is 3.1510"" AL..."," In the case of $^{12}$ CO(J = 2–1), the average $M_{\rm H_{2}}$ of the ring is about $\times10^{7}$ $_{\odot}$ within one synthesized beam, and the corresponding value for the nucleus is $\times10^{7}$ $_{\odot}$."
 However. iu Paper E we derived the intensity ratio of ο 22 L/P COW 2 10) is ~2 for the nucleus in the lower resolution map.," However, in Paper I we derived the intensity ratio of $^{12}$ CO(J = $^{12}$ CO(J = 1–0) is $\sim$ 2 for the nucleus in the lower resolution map."
" Therefore the of the nucleus is possibly smaller than that derived Afi,above.", Therefore the $M_{\rm H_{2}}$ of the nucleus is possibly smaller than that derived above.
 The conveutional Galactic New of 31029 cm2 (Klaus +) bis often suggested to be overestimated in the galactic center and the starburst cuviromment by a factor of 2 to 5 (e.g...Maloney&Black1988:MeierTuner 2001).," The conventional Galactic $X_{\rm CO}$ of $\times10^{20}$ $^{-2}$ (K km $^{-1}$ $^{-1}$ is often suggested to be overestimated in the galactic center and the starburst environment by a factor of 2 to 5 \citep[e.g.,][]{maloney88,meier01}."
". Therefore. our estimated Mj, iu the ring and the nucleus might be sualler at least by a factor of 2,"," Therefore, our estimated $M_{\rm H_{2}}$ in the ring and the nucleus might be smaller at least by a factor of 2."
 Asstuning the ο = 21) emission js optically thin. aud the @CO/IMs abuudauce of «10.© (Solomonetal L979).," Assuming the $^{13}$ CO(J = 2–1) emission is optically thin, and the $^{13}$ $_{2}$ abundance of $\times10^{-6}$ \citep{sol79}."
". We calculate the Ady, with excitation teiiperature (Zi) of 20 Ik aud 50 K in LTE condition.", We calculate the $M_{\rm H_{2}}$ with excitation temperature $T_{\rm ex}$ ) of 20 K and 50 K in LTE condition.
" The Mj, of the nucleus averaged over one svutlicsized bean are (0.042.5) NE. and (L340.1) &10*. MT, for 20 Is aud 50 EK. respectively."," The $M_{\rm H_{2}}$ of the nucleus averaged over one synthesized beam are $\pm$ $\times10^{6}$ $_{\odot}$ and $\pm$ $\times10^{7}$ $_{\odot}$ for 20 K and 50 K, respectively."
"4109. The My, of the ring averaged over oue svuthesized beam are 100 ME. and 1.100 4107 ML. for 20 K and 50 K. respectively,"," The $M_{\rm H_{2}}$ of the ring averaged over one synthesized beam are $\pm$ $\times10^{6}$ $_{\odot}$ and $\pm$ $\times10^{7}$ $_{\odot}$ for 20 K and 50 K, respectively."
 With the assumption of coustant CO/TL abundance. if the T;4 is < 20 EK. then the conversion ‘actor we adopted for the °CO line is overestimated by a actor of ~3 to 5.," With the assumption of constant $^{13}$ $_{2}$ abundance, if the $T_{\rm ex}$ is $\le$ 20 K, then the conversion factor we adopted for the $^{12}$ CO line is overestimated by a factor of $\sim$ 3 to 5."
 The overestimation is smaller (a factor of 2 to 3) if the eas is as warm as 50 I. We measured ho οσο = 21) fus lugher than 23e to derive the otal II» mass of the nucleus aud the ring., The overestimation is smaller (a factor of 2 to 3) if the gas is as warm as 50 K. We measured the $^{12}$ CO(J = 2–1) flux higher than $\sigma$ to derive the total $_{2}$ mass of the nucleus and the ring.
 The total fiux the nueleus aud the ring are 190.2+59.3 Jv lan s| id 2902.2+ 176.8 Jy hans 1. respectively.," The total flux of the nucleus and the ring are $\pm$ 59.3 Jy km $^{-1}$ and $\pm$ 476.8 Jy km $^{-1}$, respectively."
" If we adopted he intensity ratios of ΟΙ = 21) aud ? COW = LO) aro 1940.2 aud 1.340.2 for the nucleus aud the ring (Paper I). the Mj, of the uncleus aud the ring are (1.6250.3) & 105 AL. and 140.3) 10 AL... respectively,"," If we adopted the intensity ratios of $^{12}$ CO(J = 2–1) and $^{12}$ CO(J = 1–0) are $\pm$ 0.2 and $\pm$ 0.2 for the nucleus and the ring (Paper I), the $M_{\rm H_2}$ of the nucleus and the ring are $\pm$ $\times$ $^{8}$ $_{\odot}$ and $\pm$ $\times$ $^{9}$ $_{\odot}$, respectively."
" Thus the nucleus aud the ving account for aud of the Mj, within the 2 kpc ciremumuclear region. respectively."," Thus the nucleus and the ring account for and of the $M_{\rm H_2}$ within the 2 kpc circumnuclear region, respectively."
 We caleulate the virial mass of individual clips by where s ds the radius of the chuup derived in Sect. 3.2..," We calculate the virial mass of individual clumps by where ${\it r}$ is the radius of the clump derived in Sect. \ref{sect-clumps},"
" and cg, is the three dimensional iutriusic root mncan square velocity dispersion. which is equal to VIENyon."," and $\sigma_{\rm rms}$ is the three dimensional intrinsic root mean square velocity dispersion, which is equal to $\sqrt{(3/8ln2)}$$\sigma_{\rm int}$."
" Siuce we observe the og, iu one dimension. we need to multiply the observed a; by 3."," Since we observe the $\sigma_{\rm int}$ in one dimension, we need to multiply the observed $\sigma_{\rm int}^{2}$ by 3."
 The radius is adopted with the size of the clumps assundue that ej is isotropic., The radius is adopted with the size of the clumps assuming that $\sigma_{\rm int}$ is isotropic.
 The results are shown iu Table 3.., The results are shown in Table \ref{t.vir}.
 Note that the virial mass is for CALAs not GAICs. and we assume that the CATAs are bounded structures.," Note that the virial mass is for GMAs not GMCs, and we assume that the GMAs are bounded structures."
" The ratio of the virtal mass Au and Mia, are shown in Table 3..", The ratio of the virial mass $M_{\rm vir}$ and $M_{\rm gas}$ are shown in Table \ref{t.vir}.
 We found that the narrow line clumps have Aap Aden. that are more or less about unity. but are larger than unity for the broad line aud dust lane chumps.," We found that the narrow line clumps have $M_{\rm vir}$ $M_{\rm gas}$ that are more or less about unity, but are larger than unity for the broad line and dust lane clumps."
 We plot the general propertics of the molecular eas mass of individual molecular chuups in Fieure 7.., We plot the general properties of the molecular gas mass of individual molecular clumps in Figure \ref{fig-clump-prop}.
 Iu Figure Taa. the histogram of the gas mass inteerated over their size seenis to be a power law with a sharp drop at the low mass cud.," In Figure \ref{fig-clump-prop}a a, the histogram of the gas mass integrated over their size seems to be a power law with a sharp drop at the low mass end."
 This is due to the sensitivity limit. since the corresponding 30 mass limit is ~27«10 AL. for a chump with a diameter of 3733 (the average of the clumps in Table 3)).," This is due to the sensitivity limit, since the corresponding $\sigma$ mass limit is $\sim27\times10^{6}$ $_{\odot}$ for a clump with a diameter of 3 (the average of the clumps in Table \ref{t.vir}) )."
" Iu Figure το, the FWIAL intrinsic Lue widths have a weak correlation with eas mass in the rus."," In Figure \ref{fig-clump-prop}b b, the FWHM intrinsic line widths have a weak correlation with gas mass in the ring."
" In Figure Saa. we show the azimuthal vatlation of Sy, calculated at the cutission peaks of the chuups inteerated over one svuthesized beam."," In Figure \ref{fig-clump-peak}a a, we show the azimuthal variation of $\Sigma_{\rm H_{2}}$ calculated at the emission peaks of the clumps integrated over one synthesized beam."
 If we asstuue that NGC 1097 has a trailing spiral. then the direction of rotation is clockwisefrom cast (07) to north (907).," If we assume that NGC 1097 has a trailing spiral, then the direction of rotation is clockwisefrom east $\degr$ ) to north $\degr$ )."
" The Xj, iu the orbit crowding region is roughly from to157.. and from to 225°."," The $\Sigma_{\rm H_{2}}$ in the orbit crowding region is roughly from to, and from to ."
".. Note that the dust lane chips typically have Xj, similar to the narrow lue ring clumps. which are lower than the broad line ving clumps in the orbit crowding region."," Note that the dust lane clumps typically have $\Sigma_{\rm H_{2}}$ similar to the narrow line ring clumps, which are lower than the broad line ring clumps in the orbit crowding region."
" The average My, of the narrow liue ring and dust laue chuups is ~1500 M. ὃν and ~2300 7 for the broad lue ving chumps."," The average $\Sigma_{\rm H_{2}}$ of the narrow line ring and dust lane clumps is $\sim1800$ $_{\odot}$ $^{-2}$, and $\sim2300$ $_{\odot}$ $^{-2}$ for the broad line ring clumps."
 In Figure 8bb. the velocity dispersion of chuups located at the orbit crowding region aud the dust lane are larger than that of the narrow liue ring chuups.," In Figure \ref{fig-clump-peak}b b, the velocity dispersion of clumps located at the orbit crowding region and the dust lane are larger than that of the narrow line ring clumps."
 The average velocity clispersion of the narrow line clumps is ~50 kins +. and ~90 kins + for the broad line rine/dust lane clumps.," The average velocity dispersion of the narrow line clumps is $\sim50$ km $^{-1}$, and $\sim90$ km $^{-1}$ for the broad line ring/dust lane clumps."
" Caven the similar line brightuesses (Table 2)) of the peaks. the increased “yy, are probably the results of increased lue widths as it is inclicated in equation (2)."," Given the similar line brightnesses (Table \ref{t.clumps}) ) of the peaks, the increased $\Sigma_{\rm H_{2}}$ are probably the results of increased line widths as it is indicated in equation (2)."
" Ou theother haud. since the Ny is proportional to the nuuber deusity of gas aud line-ofsight path G.c.. optical depth). the higher “yy, of the broad line clumps may be due to either a larger uunber density. or a larecr Luc-ofsight path."," On theother hand, since the $N_{\rm H_2}$ is proportional to the number density of gas and line-of-sight path (i.e., optical depth), the higher $\Sigma_{\rm H_2}$ of the broad line clumps may be due to either a larger number density, or a larger line-of-sight path."
 We will discuss these effects iu the discussio-, We will discuss these effects in the discussion.
 To check how the star formation properties are associated with the molecular clumps in the rine. we compare the G-cin radio coutimmum. V-band. aud Pan images with the molecular gas nuages.," To check how the star formation properties are associated with the molecular clumps in the ring, we compare the 6-cm radio continuum, V-band, and ${\alpha}$ images with the molecular gas images."
 The G-cni radio contimuni sources are selected from the intensity peaks (IIuuuucletal.1987) at an angular resolution of 2755.," The 6-cm radio continuum sources are selected from the intensity peaks \citep{hum87}
 at an angular resolution of 5."
 The V-baud clusters (< 13 mas) are selected from the TST E555W. imaee (Barthotal. 1995).., The V-band clusters $<$ 13 mag) are selected from the HST F555W image \citep{barth95}. .
 These V-band selected clustershave typical size of 2 pe (07003). and are suspected to be super star clusters by Barthetal. (1995)...," These V-band selected clustershave typical size of 2 pc 03), and are suspected to be super star clusters by \citet{barth95}. ."
 Pana clusters are identified a our IIST FIstN nuage described, $\alpha$ clusters are identified in our HST F187N image described
which is the desired result.,which is the desired result.
 Iu this derivation we have mace use of the fact that liebt propagates along uull ecodesics. in a space whose curvature enters implicitly twough the function e.," In this derivation we have made use of the fact that light propagates along null geodesics, in a space whose curvature enters implicitly through the function $a$."
 In effect. we have constructed the quadrilateral bounded by two null ravs (on either side) aud two wavelengths (at either cud).," In effect, we have constructed the quadrilateral bounded by two null rays (on either side) and two wavelengths (at either end)."
 Iu a fow-dimenusionally flat (Miukowskv) universe both euds will have the same Iereth and there is no cosmological redshitt., In a four-dimensionally flat (Minkowsky) universe both ends will have the same length and there is no cosmological redshift.
 In a universe with any curvature. the null ravs will diverse (or foctus).," In a universe with any curvature, the null rays will diverge (or focus)."
 Since the scale factor e measures the proper istance between geodesics follow by comoving observers. it also measures the curvature of the universe oo La particular wav. a way couvenieut for computing the ne change in proper distance between the euds of +t1e quadrilateral.," Since the scale factor $a$ measures the proper distance between geodesics followed by comoving observers, it also measures the curvature of the universe in a particular way, a way convenient for computing the net change in proper distance between the ends of the quadrilateral."
 This is the plivsical source of the cosmological redshit: the propagation of null ravs through curved spacetime., This is the physical source of the cosmological redshift: the propagation of null rays through curved spacetime.
 Expansion of a set of comoving observers. or cMf cspace’ itself (whatever that nüght mean). is another effect of curvature. not the cause.," Expansion of a set of comoving observers, or of `space' itself (whatever that might mean), is another effect of curvature, not the cause."
 The concept of space’ certainly cliiuged drasically between classical Newtouian dyvuiuuies and Ceneral Relativity., The concept of `space' certainly changed drastically between classical Newtonian dynamics and General Relativity.
 However similar (sometimes ideutical) the effects of eravity calculated dy each method. the nuderlving wavs of thinking are quite differeit.," However similar (sometimes identical) the effects of gravity calculated by each method, the underlying ways of thinking are quite different."
 The fact that space in GR has an active role in dynamics. rowever. does not mean it las the attributes of a plvsical object.," The fact that space in GR has an active role in dynamics, however, does not mean it has the attributes of a physical object."
 It. does not act like a viscous fluid. drawing all bodies iuto the ITubble flow. even asvinptotically: it does not affect thiugs by ‘expanding. nor bv accelerating this expansion.," It does not act like a viscous fluid, drawing all bodies into the Hubble flow, even asymptotically; it does not affect things by `expanding,' nor by accelerating this expansion."
 Iu fact. if oue looks at space itseIf apart from the associated fluid of cosinological uatter. it js nof at all certain what the expausio1 of space 110111.," In fact, if one looks at space itself apart from the associated fluid of cosmological matter, it is not at all certain what `the expansion of space' means."
 The effect on space on matter is throug Lits curvature., The effect on space on matter is through its curvature.
 Iu the mantra of MGsner. Thorne Wheeler’. uatter tells space how to curve: space tells matterut how to uxwe.," In the mantra of Misner, Thorne $^{5}$, matter tells space how to curve; space tells matter how to move."
 A free particle in a cosmological background noves according to the curvature of space: lie curvature is produced by the background matter., A free particle in a cosmological background moves according to the curvature of space; the curvature is produced by the background matter.
 Motiou (ft the background matter (as long as it renmadus nonrelativistic. contributing nothing to the stress-energv censor) has no direct effect ou curvature: looked at this wav. it would be stranee indeed if a free particle were somehow coustrained to follow that motion.," Motion of the background matter (as long as it remains nonrelativistic, contributing nothing to the stress-energy tensor) has no direct effect on curvature; looked at this way, it would be strange indeed if a free particle were somehow constrained to follow that motion."
 The motion of the backeround. however. isalso determined by he curvature of space.," The motion of the background, however, is determined by the curvature of space."
 This inicaus that the backeround eeodesies 1ieasure the curvature (n a way which is xwtieularlv couvenieut for auvone trving to caculate a redshift)., This means that the background geodesics measure the curvature (in a way which is particularly convenient for anyone trying to calculate a redshift).
 The author would like to thank the referees for claxifviug several poiuts in this exposition (clavity being of particular iuiportance iu this matter)., The author would like to thank the referees for clarifying several points in this exposition (clarity being of particular importance in this matter).
 Especially. the matter of distances (radar. peculiu. proper aud otherwise) require special care. and their coments helped ereatlv.," Especially, the matter of distances (radar, peculiar, proper and otherwise) require special care, and their comments helped greatly."
Evidence for the preseuce of dust in elliptical ealaxies was eiveu by the optical observations of obscured reeious ni sole systems (see for iustauce Bertola et al.,Evidence for the presence of dust in elliptical galaxies was given by the optical observations of obscured regions in some systems (see for instance Bertola et al.
 1985. Vérron Veérron 1988) aud was finally confirmed by he FIR observations of the IRAS satellite (INuapp et al.," 1985, Vérron Vérron 1988) and was finally confirmed by the FIR observations of the IRAS satellite (Knapp et al."
 1989. Roberts et al.," 1989, Roberts et al."
 1991)., 1991).
 The low resolution IRAS data (> l arcmin) have to be processed by adequate echniques to provide detailed (1 arcmin at 100 ja) information about the dust spatial distribution iu nearby elliptical galaxies., The low resolution IRAS data $>$ 1 arcmin) have to be processed by adequate techniques to provide detailed $\sim 1$ arcmin at 100 $\mu$ m) information about the dust spatial distribution in nearby elliptical galaxies.
 In general. ouly the iutegrated flux of he detected source is available in the IBAS lauds.," In general, only the integrated flux of the detected source is available in the IRAS bands."
 Despite the intrinsic lanits due to the low resolution. παρ et al. (," Despite the intrinsic limits due to the low resolution, Knapp et al. ("
1989) showed that a siguificaut fraction (18%) of the nearby E and SO galaxies from the Revised Shapley-Aunes Catalogue (Sandage Tamunann 1981. hereafter RSA) have been detected by IRAS at 60 aud 100 jun at the hnüiting seusitivitv (about 3 times lower than in the IRAS Point Source Catalog).,"1989) showed that a significant fraction $\%$ ) of the nearby E and S0 galaxies from the Revised Shapley-Ames Catalogue (Sandage Tammann 1981, hereafter RSA) have been detected by IRAS at 60 and 100 $\mu$ m at the limiting sensitivity (about 3 times lower than in the IRAS Point Source Catalog)."
 It is not surprising that ellipticals contain dust. since he presence of dust is directly related to the stellar ornuation.," It is not surprising that ellipticals contain dust, since the presence of dust is directly related to the stellar formation."
 But the coexistence of solk particles with he dominant eas conrponenut. which in these galaxies is jeated to 10* IK aud raciates at X-ray wavelengths. is a uatter of discussion.," But the coexistence of solid particles with the dominant gas component, which in these galaxies is heated to $\sim 10^7$ K and radiates at X-ray wavelengths, is a matter of discussion."
 Iu fact. dust eraius should be quickly destroved by sputteriug (Draine Salpeter 1979) when in direct contact with the hot eas.," In fact, dust grains should be quickly destroyed by sputtering (Draine Salpeter 1979) when in direct contact with the hot gas."
 In such an euvironumeut he dust has a lifetime of 10°10* ve., In such an environment the dust has a lifetime of $10^6-10^7$ yr.
 The question thus ollows: where does the dust come frou?, The question thus follows: where does the dust come from?
 At FIR wavelengths (60 aud 100 gan) the thermal chussion is niünlv due to large grains (radius ~ει jan) which mav have different heating sources. e.g.Yo the eeneral interstellar radiation field or OB stars.," At FIR wavelengths (60 and 100 $\mu$ m) the thermal emission is mainly due to large grains (radius $\sim 0.1$ $\mu$ m) which may have different heating sources, e.g. the general interstellar radiation field or OB stars."
 In order to discriuiuate between the two different coutributions and to estimate the weigh of cach of the. several efforts were undertaken (see for istauce Calzetti et al.," In order to discriminate between the two different contributions and to estimate the weight of each of them, several efforts were undertaken (see for instance Calzetti et al."
 1995)., 1995).
 A reasonable approach is to consider the G0 juu flux entirely due to the warm dust (10 I). while the 100 pan flux should be cousidered the result of two contributions: the wart and the cold (~10 Is) dust.," A reasonable approach is to consider the 60 $\mu$ m flux entirely due to the warm dust $\sim$ 40 K), while the 100 $\mu$ m flux should be considered the result of two contributions: the warm and the cold $\sim$ 10 K) dust."
 While ai two-colmponcent dust uodel is used to explain the spectral trend for different types of ealaxies. it is rarely adopted for elliptical galaxies. which are often characterized by weak FIR euuission aud which are uot always detected im all IRAS bands.," While a two-component dust model is used to explain the spectral trend for different types of galaxies, it is rarely adopted for elliptical galaxies, which are often characterized by weak FIR emission and which are not always detected in all IRAS bands."
 Therefore. a single color temperature (from the 60 and 100 san data) is usually taken as the cust temperature.," Therefore, a single color temperature (from the 60 and 100 $\mu$ m data) is usually taken as the dust temperature."
 It follows that no information abou the cust temperature distribution aud the dust spatial distribution is available for elliptical galaxies., It follows that no information about the dust temperature distribution and the dust spatial distribution is available for elliptical galaxies.
 The availability of the ISO data will provide spectra in a wider IR waveleneth range (2.5-210 gan)., The availability of the ISO data will provide spectra in a wider IR wavelength range (2.5-240 $\mu$ m).
 Several attempts to understaud the dust nature aud origin sugeest interesting interpretations * conpanius optical and FIR data (Coudfrooij cle Jong 1995. hereafter GJ95. Tsai Mathews 1995. 1996). wostudvine the stellar conteut. or by using a severe and critical approach to the data (Bregman et al.," Several attempts to understand the dust nature and origin suggest interesting interpretations by comparing optical and FIR data (Goudfrooij de Jong 1995, hereafter GJ95, Tsai Mathews 1995, 1996), by studying the stellar content, or by using a severe and critical approach to the data (Bregman et al."
 1998)., 1998).
 Fiuallv. few elliptical galaxies have been observed at sub-uillinmeter waveleugths by Fich Iodee (1991. 1993).," Finally, few elliptical galaxies have been observed at sub-millimeter wavelengths by Fich Hodge (1991, 1993)."
 €1J95 found that the dust masses determined from the IRAS fiux deusities ave τοσο] an order of magnitude Neher than those determined from optical extinction, GJ95 found that the dust masses determined from the IRAS flux densities are ly an order of magnitude higher than those determined from optical extinction
"Cyganowskietal.(2008). identified more than 3OO ¢ralactic extended 4.5 pmi sources. naming extended: &reen objec‘Is (EGOs) or ""green fuzzies”. for the common coding of the 4.5 jm] band as green in three-color composite Infrared Array Camera images [rom the Telescope.","\citet{cyga08} identified more than 300 Galactic extended 4.5 $\mu$ m sources, naming extended green objects (EGOs) or “green fuzzies”, for the common coding of the [4.5 $\mu$ m] band as green in three-color composite Infrared Array Camera images from the Telescope."
 Xecording to t authors. an EGO is a probable massive voung stellar objec (MYSO) driving outllows.," According to the authors, an EGO is a probable massive young stellar object (MYSO) driving outflows."
 The extended emission in the 4.5 jim band is supposed to be due to IH» Gv=O0. S(9.10.11)) lines and. CO (v=1 0) band. heads. that are exciti by the shock of the outflows propagating in the interstelle mecdium(see MAN'respoctal.2004:Marstonet20tSmh&Rosen2005 .," The extended emission in the 4.5 $\mu$ m band is supposed to be due to $_{2}$ $\nu=0-0$, S(9,10,11)) lines and CO $\nu = 1-0$ ) band heads, that are excited by the shock of the outflows propagating in the interstellar medium (see \citealt{noriega04,marston04,smith05}) )."
 neRecently. DeBuizer&Vacca(20 reported the first spectroscopic identification of the origin of the 4.5 pin emission towards two EGOs using NIB ont1e Gemini North telescope.," Recently, \citet{debuizer10} reported the first direct spectroscopic identification of the origin of the 4.5 $\mu$ m emission towards two EGOs using NIRI on the Gemini North telescope."
 In one of the observed EGOs. they proved that the 4.5 jim emission is due primarily to lines of molecular hydrogen. which are collisionally excited.," In one of the observed EGOs, they proved that the 4.5 $\mu$ m emission is due primarily to lines of molecular hydrogen, which are collisionally excited."
 EGO €135.03|0.35 (hereafter EGO&35)is embedded in a cense molecular cloud at the distance of 3.5 kpe (Petriellaetal.2010). located. on the border of the infrared cust bubble N65 (Churchwellctal.2006.," EGO G35.03+0.35 (hereafter EGOg35)is embedded in a dense molecular cloud at the distance of 3.5 kpc \citep{albert10} located on the border of the infrared dust bubble N65 \citep{church06,church07}."
2007) According to Petriellaetal.(2010) there are several voung stellar object (YSO) candidates. around: N65. being IGOs35 the most prominent source.," According to \citet{albert10} there are several young stellar object (YSO) candidates around N65, being EGOg35 the most prominent source."
 Using UR ancl sub-mam fluxes measured from this EGO. the authors performed. an spectral energy distribution (SED). showing that this source is indeed a massive stellar object at the earlier stages of evolution wi outIlowing activity.," Using IR and sub-mm fluxes measured from this EGO, the authors performed an spectral energy distribution (SED), showing that this source is indeed a massive stellar object at the earlier stages of evolution with outflowing activity."
 The 4.5 pam emission of this EGO has bipolar morphology. with one lobe to the NE and other to the SW.," The 4.5 $\mu$ m emission of this EGO has a bipolar morphology, with one lobe to the NE and the other to the SW."
 The source presents maser emission several molecular lines (see c.g. Forster&Caswell19Caswelletal.1995:Ixurtz&Llofner.," The source presents maser emission of several molecular lines (see e.g. \citealt{forster89,caswell95,kurtz05}) )."
 2005)). Recon CIE;OLI maser emission at 6.7 and 44 ClIHIz has been detected and analvsed by Cyvganowskictal.," Recently, $_{3}$ OH maser emission at 6.7 and 44 GHz has been detected and analysed by \citet{cyga09}. ."
 The 6.7 Cllz, The 6.7 GHz
We have assuued that the S98 sample is α fair suuple of galaxy clusters. aud have considered it to obtain an average nuniber of observed leused sources due to foreground galaxy clusters.,"We have assumed that the S98 sample is a fair sample of galaxy clusters, and have considered it to obtain an average number of observed lensed sources due to foreground galaxy clusters."
 This assumption is likely to be false given biases and systematic effects iu the cluster sample selection., This assumption is likely to be false given biases and systematic effects in the cluster sample selection.
 Our treatment of pointed cluster observations as a series of random untargeted observations is likely to create an additional systematic bias. but such a bias is not expected to uuderestimate the current upper limit.," Our treatment of pointed cluster observations as a series of random untargeted observations is likely to create an additional systematic bias, but such a bias is not expected to underestimate the current upper limit."
 We have also considered a low value for the magnification bias such that the upper Iiuüt on backeround source redshift is overestimated., We have also considered a low value for the magnification bias such that the upper limit on background source redshift is overestimated.
" ΠΠ. for example. the true maeuification bias is 1.FL. then the upper limit on (2) decreases to 2.6 frou 3.1 iu a cosinology of Q,,—0.3 aud Q4=0.7."," If, for example, the true magnification bias is 1.4, then the upper limit on $\langle z \rangle$ decreases to 2.6 from 3.1 in a cosmology of $\Omega_m=0.3$ and $\Omega_\Lambda=0.7$."
 Other uncertainties include the determination of cluster abuudances eiven svsteimatic aud statistical uncertainties involved with the PS calculation. resulting from errors due to ay etc.," Other uncertainties include the determination of cluster abundances given systematic and statistical uncertainties involved with the PS calculation, resulting from errors due to $\sigma_8$ etc."
 We have tried to compcusate for such errors by considering a statistical mucertainty in the derivation of CF(zy)., We have tried to compensate for such errors by considering a statistical uncertainty in the derivation of $\langle F(z_l) \rangle$.
 In general. it is likely that we have overestimated the upper limit. since most of the svstematic effects tend to bias our results such that we wucerestimate the expected lousing rate.," In general, it is likely that we have overestimated the upper limit, since most of the systematic effects tend to bias our results such that we underestimate the expected lensing rate."
 We have derived upper Inuits ou the redshift distribution of ΠΕ sources by comparing statistics of leused sources towards a suuple of galaxy clusters to unuleused. sources., We have derived upper limits on the redshift distribution of submm sources by comparing statistics of lensed sources towards a sample of galaxy clusters to unlensed sources.
" Our derived limits depends oi cosinologyv. iud if Q,,=0.3 and O4=(0.7. as currently sueeested bw various cosmological probes. at the level the average redshift of subnuu sources is less than 3.1."," Our derived limits depends on cosmology, and if $\Omega_m=0.3$ and $\Omega_\Lambda=0.7$, as currently suggested by various cosmological probes, at the level the average redshift of submm sources is less than 3.1."
 Such au upper limit is consistent with the redshift distribution predicted for subnuu sources based on starformation models. where starformation historv remains constaut bevond a redshitt of 1.5. using observed f£u-ufrared aud «πια backeround radiatious.," Such an upper limit is consistent with the redshift distribution predicted for submm sources based on starformation models, where starformation history remains constant beyond a redshift of 1.5, using observed far-infrared and submm background radiations."
 The derived upper luit ou the average redshift is also cousisteut with sugeested redshift ranges based ou colors of plausible optical identifications for subi sources detected towards cluster potentials., The derived upper limit on the average redshift is also consistent with suggested redshift ranges based on colors of plausible optical identifications for submm sources detected towards cluster potentials.
 I would like to thank the anouvinous referee for constructive commuents ou the paper., I would like to thank the anonymous referee for constructive comments on the paper.
In Espositoetal.(2010) we reported on our search for periodicities made on the first Swift//XRT observation (00416485000) by calculating a fast-Fourier-transform power spectrum.,In \citet{eis10} we reported on our search for periodicities made on the first /XRT observation (00416485000) by calculating a fast-Fourier-transform power spectrum.
 A very prominent peak occurs in the spectrum of that observation at 7.5653(4) s (the quoted uncertainty indicates the Fourier period resolution)., A very prominent peak occurs in the spectrum of that observation at 7.5653(4) s (the quoted uncertainty indicates the Fourier period resolution).
" Pulsations were clearly detected also in all the others aand ddatasets, and in the oones up to 2010 October 30 (MJD 55499), when presumably the flux became too low for the PCA sensitivity."," Pulsations were clearly detected also in all the others and datasets, and in the ones up to 2010 October 30 (MJD 55499), when presumably the flux became too low for the PCA sensitivity."
" In order to obtain a refined ephemeris for the longest possible baseline, we studied the pulse phase evolution in these observations by means of an iterative phase-fitting technique (see e.g. Dall’Ossoetal. 2003))."," In order to obtain a refined ephemeris for the longest possible baseline, we studied the pulse phase evolution in these observations by means of an iterative phase-fitting technique (see e.g. \citealt{dallosso03}) )."
 The fits were carried out in the range 2-10 keV with a x? minimisation approach using (James&Roos1975)., The fits were carried out in the range 2–10 keV with a $\chi^2$ minimisation approach using \citep{james75}.
. Throughout the period covered by useful observations (~225 days) the relative phases and amplitudes were such that the phase evolution of the signal could be followed unambiguously., Throughout the period covered by useful observations $\sim$ 225 days) the relative phases and amplitudes were such that the phase evolution of the signal could be followed unambiguously.
" A second-order polynomial, as employed in the recent analysis by Gógüsetal.(2010a) over the first ~47 days, provides an unacceptable fit to the data, with x2=9.06 for 86 dof."," A second-order polynomial, as employed in the recent analysis by \citet{gogus10short} over the first $\sim$ 47 days, provides an unacceptable fit to the data, with $\chi^2_\nu=9.06$ for 86 dof."
 We tried higher-order polynomials until the addition of a further (higher-order) term was not statistically significant at more than 30 with respect to the null hypothesis (as evaluated by the Fisher test)., We tried higher-order polynomials until the addition of a further (higher-order) term was not statistically significant at more than $\sigma$ with respect to the null hypothesis (as evaluated by the Fisher test).
" The outcome of this process was a fourth-order polynomial (the improvement obtained in the fit with a fifth-order polynomial has a statistical significance of only 2.60), which we used to fit the phase shifts."," The outcome of this process was a fourth-order polynomial (the improvement obtained in the fit with a fifth-order polynomial has a statistical significance of only $\sigma$ ), which we used to fit the phase shifts."
 The resulting phase-coherent solution is given in and plotted in Fig. 4;;, The resulting phase-coherent solution is given in \\ref{timing-fit} and plotted in Fig. \ref{residuals};
" the best fit (x2=1.07 for 84 dof) gives v=0.132180571(7) Hzand v=—6.0(5)x10 Hz s!, assuming MJD 55274.0 as reference epoch."," the best fit $\chi^2_\nu=1.07$ for 84 dof) gives $\nu=0.132\,180\,571(7)$ Hz and $\dot{\nu}=-6.0(5)\times10^{-14}$ Hz $^{-1}$, assuming MJD 55274.0 as reference epoch."
 We have checked14 that the positional uncertainty (03; Gógügetal. 2010a)) does not significantly affect the rotational parameters resulting from our analysis., We have checked that the positional uncertainty $0\farcs3$; \citealt{gogus10short}) ) does not significantly affect the rotational parameters resulting from our analysis.
 In Fig., In Fig.
 5 we show the three llight curves obtained folding the high time-resolution (nominal frame time of 73.4 ms) EPIC-pn data at our phase-coherent ephemeris., \ref{pprofile} we show the three light curves obtained folding the high time-resolution (nominal frame time of 73.4 ms) EPIC-pn data at our phase-coherent ephemeris.
" The pulse profile is sinusoidal and the pulsed fraction, that we define as the semi-amplitude of sinusoidal modulation divided by the mean source count rate, was consistent through the first two observations [(69.0+1.4)% and (67.4+1.5)%,, respectively], while it decreased to (57.0+1.7)% in the third one."," The pulse profile is sinusoidal and the pulsed fraction, that we define as the semi-amplitude of sinusoidal modulation divided by the mean source count rate, was consistent through the first two observations $69.0\pm1.4$ and $67.4\pm1.5$, respectively], while it decreased to $57.0\pm1.7$ in the third one."
 We have investigated the morphology of the pulse phase distribution as a function of energy by comparing the ppulse profiles in different energy bands., We have investigated the morphology of the pulse phase distribution as a function of energy by comparing the pulse profiles in different energy bands.
" The measured pulsed fractions in the three observations are (66+2)%, (64+2)%, and (56+2)% in the soft (2-5 keV) energy band, and (75+2)%, (71+2)%, and (59+3)% in the hard (5-10 keV) band."," The measured pulsed fractions in the three observations are $(66\pm2)\%$, $(64\pm2)\%$, and $(56\pm2)\%$ in the soft (2–5 keV) energy band, and $(75\pm2)\%$, $(71\pm2)\%$, and $(59\pm3)\%$ in the hard (5–10 keV) band."
" Apart from this marginal indication for an increasing trend of the pulsed fraction with energy, no significant pulse shape variations (such as phase shifts of the maxima) were found as a function of energy by cross-correlating or comparing through a two-sided Smirnov test the soft and hard folded profiles."," Apart from this marginal indication for an increasing trend of the pulsed fraction with energy, no significant pulse shape variations (such as phase shifts of the maxima) were found as a function of energy by cross-correlating or comparing through a two-sided Kolmogorov--Smirnov test the soft and hard folded profiles."
 The 2-10 keV pulsed fractions measured in the individual observations obtained with aand aare shown in the bottom panel of Fig., The 2–10 keV pulsed fractions measured in the individual observations obtained with and are shown in the bottom panel of Fig.
 3 (we did not consider the ddata as the non-imaging PCA instrument does not ensure reliable background subtraction)., \ref{history} (we did not consider the data as the non-imaging PCA instrument does not ensure reliable background subtraction).
 A constant fit of the pulsed fraction values derived with (and shown in Fig. 3)), A constant fit of the pulsed fraction values derived with (and shown in Fig. \ref{history}) )
" does not adequately describe the data (x2,=4.52 for 24 dof); however no particular trend is apparent in the pulsed fraction evolution and the simple functions we tried do not yield significantly better fits.", does not adequately describe the data $\chi^2_\nu=4.52$ for 24 dof); however no particular trend is apparent in the pulsed fraction evolution and the simple functions we tried do not yield significantly better fits.
The SAIC has been the subject of extensive monitoring using the RAPE Proportional Counter Array (PCA) over the last LO vears.,The SMC has been the subject of extensive monitoring using the RXTE Proportional Counter Array (PCA) over the last 10 years.
 The PCA instrument has a Full Width Llalf Maxinuun CENLIM) field of view of 1 and data in the enerev range 3-10 keV were used.," The PCA instrument has a Full Width Half Maximum (FWHM) field of view of $1\,^{\circ}$ and data in the energy range 3-10 keV were used."
 Most. of the observations were pointed at the Dar region of the SALC where the majority of the known X-ray. pulsar systems are located., Most of the observations were pointed at the Bar region of the SMC where the majority of the known X-ray pulsar systems are located.
 Sources were identified [rom their pulse periods using Lonib-Scarele (Lomb 1976. Scargle 1982) power spectral analysis of the data sets.," Sources were identified from their pulse periods using Lomb-Scargle (Lomb 1976, Scargle 1982) power spectral analysis of the data sets."
 Laveock et al. (, Laycock et al. (
2005) and Galache et al. (,2005) and Galache et al. (
2008) present. full details of the data analvsis approach that has been used to determine which pulsars were active during each observation.,2008) present full details of the data analysis approach that has been used to determine which pulsars were active during each observation.
 In their work. for each X-ray. outburst. the pulse amplitude and. history of the pulse periods were determined.," In their work, for each X-ray outburst, the pulse amplitude and history of the pulse periods were determined."
 Vhose results are used. in. this work., Those results are used in this work.
 Since a database of 210. vears of observations of the SAIC exists ib was therefore possible to use these data to search for evidence of spin period changes in the svstems., Since a database of $\ge$ 10 years of observations of the SMC exists it was therefore possible to use these data to search for evidence of spin period changes in the systems.
 The PCA is a collimated instrument. therefore interpreting the strength and significance. of the signal depends upon the position of the source within the field. of view.," The PCA is a collimated instrument, therefore interpreting the strength and significance of the signal depends upon the position of the source within the field of view."
 In all the objects presented. here the target was assumed to be located at the position of the known optical counterpart., In all the objects presented here the target was assumed to be located at the position of the known optical counterpart.
 Only observations that hac a collimator response 225'4 and a detection signilicance of 99% were used in this work., Only observations that had a collimator response $\ge$ and a detection significance of $\ge$ were used in this work.
 A total of 24 sources were chosen for this study., A total of 24 sources were chosen for this study.
 In each case three possible measurements of period changes were obtained: Li was not always possible to determine both a Short? and Lonel? for every source in this work due to several possible reasons: one being the observational coverage and another the activity history of the system., In each case three possible measurements of period changes were obtained: It was not always possible to determine both a $\dot{P}$ and $\dot{P}$ for every source in this work due to several possible reasons; one being the observational coverage and another the activity history of the system.
 Details of the recorded spin period changes are given in Table 1.., Details of the recorded spin period changes are given in Table \ref{rxte}.
 Full records of the behaviour of cach source may be found. in Calache et al. (, Full records of the behaviour of each source may be found in Galache et al. (
2008).,2008).
 The strong link between the equilibrium spin. period and the rate of change of spin period seen during outbursts is shown in Figure 2.., The strong link between the equilibrium spin period and the rate of change of spin period seen during outbursts is shown in Figure \ref{fig2}.
 In this figure. the straight line represents D — kP? αν predicted for accretion from a disk on to a neutron star (see Equation 15 in Chosh Lamb. 1979).," In this figure, the straight line represents $\dot{P}$ = $P^{2}$ - as predicted for accretion from a disk on to a neutron star (see Equation 15 in Ghosh Lamb, 1979)."
 ]t is interesting to note that the four spin-down systems (SNPS.80. SXP59.0. SNPI44 SXDP1323) fit comfortably n this relationship alongside the much larger number of μα?nn-up svstenis.," It is interesting to note that the four spin-down systems (SXP8.80, SXP59.0, SXP144 SXP1323) fit comfortably on this relationship alongside the much larger number of spin-up systems."
" To pursue the understanding of the aceretion process further. a value for the X-ray luminositv. £,. is neeclec uring cach outburst."," To pursue the understanding of the accretion process further, a value for the X-ray luminosity, $L_x$, is needed during each outburst."
 So the X-ray luminosity was calculated. from. the observed peak pulse amplitude in counts/pcu/s that occurred. during. the outburst tha xoduced the Short? values. listed in. Table 1., So the X-ray luminosity was calculated from the observed peak pulse amplitude in counts/pcu/s that occurred during the outburst that produced the $\dot{P}$ values listed in Table \ref{rxte}.
 This amplitude is converted to. luminosity assuming that INTE count/peu/s = 10% erg/s at the distance of Gükpc to the SMC (though the depth of the sources within he SMC is unknown and alfect this distance by up to +10kpe)., This amplitude is converted to luminosity assuming that 1 RXTE count/pcu/s = $\times 10^{37}$ erg/s at the distance of 60kpc to the SMC (though the depth of the sources within the SMC is unknown and affect this distance by up to $\pm$ 10kpc).
 Phe X-ray spectrum was assumed to be a power aw with a photon index = 1.5 and an Ny=610ο7., The X-ray spectrum was assumed to be a power law with a photon index = 1.5 and an $N_{H}=6\times10^{20}cm^{-2}$.
 Furthermore it was assumed that there was an average οσο fraction. of for all the systems anc hence the correct total [lux is 3 times the pulse component., Furthermore it was assumed that there was an average pulse fraction of for all the systems and hence the correct total flux is 3 times the pulse component.
 Thus he values shown in Table 1... were determined using the relationship: 93 cre where 2 = RXTE counts bs ! Note that though the values for L. obviously scale, Thus the values shown in Table \ref{rxte} were determined using the relationship: $L_X$ = 0.4 $\times\ 10^{37}$ $\times$ $R$ erg $^{-1}$ where $R$ = RXTE counts $^{-1}$ $^{-1}$ Note that though the values for $L_x$ obviously scale
at removing CSS objects. most of which seem to show a Hattening in their spectra at low frequencies. and. therefore emit relatively low power at 151 MlIz in the rest. frame.,"at removing CSS objects, most of which seem to show a flattening in their spectra at low frequencies, and therefore emit relatively low power at 151 MHz in the rest frame."
" This produces a ""cleaner"" looking correlation. indeed: so good that we have felt confident using parametric methods to evaluate the slope of the correlation."," This produces a “cleaner” looking correlation, indeed so good that we have felt confident using parametric methods to evaluate the slope of the correlation."
 Phe slopes were determined. using the Estimation-Maximisation method. as implemented in the task (Isobe Feigelson 1990) which is able to deal with the size limits in the data by assuming a normal cistribution for the correlation residuals., The slopes were determined using the Estimation-Maximisation method as implemented in the task (Isobe Feigelson 1990) which is able to deal with the size limits in the data by assuming a normal distribution for the correlation residuals.
 As only one NEC* and four objects had size limits this is unlikely to have allected the correlations., As only one NEC* and four objects had size limits this is unlikely to have affected the correlations.
 The correlation. slopes for the NEC* |. 307 samples are [latter than the Blundell et ((1999) values., The correlation slopes for the NEC* + 3C* samples are flatter than the Blundell et (1999) values.
 Phoush his is only marginallvo significant.C» statistically. it is clear rom examination of 99 that this cllect is due to arger numbers of high redshift CSS sources in the 151-MllIz (observed-frame) selected. samples.," Though this is only marginally significant statistically, it is clear from examination of 9 that this effect is due to larger numbers of high redshift CSS sources in the 151-MHz (observed-frame) selected samples."
 “Phe small size of these objects (C; 30kpc) means that most of them are probably confined by the interstellar medium of the host galaxy. anc his high density environment. will make the conversion of jet kinetic power into radio emission. particularly. efficient.," The small size of these objects $\stackrel{<}{_{\sim}} 30$ kpc) means that most of them are probably confined by the interstellar medium of the host galaxy, and this high density environment will make the conversion of jet kinetic power into radio emission particularly efficient."
 A steep jet luminosity function IxIxaiser Alexander 998) will then lead to them being preferentially include in [lux-limited. samples. particularly at high rest-frame selection frequencies where their spectra are vet το star Hattening.," A steep jet luminosity function Kaiser Alexander 1998) will then lead to them being preferentially included in flux-limited samples, particularly at high rest-frame selection frequencies where their spectra are yet to start flattening."
 For the sample we obtain slopes very similar to those of Blundell ct (probably in part because we use the same LIUL sample for comparison)., For the sample we obtain slopes very similar to those of Blundell et (probably in part because we use the same LRL sample for comparison).
 We thus find that the dependence of radio source size on redshift is fairlv weak. with radio source size evolving approximately with the scale factor of the Universe. (1.12).!.," We thus find that the dependence of radio source size on redshift is fairly weak, with radio source size evolving approximately with the scale factor of the Universe, $(1+z)^{-1}$."
 In particular the upper envelope of the size distribution is very weakly dependent on redshift., In particular the upper envelope of the size distribution is very weakly dependent on redshift.
 The trend. of decreasing size with redshift has usually been assumed to be a straightforward. consequence of increasing environmental density with redshift. combined with a finite source lifetime SSubramanian Swarup 1990).," The trend of decreasing size with redshift has usually been assumed to be a straightforward consequence of increasing environmental density with redshift, combined with a finite source lifetime Subramanian Swarup 1990)."
 Llowever. the strong selection pressures. associated with the steep radio luminosity function mean that what is being measured is the size at which. for a given redshift. the average racio source reaches its maximum luminosity.," However, the strong selection pressures associated with the steep radio luminosity function mean that what is being measured is the size at which, for a given redshift, the average radio source reaches its maximum luminosity."
 Hence if as argued. by Dlundell et al. (," Hence if, as argued by Blundell et al. ("
1999). high. redshift. high luminosity sources reach their maximum luminosities earlier in their lives than low redshift ones due to a combination highere inverse-Compton losses and a correlation of injection spectral index with radio Luminosity. this is also a possible explanation.,"1999), high redshift, high luminosity sources reach their maximum luminosities earlier in their lives than low redshift ones due to a combination higher inverse-Compton losses and a correlation of injection spectral index with radio luminosity, this is also a possible explanation."
" Itecently. Ixaiser Alexander (1998) have also suggested. that if either the lifetimes of radio sources or the shape of the density. distribution of their environments allects the powers ofthe jets they can explain both the epoch dependence of linear size and the rapid evolution in. the luminosity function of the FRI population. at least. within 3€. One observational way in which modelling uncertainties could. be addressed. is by finding ""σπα raclio sources and estimating their ages: this may be the only way of constraining whether finite source lifetimes are likely to play an important rolle in evolutionary scenarios."," Recently, Kaiser Alexander (1998) have also suggested that if either the lifetimes of radio sources or the shape of the density distribution of their environments affects the powers of the jets they can explain both the epoch dependence of linear size and the rapid evolution in the luminosity function of the FRII population, at least within 3C. One observational way in which modelling uncertainties could be addressed is by finding “dead” radio sources and estimating their ages; this may be the only way of constraining whether finite source lifetimes are likely to play an important rôlle in evolutionary scenarios."
 The influence of the CSS sources in driving the recdshilt- correlation in Ilux-density limited samples of ETUL racio sources could explain why much steeper slopes are seen when samples are selected at higher frequency. because these samples (vpically contain substantially larger. fractionsga of CSS sources than LItL.," The influence of the CSS sources in driving the redshift-size correlation in flux-density limited samples of FRII radio sources could explain why much steeper slopes are seen when samples are selected at higher frequency, because these samples typically contain substantially larger fractions of CSS sources than LRL."
 Phe CSS sources could also account for the radio luminosity dependence too., The CSS sources could also account for the radio luminosity dependence too.
 The most likely mechanisms for spectral Uattening (thermal absorption: Acknell ct 11997. or synchrotron scll-absorption) are both likely to be Luminosity dependent. thermal absorption through the emission line luminosity radio Luminosity correlation and svnchrotron self-absorption more clirectly.," The most likely mechanisms for spectral flattening (thermal absorption; Bicknell et 1997, or synchrotron self-absorption) are both likely to be luminosity dependent, thermal absorption through the emission line luminosity – radio luminosity correlation and synchrotron self-absorption more directly."
 Ixapahi (1989) removed. CSS sources from his study of size evolution. and found that the best-fit value of 7 was reduced (from 3.52:0.5 to 2.50.5) and the dependence on luminosity was weakened. (ο was reduced from 0.4—0.1 to 0.2£ 0.1).," Kapahi (1989) removed CSS sources from his study of size evolution, and found that the best-fit value of $\eta$ was reduced (from $3.5 \pm 0.5$ to $2.5 \pm 0.5$ ) and the dependence on luminosity was weakened $\epsilon$ was reduced from $0.4 \pm 0.1$ to $0.2 \pm 0.1$ )."
 These correlations are still steeper than we observe for the NEC* | 3€7 sample (7=1.6c0.3. οz 0). although the errors on the slopes both. in our study. ancl that of Ixapahi's are sullicient for the discrepancies to be due to," These correlations are still steeper than we observe for the NEC* + 3C* sample $\eta=1.6 \pm 0.3$, $\epsilon \approx 0$ ), although the errors on the slopes both in our study and that of Kapahi's are sufficient for the discrepancies to be due to"
include au explicit fragiucutation process.,include an explicit fragmentation process.
 Although it may seem uuwielkdv. it is actually straightforward to construct a plivsicallv iiotivated uodel where these processes are unavoldablv realized.," Although it may seem unwieldy, it is actually straightforward to construct a physically motivated model where these processes are unavoidably realized."
" Of course. in order to describe magnetic field evolution both at large euerev and leneth scales, aud at the extremely high magnetic Revnuolds nunber relevant to the corona. it is necessary to uake simplifications of the underline plyvsical aws."," Of course, in order to describe magnetic field evolution both at large energy and length scales, and at the extremely high magnetic Reynolds number relevant to the corona, it is necessary to make simplifications of the underlying physical laws."
 Nevertheless. the model does retain certain physical features that are known to be important.," Nevertheless, the model does retain certain physical features that are known to be important."
 Primarily. it keeps track of the eeometrical constraints that the hieh conductivity of the corona aud the phiyvsies of reconnection inpose on maenetic field evolution.," Primarily, it keeps track of the geometrical constraints that the high conductivity of the corona and the physics of reconnection impose on magnetic field evolution."
 Magnetic flux is frozen iuto the plasiia auc is constrained to move with it., Magnetic flux is frozen into the plasma and is constrained to move with it.
 IHowever. when magnetic ficld gradieuts are steep reconnection can occur. changing the gcometry of he magnetic field structure.," However, when magnetic field gradients are steep reconnection can occur, changing the geometry of the magnetic field structure."
 Sccoudly. diffusion of ootpoints and flux emergence are. cousidered to ος the dominant sources driving coronal magnetic held energy.," Secondly, diffusion of footpoints and flux emergence are considered to be the dominant sources driving coronal magnetic field energy."
 The fundamental eutitv iu the model is a directed loop which traces the iuidline of a fiux ube and is anchored to a surface at two opposite olazity footpoints., The fundamental entity in the model is a directed loop which traces the midline of a flux tube and is anchored to a surface at two opposite polarity footpoints.
 A of collectionthese loops aud heir footpoiuts elves a representation of a coronal field naeneticstructure., A collection of these loops and their footpoints gives a representation of a coronal magnetic field structure.
 It is able to describe fields hat are very complicated or interwoven. like the naguetic carpet.," It is able to describe fields that are very complicated or interwoven, like the magnetic carpet."
 A snapshot of a coufiguration iu he steady-state is shown iu Fig. 1.., A snapshot of a configuration in the steady-state is shown in Fig. \ref{loops}.
 The uuuber of oops connecting any pair of footpoiuts is indicated * a color-codius., The number of loops connecting any pair of footpoints is indicated by a color-coding.
 It is evident that both the iuuber of loops attached to a footpoiut aud the streneth of those loops vary over a broad range., It is evident that both the number of loops attached to a footpoint and the strength of those loops vary over a broad range.
 Loops injected at small length scales are stretched aud shrunk as their footpoiuts. diffuse over the surface., Loops injected at small length scales are stretched and shrunk as their footpoints diffuse over the surface.
 Loops subimierge when their ootpoits approach closelv., Loops submerge when their footpoints approach closely.
" Nearby footpoiuts of the same polarity coalesce. to form maguetic rasnients, which can themselves coalesce to orn ever larger concentrations of flux. such as pores and sunspots."," Nearby footpoints of the same polarity coalesce, to form magnetic fragments, which can themselves coalesce to form ever larger concentrations of flux, such as pores and sunspots."
 Conversely. opposite »olaritv footpoiuts may cancel (though see section 2.3).," Conversely, opposite polarity footpoints may cancel (though see section 2.3)."
 Loops can reconnect when they collice in three dimensional space. thereby releasing nagnetic energev.," Loops can reconnect when they collide in three dimensional space, thereby releasing magnetic energy."
 Reconnection of flux tubes or concentration cancelation may trigger a cascade of further recounectiou. represcuting a flare.," Reconnection of flux tubes or concentration cancelation may trigger a cascade of further reconnection, representing a flare."
 Each flux tube is represented by au infinitesimally thin. directed loop which traces its midline.," Each flux tube is represented by an infinitesimally thin, directed loop which traces its midline."
 A coronal magnetic field is described bv a configuration of many of these loops., A coronal magnetic field is described by a configuration of many of these loops.
 For simplicity cach loop is cousidered to be a semicircle emereme perpendicular to the Gey) plane., For simplicity each loop is considered to be a semicircle emerging perpendicular to the $(xy)$ plane.
 Every loop has a positive footpoiut. where magnetic fiux emerges from the photosphere. aud a negative oue. Where flux returns.," Every loop has a positive footpoint, where magnetic flux emerges from the photosphere, and a negative one, where flux returns."
 The size of the svstem in the Gey) plane. which represeuts a region of the photospheric surface. is Lo.EL.," The size of the system in the $(xy)$ plane, which represents a region of the photospheric surface, is $L
\times L$."
" Loops are labeled by an integer. s. and the positions of the two footpoiuts of the nth loop arelabeled in the Cry) plane by vr! and r,."," Loops are labeled by an integer, $n$, and the positions of the two footpoints of the $n$ th loop arelabeled in the $(xy)$ plane by ${\bf r}^+_n$ and ${\bf r}^-_n$."
" The footpoiut separation of a flux tube is thon d,=|r/—r,| and the length of a flux tube is 7,=4|r!x,| A footpoiut labels the center of a magnetic concentration."," The footpoint separation of a flux tube is then $d_n= |{\bf
r}^+_n -{\bf r}^-_n|$ and the length of a flux tube is $l_n= {\pi\over 2} |{\bf
r}^+_n -{\bf r}^-_n|$ A footpoint labels the center of a magnetic concentration."
 Due to coalescence. described in step 6. a footpoiut can lave amore than one loop attached to it.," Due to coalescence, described in step 6, a footpoint can have more than one loop attached to it."
 This meaus that the magnetic flux from amy concentration. or the flux connecting auv two concentrations. can be arbitrarily large. despite he fact that the individual loops. defined im step l represent sinall. quantized uuits of flux.," This means that the magnetic flux from any concentration, or the flux connecting any two concentrations, can be arbitrarily large, despite the fact that the individual loops, defined in step 1 represent small, quantized units of flux."
 We do rot attach a surface area to the concentrations. just as we do not attach a width to flux tubes.," We do not attach a surface area to the concentrations, just as we do not attach a width to flux tubes."
 A concentration in the model is completely described * a suele footpoiut located at its center aud he number of loops counected to the footpoiut represents the total flux of the fragment., A concentration in the model is completely described by a single footpoint located at its center and the number of loops connected to the footpoint represents the total flux of the fragment.
 To conrpare with observations. one loop iu the model should be set equal to the minim threshold for flux to be included in the data set.," To compare with observations, one loop in the model should be set equal to the minimum threshold for flux to be included in the data set."
 This aud the next term represent the driving terms for magnetic fiux., This and the next term represent the driving terms for magnetic flux.
 The diffuxou of footpoiuts. and therefore the concentrations they represent. is described as a random walk on the two dimensional (ry) pluie.," The diffusion of footpoints, and therefore the concentrations they represent, is described as a random walk on the two dimensional $(xy)$ plane."
 At an update step. an arbitrary footpoiutis chosen at random and its position is moved. r»|Ar.," At an update step, an arbitrary footpointis chosen at random and its position is moved, ${\bf r} \rightarrow {\bf r} + \Delta {\bf
r}$."
 The vector Ar has leneth and anele chosen randomly from uniforii distributions between 0 aud 1. aud 0 aud 27. respectively.," The vector $\Delta{\bf r}$ has length and angle chosen randomly from uniform distributions between 0 and $1$ , and 0 and $2\pi$ , respectively."
 If the, If the
lines were covered by using two correlator units of 80 MHz with a spectral resolution of 0.3125 MHz (0.97 Κπις-!)) and by two correlator units of 320 MHz.,lines were covered by using two correlator units of 80 MHz with a spectral resolution of 0.3125 MHz (0.97 ) and by two correlator units of 320 MHz.
" The latter were used, free of line, to produce the continuum image with a total bandwith of ~450 MHz."," The latter were used, free of line, to produce the continuum image with a total bandwith of $\sim$ 450 MHz."
" At 1 mm, the CH3OH 5;— 4; v,=0 lines were observed with two units of 160 MHz with a spectral resolution of 1.250 MHz (1.55 kms7!))."," At 1 mm, the $_{3}$ OH $_{k} \rightarrow$ $_{k}$ $v_{t} = 0$ lines were observed with two units of 160 MHz with a spectral resolution of 1.250 MHz (1.55 )."
" The remaining units of 80, 160, and 320 MHz were placed in such a way that a frequency range free of lines could be used to measure the continuum flux."," The remaining units of 80, 160, and 320 MHz were placed in such a way that a frequency range free of lines could be used to measure the continuum flux."
 The total bandwidth of both sidebands was ~700 MHz., The total bandwidth of both sidebands was $\sim$ 700 MHz.
 The phase and amplitude were calibrated with observations of the object 1730-130., The phase and amplitude were calibrated with observations of the object 1730-130.
 The bandpass calibration was done with 3C 273., The bandpass calibration was done with 3C 273.
 MWC 349 was used as primary flux calibrator of the 3 and 1 mm data (see Table 1))., MWC 349 was used as primary flux calibrator of the 3 and 1 mm data (see Table \ref{Tobs}) ).
" We estimate the final flux density accuracy to be ~5% and for the 3mm and 1mm data, respectively."," We estimate the final flux density accuracy to be $\sim$ and for the 3mm and 1mm data, respectively."
 Continuum images were subtracted from the line data in the visibility plane., Continuum images were subtracted from the line data in the visibility plane.
" The combination of the C and D configurations provides angular scales from 1""788-12""66 (1 mm) and 4744-32755 (3 mm), i.e., providing information on spatial scales of 0.03—0.22 pc and 0.08—0.57 pc, respectively."," The combination of the C and D configurations provides angular scales from 6 (1 mm) and 5 (3 mm), i.e., providing information on spatial scales of 0.03–0.22 pc and 0.08–0.57 pc, respectively."
 The calibration and data reduction were made within the GILDAS at IRAM Grenoble., The calibration and data reduction were made within the GILDAS at IRAM Grenoble.
 Images were created with natural weighting and CLEANed using the standard Hóggbom algorithm., Images were created with natural weighting and CLEANed using the standard Höggbom algorithm.
" G11.11P1 was observed in the 2; —1; νι=0 (3 mm) and 5,— 44 (1 mm) CH30H bands with the IRAM 30m telescope.", G11.11P1 was observed in the $_{k}\rightarrow$ $_{k}$ $v_{t} = 0$ (3 mm) and $_{k} \rightarrow$ $_{k}$ (1 mm) $_{3}$ OH bands with the IRAM 30m telescope.
 An area of x was mapped with the SIS B100 receiver at 3 mm and x with the HERA receiver at 1 mm with a bandwidth of «140 MHz., An area of $\times$ was mapped with the SIS B100 receiver at 3 mm and $\times$ with the HERA receiver at 1 mm with a bandwidth of $\sim$ 140 MHz.
" Sampling at 3 mm was of while at 1 mm was of6"".", Sampling at 3 mm was of while at 1 mm was of.
. Conversion from antenna temperature to main-beam brightness temperature (Twp) was performed by using a beam efficiency of 0.78 at 3 mm and 0.52 at | mm., Conversion from antenna temperature to main-beam brightness temperature $T_{\rm MB}$ ) was performed by using a beam efficiency of 0.78 at 3 mm and 0.52 at 1 mm.
" The rms at 1 mm is 0.3 K and at 3 mm is 0.02 K. The data were reduced with the CLASS program, part of the GILDAS software package."," The $rms$ at 1 mm is 0.3 K and at 3 mm is 0.02 K. The data were reduced with the CLASS program, part of the GILDAS software package."
and galactic bulges.,and galactic bulges.
 This diagnostic should be more effective in metal-rich populations (|M/II] =—0.5) because there are more RGBB stars at higher metallicity (for fixed age and helium enrichment) rendering anv given [huetuation more statistically significant., This diagnostic should be more effective in metal-rich populations ([M/H] $\gtrsim -0.5$ ) because there are more RGBB stars at higher metallicity (for fixed age and helium enrichment) rendering any given fluctuation more statistically significant.
 However. recent research has demonstrated that theory may inaccurately overestimate the RGBB huninosity bv 20.2 mag (Cassisietal.2011).. which was estimated bv comparing the dillerence in brightness between (he RGBB aud (the main-sequence turn-olf for a sample of 15 Galactic elobular clusters. aud that predicted. given (the estimated. ages and measured metallicities of the cluster.," However, recent research has demonstrated that theory may inaccurately overestimate the RGBB luminosity by $\sim$ 0.2 mag \citep{2011A&A...527A..59C}, which was estimated by comparing the difference in brightness between the RGBB and the main-sequence turn-off for a sample of 15 Galactic globular clusters, and that predicted given the estimated ages and measured metallicities of the cluster."
 If (here are small svstematic errors in (he theoretical predictions for RGDD unminositw. (here may also be errors in the predictions for the RGDD lifetime.," If there are small systematic errors in the theoretical predictions for RGBB luminosity, there may also be errors in the predictions for the RGBB lifetime."
 On that vole. our brightness gradients for the IB. (0.072£0.009) mag !. and for the RGDD. (0.0832:0.023) mag dex... vield a ratio of dVice;/dVig = (1.1540.35).," On that note, our brightness gradients for the HB, $(0.072 \pm 0.009)$ mag $^{-1}$ , and for the RGBB, $(0.083 \pm 0.023)$ mag $^{-1}$, yield a ratio of $dV_{RGBB}/dV_{HB}$ = $(1.15 \pm 0.35)$."
 This is unfortunately insufficientlv precise to lest a stellar theory. prediction of Cassisi&Salaris(1997).., This is unfortunately insufficiently precise to test a stellar theory prediction of \citet{1997MNRAS.285..593C}.
 Their nodels predicted that that the RGBB Iumninosity should respond more steeply to varying initial helium abundance than the ZAIID bhuninosity. and as such the difference in their brightuess should decrease by —0.011 mag for every increase of 0.01 in Y. suggesting a ralio dVpyc;pgp/dViig~2.," Their models predicted that that the RGBB luminosity should respond more steeply to varying initial helium abundance than the ZAHB luminosity, and as such the difference in their V-band brightness should decrease by $\sim$ 0.011 mag for every increase of 0.01 in Y, suggesting a ratio $dV_{RGBB}/dV_{HB} \sim 2$."
 A more precise comparison is within reach if uniform photometry is obtained over the entirety of the cluster., A more precise comparison is within reach if uniform photometry is obtained over the entirety of the cluster.
 DAMN andAG were partially supported by the NSF erant AST-O757888., DMN andAG were partially supported by the NSF grant AST-0757888.
andP3**(k)-,and _v k).
"ho(B,,o Finally, the Fisher matrix »).(63)for 6 and o, is specified by where p,=f,0, indicates the parameter vector and the partial derivatives can be taken either numerically or in some cases analytically (if for example the distribution f(v) is either a Gaussian or an exponential)."," Finally, the Fisher matrix for $\beta$ and $\sigma_v$ is specified by where $p_a=\beta,\sigma_v$ indicates the parameter vector and the partial derivatives can be taken either numerically or in some cases analytically (if for example the distribution $f(v)$ is either a Gaussian or an exponential)."
" We use as limits of the sum kmin=2-/(V.)U/9 and utilize a variable maximum wavenumber kma, to test the magnitude and trend of the errors."," We use as limits of the sum $k_{\mathrm
{min}}=2\pi/(V_s)^{(1/3)}$ and utilize a variable maximum wavenumber $k_{max}$ to test the magnitude and trend of the errors."
" In reffig:fisher we show the percent fractional error for marginalized errors on ( for a survey centered at redshift z=1, of volume V,=1(h-! full sky, with enough dark matter particles to Gpc),dominate shot noise."," In \\ref{fig:fisher} we show the percent fractional error for marginalized errors on $\beta$ for a survey centered at redshift $z=1$, of volume $V_s=1\,\left( h^{-1}\,\mathrm{Gpc} \right)^3$, full sky, with enough dark matter particles to dominate shot noise."
 Our Fisher matrix study is based on Fourier space and to relate it to configuration space we need to link wavenumbers to distances., Our Fisher matrix study is based on Fourier space and to relate it to configuration space we need to link wavenumbers to distances.
 We find that the scaling kimax=2/Tmin gives the Fisher errors of similar magnitude and trend to the errors estimated from the simulations., We find that the scaling $k_{\mathrm{max}}=2 \pi/r_{\mathrm{min}}$ gives the Fisher errors of similar magnitude and trend to the errors estimated from the simulations.
 While a scaling kmax=1.57/rmin assures that the magnitude and the trends of both the errors are closely matched., While a scaling $k_{\mathrm{max}}=1.5 \pi/r_{\mathrm{min}}$ assures that the magnitude and the trends of both the errors are closely matched.
 We note that a similar scaling kimax© was chosen by tuning in Okumura&Jing(2011) and 7/TminReid&White(2011) to provide a precise link between Fourier and configuration space., We note that a similar scaling $k_{\mathrm{max}}\approx \pi/r_{\mathrm{min}}$ was chosen by tuning in \cite{2011ApJ...726....5O} and \cite{2011MNRAS.417.1913R} to provide a precise link between Fourier and configuration space.
 The constraint on 6 can be used to extract information on the growth function if the galaxy linear bias is measured independently., The constraint on $\beta$ can be used to extract information on the growth function if the galaxy linear bias is measured independently.
" For instance, the amplitude of the matter fluctuations quantified by of’(z=0) and measured from a cosmic microwave background radiation experiment could be first scaled to z=1, and then the bias could be estimated by b=of(z1)/eg(z=1); the superscripts m and g indicating, as usual, matter and galaxies."," For instance, the amplitude of the matter fluctuations quantified by $\sigma_8^{m}(z=0)$ and measured from a cosmic microwave background radiation experiment could be first scaled to $z=1$, and then the bias could be estimated by $b=\sigma_8^{g}(z=1)/\sigma_8^{m}(z=1)$ ; the superscripts $m$ and $g$ indicating, as usual, matter and galaxies."
" This procedure was followed recently, for example, in Guzzoetal.(2008),, where the fluctuation amplitude derived by the (Spergel2007) was utilized."," This procedure was followed recently, for example, in \cite{2008Natur.451..541G}, where the fluctuation amplitude derived by the \citep{2007ApJS..170..377S} was utilized."
" We based our analysis on basic RSD streaming models, frequently used as benchmark models in the literature, and often shown, for some galaxy groups, to work reasonably well on quasi-linear scales."," We based our analysis on basic RSD streaming models, frequently used as benchmark models in the literature, and often shown, for some galaxy groups, to work reasonably well on quasi-linear scales."
 We have introduced a new method to determine the RSD 58 parameter in configuration space on quasi-linear scales., We have introduced a new method to determine the RSD $\beta$ parameter in configuration space on quasi-linear scales.
" The statistics D of ((42)), presents the following advantages: it is not necessary to specify the linear theory predictions for the correlation functions; there is no dependence on the mean galaxy density; only convolutions are involved, assuring stability; it does not rely on galaxy linear bias, apart from the implicit dependence within 8."," The statistics $D$ of \ref{eq:estD}) ), presents the following advantages: it is not necessary to specify the linear theory predictions for the correlation functions; there is no dependence on the mean galaxy density; only convolutions are involved, assuring stability; it does not rely on galaxy linear bias, apart from the implicit dependence within $\beta$."
 Random errors on the determinations of the galaxies redshifts are effectively incorporated in the model., Random errors on the determinations of the galaxies redshifts are effectively incorporated in the model.
" The method can actually be considered more general than the streaming model because it is insensitive to a factor multiplying the power spectrum that depends only on the wavenumber, which could arise from a quasi-linear correction and/or a scale-dependent galaxy bias."," The method can actually be considered more general than the streaming model because it is insensitive to a factor multiplying the power spectrum that depends only on the wavenumber, which could arise from a quasi-linear correction and/or a scale-dependent galaxy bias."
" Furthermore perturbation theory might have better convergence properties in configuration space, where our estimator operates, compared to Fourier space etal.2008;MatsubaraTaruya 2009)..(Sánchez"," Furthermore perturbation theory might have better convergence properties in configuration space, where our estimator operates, compared to Fourier space \citep{2008MNRAS.390.1470S,2008PhRvD..77f3530M,2009PhRvD..80l3503T}."
" In addition, we have noticed a baryonic feature at about 100A-!Mpc in the normalized quadrupole in configuration space, that should not have a major impact on the extraction of 8 with our method."," In addition, we have noticed a baryonic feature at about $100\,h^{-1}\,\mathrm{Mpc}$ in the normalized quadrupole in configuration space, that should not have a major impact on the extraction of $\beta$ with our method."
" We have also carried out analyses based on the N- body simulations of Sato&Matsubara and the Fisher matrix method, finding that errors(2011) of a few percent on f are feasible with a full sky, 1(5-7!Gpc)? survey centered at a redshift of unity and with negligible shot noise and that for minimum separations such that Smin>35h-1Mpc there is no evidence for a bias in the estimation of 5."," We have also carried out analyses based on the $N$ -body simulations of \cite{2011PhRvD..84d3501S} and the Fisher matrix method, finding that errors of a few percent on $\beta$ are feasible with a full sky, $1\,(h^{-1}\,\mathrm{Gpc})^3$ survey centered at a redshift of unity and with negligible shot noise and that for minimum separations such that $s_{\mathrm{min}}>35\,h^{-1}\mathrm{Mpc}$ there is no evidence for a bias in the estimation of $\beta$."
" It is going to be crucial to be able to measure and interpret in detail the growth of fluctuations, because the information contained therein could indicate deviations from general relativity and constrain the cosmic expansion history."," It is going to be crucial to be able to measure and interpret in detail the growth of fluctuations, because the information contained therein could indicate deviations from general relativity and constrain the cosmic expansion history."
" Since subtleties in thesystematics of cosmological observations do not guarantee that the theoretical equivalence of Fourier and configuration space can be easily reproduced in practical observations, it should be convenient to have the widest possible arsenal of methods, involving both points of view."," Since subtleties in thesystematics of cosmological observations do not guarantee that the theoretical equivalence of Fourier and configuration space can be easily reproduced in practical observations, it should be convenient to have the widest possible arsenal of methods, involving both points of view."
path.,path.
" As already indicated before, what is model dependent are the amplitudes of the departure from the cosmic ratios."," As already indicated before, what is model dependent are the amplitudes of the departure from the cosmic ratios."
" Figure 3 shows that the observed points may be reproduced from models having differing velocities and masses, with and without magnetic field; however, by itself this diagram cannot help in disentangling the various possibilities."," Figure \ref{cnomix920} shows that the observed points may be reproduced from models having differing velocities and masses, with and without magnetic field; however, by itself this diagram cannot help in disentangling the various possibilities."
 The same evolutionary tracks are plotted in the N/C vs. logg diagram in Fig. 4.., The same evolutionary tracks are plotted in the $N/C$ vs. $\log g$ diagram in Fig. \ref{cnomix_gnc}.
" We see that, except for the ⊱↥∂↕⊱↕−∐⊃⊖⊖↕⊙⊖⊱∂∐⊂⊔−∐⊃∐∠⇂⊖∠⇂∍⊱≼⊤⊱∁⊙⋝∇∇∐↥∁∤∏⊱∂≣⊜∐−uineslowrotatorwitharatherstrongmagneticfield? of probably fossil2006), the other points may be accounted for by models having low initial rotation."," We see that, except for the stars 61068 and 149438 \citep[$\tau$\,Sco, which is a genuine slow rotator with a
rather strong magnetic field\footnote{A magnetic field is present also
in HD\,74575 \citep{hubrig09}.} of probably fossil, the other points may be accounted for by models having low initial rotation."
 661068 is marginally compatible with the Μο model ∂↥∨≧≖≟↕≖∶∶∍∍⊙⊙↕⊂⇂⋜∐↕⊱⊱−↕∁⊙∐↕↧⊃∐↥⊜⊂↥∇∇↥↥∐∐↕∂≣∐⊜↥↥∁∅∁↕⊂↥⋝∇∨∐↥↕⊜ T SSco challenges present stellar models.," 61068 is marginally compatible with the $M_\odot$ model at $v_\mathrm{rot}^\mathrm{ini}$ $^{-1}$ computed with magnetic field, while $\tau$ Sco challenges present stellar models."
" It has a behaviour that may be explained by a homogeneous evolution, but this still has to be confirmed by further computations."," It has a behaviour that may be explained by a homogeneous evolution, but this still has to be confirmed by further computations."
 This is a very interesting star that certainly deserves further inspections both from the observational and theoretical points of view., This is a very interesting star that certainly deserves further inspections both from the observational and theoretical points of view.
 The behaviour of light element abundances in the whole star sample is shown in Fig. 5.., The behaviour of light element abundances in the whole star sample is shown in Fig. \ref{ourcnomix}.
" In contrast to the literature values (Fig. 1)),"," In contrast to the literature values (Fig. \ref{litsummary}) ),"
" a clear and tight trend is found, confirming the predicted locus of N/O-N/C abundance ratios."," a clear and tight trend is found, confirming the predicted locus of $N/O$ $N/C$ abundance ratios."
" However, as already indicated above, the MM03 models for rotating stars with mass loss evolving towards the red supergiant stage (solid line in Fig. 5))"," However, as already indicated above, the MM03 models for rotating stars with mass loss evolving towards the red supergiant stage (solid line in Fig. \ref{ourcnomix}) )"
" predict mixing that is too low, i.e. too low f (Sect. 2)),"," predict mixing that is too low, i.e. too low $f$ (Sect. \ref{theory}) ),"
 in particular for most of the supergiants., in particular for most of the supergiants.
 A, A
observed: near-infrarecl colors is discrepant. with the mass inferred from the companion's Luminosity and the age of the primary (Mohantyοἱal.2007).,observed near-infrared colors is discrepant with the mass inferred from the companion's luminosity and the age of the primary \citep{Moh07}.
. Discrepancies such as these may arise because voung exoplanets exist in a gravitv-elfective temperature (g.Liar) regime in which both the evolutionary ancl atmospheric models have vet to be validated.," Discrepancies such as these may arise because young exoplanets exist in a gravity-effective temperature $g,T_{\rm eff}$ ) regime in which both the evolutionary and atmospheric models have yet to be validated."
 Fits of photometry and spectroscopy to predictions of atmosphere mocels depend upon the veracity. of the models. themselves. which—in the Tir range of interestin turn sensitively depend. upon model cloud. profiles. which are as vet uncertain.," Fits of photometry and spectroscopy to predictions of atmosphere models depend upon the veracity of the models themselves, which—in the $T_{\rm eff}$ range of interest—in turn sensitively depend upon model cloud profiles, which are as yet uncertain."
 Extensive experience with fitting models to brown dwarf spectra and photometry (Cushingctal.2008:Stephenset2009) reveals that while effective temperature can be fairly tightly constrained. gravity determinations are usually less precise. uncertain in some cases by almost an order of magnitude in ο.," Extensive experience with fitting models to brown dwarf spectra and photometry \citep{Cus08, Ste09} reveals that while effective temperature can be fairly tightly constrained, gravity determinations are usually less precise, uncertain in some cases by almost an order of magnitude in $g$."
 While there are low gravity spectral indicators recognized [rom surveys of voung objects (Cruzetal.2009:Kirkpatricketal.2006). these have vet to be calibrated by studies of binary objects which allow independent. measures of mass.," While there are low gravity spectral indicators recognized from surveys of young objects \citep{cruz09, 
kirk06}
 these have yet to be calibrated by studies of binary objects which allow independent measures of mass."
" Ideally for a single object with a given radius 2. evolution moclel luminosity. which (for a known parallax) constrains BT, would be fully consistent with (g. constraints from atmosphere model fitting."," Ideally for a single object with a given radius $R$, evolution model luminosity, which (for a known parallax) constrains $R^2T_{\rm eff}^4$ would be fully consistent with $(g, T_{\rm eff})$ constraints from atmosphere model fitting."
 But. as noted τι).above. this is often not in fact the case. as the derived. luminosity. mass. and radii of the companion to 2MI207 b as well as the LR 8799 planets are not fully internally self-consistent with standard evolution mocels.," But, as noted above, this is often not in fact the case, as the derived luminosity, mass, and radii of the companion to 2M1207 b as well as the HR 8799 planets are not fully internally self-consistent with standard evolution models."
 Given the likely future ubiquity of direct. detections of voung. hot Jupiters and the clear need. for actelitiona independent methods to constrain planet properties. we have explored. the utility of polarization as an acdcditiona method for characterizing scl-luminous planets.," Given the likely future ubiquity of direct detections of young, hot Jupiters and the clear need for additional independent methods to constrain planet properties, we have explored the utility of polarization as an additional method for characterizing self-luminous planets."
 Polarization of cosc-in giant exoplanets whose ho atmosphere favours the presence of silicate condensates. is discussed by Seager.Wütnev&Sassclov(2000) and by Sengupta&Maiti(2006).," 	 Polarization of close-in giant exoplanets whose hot atmosphere favours the presence of silicate condensates, is discussed by \cite{Sea00} and by \cite{sengupta}."
. While these authors considere the polarization of the combined light from an unresolvec system of star and planet. Stam.Hovenier&Waters(2004) presented. the polarization. of the rellected. light. of a resolved. clirecthy-imaged Jupiter-like exoplanet.," While these authors considered the polarization of the combined light from an unresolved system of star and planet, \cite{Sta04} presented the polarization of the reflected light of a resolved, directly-imaged Jupiter-like exoplanet."
 Since the polarized light of a close-in exoplanet is combined with the unpolarized continuum Dux of the star which cannot. be resolved. the amount ofobservable polarization in such case is extrmeley low — of the order of magnitude of planct-to-star Hux ratio.," Since the polarized light of a close-in exoplanet is combined with the unpolarized continuum flux of the star which cannot be resolved, the amount of observable polarization in such case is extrmeley low – of the order of magnitude of planet-to-star flux ratio."
 Polarization measurements of. directlv-imaged exoplanets in reflected light is also challenging., Polarization measurements of directly-imaged exoplanets in reflected light is also challenging.
 The removal of scattered. light. from the primary star must. be precise in both polarization channels so that the planets intrinsic polarization. (which is a dillerential. measurement) can be accurately determined., The removal of scattered light from the primary star must be precise in both polarization channels so that the planet's intrinsic polarization (which is a differential measurement) can be accurately determined.
 In any case no extrasolar planet has vet been imaged in scattered light. an accomplishment that will likely require a space-based coronagraph (e.g...Boc-calettictal. 2011).," In any case no extrasolar planet has yet been imaged in scattered light, an accomplishment that will likely require a space-based coronagraph \citep[e.g.,][]{Boc11}."
. Measuring polarization of thermally emitted radiationas we propose hereis also cliflieult but does not require a planet to be close to the star (where the starlight suppression is most difficult) so that it is bright in reflected light., Measuring polarization of thermally emitted radiation—as we propose here—is also difficult but does not require a planet to be close to the star (where the starlight suppression is most difficult) so that it is bright in reflected light.
 Furthermore extrasolar planets have already been imaged which raises the possibility. of polarization observations., Furthermore extrasolar planets have already been imaged which raises the possibility of polarization observations.
 lt is clear from comparisons of model spectra to data that most of the exoplanets directly. imaged to date have dusty atmospheres (Maroisetal.2008:Bowler2010:2011:Skemerctal.2011).," It is clear from comparisons of model spectra to data that most of the exoplanets directly imaged to date have dusty atmospheres \citep{Mar08,Bow10,Laf10,Bar11,Cur11,Ske11}."
. Clear atmospheres lacking dust grains can be polarized. but only at blue optical wavelengths where gaseous Ravleigh scattering is important. (Sengupta&Marley.2009).," Clear atmospheres lacking dust grains can be polarized, but only at blue optical wavelengths where gaseous Rayleigh scattering is important \citep{Sen09}."
. Since even the hottest voung exoplanets will not emit significantly in the blue. grain scattering must be present for there to be measurable polarization in the near-infrared where warm giant planets are bright (Sengupta&Marley.2010).," Since even the hottest young exoplanets will not emit significantly in the blue, grain scattering must be present for there to be measurable polarization in the near-infrared where warm giant planets are bright \citep{Sen10}."
.. Phere are two temperature ranges within which we expect a gas giant exoplanet to possess significant atmospheric condensates., There are two temperature ranges within which we expect a gas giant exoplanet to possess significant atmospheric condensates.
 The first is L-dwarf like planets (roughly 1000.«Zi24001) where iron and silicate erains condense in the observable atmosphere.," The first is L-dwarf like planets (roughly $1000 < 
T_{\rm eff} < 2400\,\rm K$ ) where iron and silicate grains condense in the observable atmosphere."
 “Phe lower end. of this range in the planetary mass regime is as vet uncertain., The lower end of this range in the planetary mass regime is as yet uncertain.
 The second temperature range occurs in cool planets with atmospheric water clouds (Zar<OI).," The second temperature range occurs in cool planets with atmospheric water clouds $T_{\rm eff} < 400\,\rm K$ )."
 There have vet been no confirmed detections of such planets., There have yet been no confirmed detections of such planets.
 Here we will focus on the first category. since such objects are brighter. more casily detectable. and the comparison to the ficlcl chwarls is possible.," Here we will focus on the first category since such objects are brighter, more easily detectable, and the comparison to the field dwarfs is possible."
 Although survey sizes are fairly small. linear polarization of field. L dwarfs has been detected.," Although survey sizes are fairly small, linear polarization of field L dwarfs has been detected."
 Menardetal.(2002). and ZapateroOsorioetal.(2005) both report that a fraction of L. dwarfs. particularly the later. dustier spectral types. are intrinsically polarized.," \cite{Men02} and \cite{Oso05} both report that a fraction of L dwarfs, particularly the later, dustier spectral types, are intrinsically polarized."
 Sengupta&Marley(2010). find that the observed. polarization. can plausibly arise from. emission. of cloudy. oblate cdwarls. although to produce the required oblateness or more) the chwarls must have fairly low eravity for a field dwarf (g300ms 7) and rapid rotation.," \cite{Sen10} find that the observed polarization can plausibly arise from emission of cloudy, oblate dwarfs, although to produce the required oblateness or more) the dwarfs must have fairly low gravity for a field dwarf $g\sim 300\,\rm m\,s^{-2}$ ) and rapid rotation."
 The required rotation periods are brisk. as little as 2 hours or less. but are Compatible with observed: rotational velocities in at east some cases (see Sengupta&Alarles(2010). for a discussion).," The required rotation periods are brisk, as little as 2 hours or less, but are compatible with observed rotational velocities in at least some cases (see \cite{Sen10} for a discussion)."
 Sengupta&Marley(2010) further find. that he near-infrared polarization is greatest at Zip16001x where their model condensate couds are both optically hick and still. prominent enough. in the. photosphere to maximally alfect the polarization.," \cite{Sen10} further find that the near-infrared polarization is greatest at $T_{\rm eff} \sim 1600\,\rm K$ where their model condensate clouds are both optically thick and still prominent enough in the photosphere to maximally affect the polarization."
 Surface inhomogeneities can also give rise to à ne x»larization (Menard.&Delfosse2004) απ experience rom the solar svstem confirms that irregularly spaced clouds are to be expected., Surface inhomogeneities can also give rise to a net polarization \citep{Men04} and experience from the solar system confirms that irregularly spaced clouds are to be expected.
 Both Jupiters and Saturn's therma emission in the five-micron spectral window is strongly modulated by horizontally inhomogeneous cloud. cover ane it would not be surprising to find similar morphology in the atmospheres of exoplanets., Both Jupiter's and Saturn's thermal emission in the five-micron spectral window is strongly modulated by horizontally inhomogeneous cloud cover and it would not be surprising to find similar morphology in the atmospheres of exoplanets.
 La the presence. of surface inhomogeneity. the asymmetry that produces. the net non-zero clisk-integrated polarization would. increase and. hence a combination of oblate photosphere ane surface inhomogeneity can give rise to detectable levels of »olawization.," In the presence of surface inhomogeneity, the asymmetry that produces the net non-zero disk-integrated polarization would increase and hence a combination of oblate photosphere and surface inhomogeneity can give rise to detectable levels of polarization."
 Exoplanets are even better candidates than L dwarfs to jwe an oblate shape and be polarized., Exoplanets are even better candidates than L dwarfs to have an oblate shape and be polarized.
 With a lower mass and roughly the same radius as a brown cwarf (and thus a ower eravitv). a rapidly rotating planet can be significantly oblate and consequently. produce. a polarization signal even without surface inhomogencities.," With a lower mass and roughly the same radius as a brown dwarf (and thus a lower gravity), a rapidly rotating planet can be significantly oblate and consequently produce a polarization signal even without surface inhomogeneities."
 Llere we explore he conditions under which the thermal emission from a warm. voung exoplanet may be polarized ancl consicler the scientific value. of measuring exoplanet polarization.," Here we explore the conditions under which the thermal emission from a warm, young exoplanet may be polarized and consider the scientific value of measuring exoplanet polarization."
 We irst look at the issue of oblateness. then present a set of," We first look at the issue of oblateness, then present a set of"
insullicient producers of Ii with respect to observations (Vinimes et al.,insufficient producers of $^{47}$ Ti with respect to observations (Timmes et al.
 1995)., 1995).
 The production of Pi is traced to the nueleosvnthesis of Cr in stellar explosions: two subsequent J-decays alter its ejection leads to “Ti., The production of $^{48}$ Ti is traced to the nucleosynthesis of $^{48}$ Cr in stellar explosions; two subsequent $\beta$ -decays after its ejection leads to $^{48}$ Ti.
 Vhis occurs mostly in explosive silicon-burning ancl curing helium. fusion., This occurs mostly in explosive silicon-burning and during helium fusion.
 The latter could occur either by the a-rich [reezeout. of shock-decomposed nuclei near the core of a Type LE supernova. or as part of explosive helium burning associated with Type la supernovae.," The latter could occur either by the $\alpha$ -rich freezeout of shock-decomposed nuclei near the core of a Type II supernova, or as part of explosive helium burning associated with Type Ia supernovae."
 TN ↓↓↓↓⊳∖↓≻↓⋅⋯⊔∐⇍∢⊾∠⇂⊔↓⋜⊔⊔↓∙∖⇁∣⋡∙∖⇁↥↓↥⋖⋅⊔⋯⇍↓⋖⋅∪≱∖∙∖⇁⊔↥⇂↥⋖⋅≱∖↓≱∖∪⇂. . . 1054 . ↓⋅⋯⊔∪⋯∼⊔∖⇁⋖⊾≺↓⋅↓⊔⊳∖↿⋖⋅∐⋜⊔⋅∢⋅⇀∖↓≻⇂∪≱∖↓∪⊔≱∖∪⇂∣⋯↿⇂↥⇂∙∖⇁↓≻∢⋅≱∖∪⇂ . supernovae., $^{49}$ Ti is produced mainly by the nucleosynthesis of radioactive $^{49}$ Cr in stellar explosions of both types of supernovae.
 rpsThe isotope. I Cr Opusis the result of the explosive. fusion⋅⋠ of pashelium as outlined. above: Cr opo:is also synthesised. during silicon-burning. as for the lighter titanium isotopes.," The isotope $^{49}$ Cr is the result of the explosive fusion of helium as outlined above; $^{49}$ Cr is also synthesised during silicon-burning, as for the lighter titanium isotopes."
" It has been suggested that ""Ei is produced. primarily in sub-Chandrasekhar mass Type Ia supernovae (Timmes et al.", It has been suggested that $^{50}$ Ti is produced primarily in sub-Chandrasekhar mass Type Ia supernovae (Timmes et al.
 1995). during which electron capture turns the composition neutron-rich.," 1995), during which electron capture turns the composition neutron-rich."
" Some fraction of the ""Ti is likely also made by slow neutron capture within the the burning shells of pre-supernovae massive stars anc ACB stars.", Some fraction of the $^{50}$ Ti is likely also made by slow neutron capture within the the burning shells of pre-supernovae massive stars and AGB stars.
" Le would appear that explosive. burning in Type Hb supernovae does not produce significant quantities of ""Ti.", It would appear that explosive burning in Type II supernovae does not produce significant quantities of $^{50}$ Ti.
" '""Fhroughout our work. we have used the (Fenner Gibson 2003: Fenner. Murphy Gibson 2005) ealactic chemical evolution package. emploving its default ~clual infall (halo | disc) mode (similar in spirit to the seminal models of Chiappini. Matteuccei Gratton 1997)."," Throughout our work, we have used the (Fenner Gibson 2003; Fenner, Murphy Gibson 2005) galactic chemical evolution package, employing its default “dual infall” (halo + disc) mode (similar in spirit to the seminal models of Chiappini, Matteucci Gratton 1997)."
 Within this framework. the halo phase occurs on a rapid timescale and cnriches the initially primordial gas to a metallicity of ~10% solar.," Within this framework, the halo phase occurs on a rapid timescale and enriches the initially primordial gas to a metallicity of $\sim$ solar."
 Phe second. (disc) phase is delaved by ~ 1 Gyr with respect to the first. ancl actsover a more prolonged. timescale.," The second (disc) phase is delayed by $\sim$ 1 Gyr with respect to the first, and actsover a more prolonged timescale."
 We assume the infall of fresh material during this second phase to be somewhat metal-enriched solar) and a-enhanced. (0.4 dex). consistent with patterns seen in metal-poor halo/thick disc stars (c.g. livan. Norris Beers 1996) and. present-cay high-velocity infalline halo gas (e.g. Gibson οἱ al.," We assume the infall of fresh material during this second phase to be somewhat metal-enriched solar) and $\alpha$ -enhanced (0.4 dex), consistent with patterns seen in metal-poor halo/thick disc stars (e.g. Ryan, Norris Beers 1996) and present-day high-velocity infalling halo gas (e.g. Gibson et al."
 2001)., 2001).
 The rate at which material is accreted. is assumed to decline exponentially., The rate at which material is accreted is assumed to decline exponentially.
 Phe evolution of total surface mass density σι(Ε.Ε) is given by: where the infall rate cocllictents ο) and Bir) are chosen in order to reproduce the present-day surface mass density of the halo and disk components. which we take to be 10 and 45 M. pe7. respectively.," The evolution of total surface mass density $\sigma_{tot}(r, t)$ is given by: where the infall rate coefficients $A(r)$ and $B(r)$ are chosen in order to reproduce the present-day surface mass density of the halo and disk components, which we take to be 10 and 45 $M_{\odot}$ $^{-2}$, respectively."
 The adopted timescales for the infall phases in the solar neighbourhood (r. =s.5kpe) are Ty= 0.05 Gye and 7p= 10.5 Gyr. with the functional form for the latter being το(11 δι127. rellecting the “inside-out” formation framework in which the dual-infall model operates.," The adopted timescales for the infall phases in the solar neighbourhood $r_{\odot}$ =8.5kpc) are $\tau_H =$ 0.05 Gyr and $\tau_D =$ 10.5 Gyr, with the functional form for the latter being $\tau_D(r)$ $1.38r-1.27$, reflecting the “inside-out” formation framework in which the dual-infall model operates."
 These timescales and cocllicicnts provide model predictions Consistent with various local observational constraints such as the metallicity cistribution function. age metallicity relation. ancl present-clayv gas surface density distribution. ancl are consistent with those eniplovecl by Fenner et al. (," These timescales and coefficients provide model predictions consistent with various local observational constraints such as the metallicity distribution function, age metallicity relation, and present-day gas surface density distribution, and are consistent with those employed by Fenner et al. ("
2005).,2005).
" We adopt oa fairly conservative star formation prescription based upon a ""Schmidt Law” of the form: where the value of the star formation elliciencyv. 7 is constrained by the present-day gas fraction (for this work. Civr 1 v=0.06 "," We adopt a fairly conservative star formation prescription based upon a “Schmidt Law” of the form: where the value of the star formation efficiency $\nu$ is constrained by the present-day gas fraction (for this work, $\nu$$=$ 0.06 $^{-1}$ )."
The shape of the initial mass function (AIL) controls the raction of material locked-up in stellar generation. which in turn determines the rate at whieh dilferent. elements are released into the ISM.," The shape of the initial mass function (IMF) controls the fraction of material locked-up in stellar generation, which in turn determines the rate at which different elements are released into the ISM."
 Our default assumption is that of the hree-component ΕΛΣ of Ixroupa. Tout. Cilmore (1993: jereafter. WLC). with lower- and upper-mass limits of 0.08 and GO AL. respectivelv: the WiC LME Lies between those of Salpeter (1955) ancl Sealo (1986). in terms of mass fraction ied up in Pvpe LE supernovae progenitors (e.g. Table 7 ol Gibson 1997).," Our default assumption is that of the three-component IMF of Kroupa, Tout Gilmore (1993; hereafter, KTG), with lower- and upper-mass limits of 0.08 and 60 $_{\odot}$, respectively; the KTG IMF lies between those of Salpeter (1955) and Scalo (1986), in terms of mass fraction tied up in Type II supernovae progenitors (e.g. Table 7 of Gibson 1997)."
 Unless otherwise stated. we assume that he mass fraction of the IME which is tied up in SNe Ia wogenitor binary systems (total binary masses in the range 16 M.) isdz such an assumption vields. within our aclopted model formalism. a clisc-averaged present-day ratio tween. SNe Ll and SNe la rates of 4.1 (consistent. with he cisc-averagecl empirical SNe LE to SNe la ratio of 3.7 van den Bereh LOSS)," Unless otherwise stated, we assume that the mass fraction of the IMF which is tied up in SNe Ia progenitor binary systems (total binary masses in the range $-$ 16 $_\odot$ ) is; such an assumption yields, within our adopted model formalism, a disc-averaged present-day ratio between SNe II and SNe Ia rates of 4.1 (consistent with the disc-averaged empirical SNe II to SNe Ia ratio of 3.7 – van den Bergh 1988)."
 We supplement this by exploring a range of single power-law LME slopes to isolate the relative contributions of low- ancl high-mass stars., We supplement this by exploring a range of single power-law IMF slopes to isolate the relative contributions of low- and high-mass stars.
 In order to sample the range of uncertainties inherent. to stellar evolution modelling. we explore the use of several sets of metallicitv-cdependent nucleosvnthetie vields. in this work those of Woosley Weaver (1995). Chielli Limongi (2004). IXobavashi et al for Ένρο LE and those of Warakas Lattanzio (2007). for low- and intermediate-mass single stars (hereafter. NN95. CLOA. Ix06. ancl IXNLOT. respectively).," In order to sample the range of uncertainties inherent to stellar evolution modelling, we explore the use of several sets of metallicity-dependent nucleosynthetic yields, in this work – those of Woosley Weaver (1995), Chieffi Limongi (2004), Kobayashi et al for Type II and those of Karakas Lattanzio (2007), for low- and intermediate-mass single stars (hereafter, WW95, CL04, K06, and KL07, respectively)."
 For Type Ia supernovae. the vields of Nomoto et al. (," For Type Ia supernovae, the yields of Nomoto et al. ("
1997) have been assumed.,1997) have been assumed.
 As these vields only have data for M < 40 M. and," As these yields only have data for M $\leq$ 40 $_{\odot}$ , and"
Finding the response of a general steady state uaguetospheric accretion flow to an orbiting protoplauet a in.verv difficultο”. problem.,Finding the response of a general steady state magnetospheric accretion flow to an orbiting protoplanet is a very difficult problem.
scr. Accordingly⋅∖ we⊳⊳⋅⋅⊾⇁ considerτο q/.cle? )ossible simplificatio-sS, Accordingly we consider possible simplifications.
 In this section we consider the calculation of the response using a local approximation., In this section we consider the calculation of the response using a local approximation.
 The idea behind this is that at any stage. the protoplauet noves relative to field lines along which the accretion flow noves. such that the speed of the protoplauct relative o the gas. denoted by |Uv| below. may become comparable to the orbital speed (see Figure 1).," The idea behind this is that at any stage, the protoplanet moves relative to field lines along which the accretion flow moves, such that the speed of the protoplanet relative to the gas, denoted by $|{\bf U} - {\bf v}|$ below, may become comparable to the orbital speed (see Figure 1)."
 The characteristic scale of interaction at which a protoplauct of. mass. Mj. would dominate: the gas flow would then ↴⋝↸∖↴∖↴↸∖⊓⋝∙↖⇁↑∐↸∖↸⊳↕⋜↧↴∖↴↴∖↴↕↸⊳⋜↧↕↕≧∪∐≼∐≓∐∪⋅↖⇁↕↸∖↥⋅⋜∥∐∏↴∖↴⋖↸∖∙∶↴∙⊾∙∐∪⋅↖⇁↕↸∖ ∙∖↽∫⇀⋅↖⇁↑↑↕↸∖↑∪∐↕∩∶≩∩∶↕⊰∏∐∎↸∖↥⋅↑↕∩≝⋊∔⋝∣⋮∠≩∐∶⊇∊↙⋮⊽⋀∐∣↗↴⋃ ⇉∙↖↖⊽↕∐∖∐↑∐↸∖↥⋅↸∖↕⋜↧⊓↖↽↸∖↴∖↴⋉∖↸∖≼⊔↴∖↴↸⊳∐⋜∐⋅⋯⊳↑↸∖↥⋅↕↴∖↴↑↕↸⊳⋜↧∐∙↖↽∪↥⋅↴⋝↕↑⋜↧↕∙ rpy is smaller by a factor ~MQ/M.1 than the radius at which the protoplauet orbits indicating a local luteraction on a characteristic time scale that is smaller than the orbital time by a similar factor.," The characteristic scale of interaction at which a protoplanet of mass, $M_p,$ would dominate the gas flow would then be set by the classical Bondi-Hoyle radius (e.g. Hoyle Lyttleton 1939; Ruffert 1996) $r_{BH} = 2GM_p/|{\bf U} - {\bf v}|^2.$ When the relative speed is characteristically orbital, $r_{BH}$ is smaller by a factor $ \sim M_p/M_* \ll 1$ than the radius at which the protoplanet orbits indicating a local interaction on a characteristic time scale that is smaller than the orbital time by a similar factor."
" The siualluess of rpuírgu~(CMM.Pe with ry being the Till radius also indicates that the interaction of the protoplauct with eas that is not bound to it is linear even iu the elaut plauct regime with A,M.~102. The margius involved are such that the interaction is expected to be local aud linear even when the relative speed is siguiBcantlv less than orbital."," The smallness of $r_{BH}/r_H \sim (M_p/M_*)^{2/3},$ with $r_H$ being the Hill radius also indicates that the interaction of the protoplanet with gas that is not bound to it is linear even in the giant planet regime with $M_p/M_* \sim 10^{-3}.$ The margins involved are such that the interaction is expected to be local and linear even when the relative speed is significantly less than orbital."
" Thislus isd unlike""↘ the1 situation.‘ im. protoplauctaryps1t disksid wherevhere theο μαιsmall relative‘clative velocitsvelocity. betweentween gasσας andai ]xotoplanetps can‘ makeUS the interaction‘ nonlinear ttat scales‘ larger‘ than' the Wall radius‘ resulting in mσα formation‘ ""(see c.g.o ""liLin NonaPavaluiPapaloizou 1993: maysDPapaloizouloi TTerquem 2006: and references therein for moreiore discusion)discussion).", This is unlike the situation in protoplanetary disks where the small relative velocity between gas and protoplanet can make the interaction nonlinear at scales larger than the Hill radius resulting in gap formation (see e.g. Lin Papaloizou 1993; Papaloizou Terquem 2006; and references therein for more discussion).
 Qn↽ the basis⋅ of. the above discussion.. we here adopt the assumption. that the protoplanet maguetosphere iuteraction. ⋅⋅is localised⋅ ium⋅ the vicinity⋅⋅⋅ of. the protoplanct ‘aud that‘ the MEMassociated characteristicUT time scaleu is| sieuificautlv shorter than the orbital period.," On the basis of the above discussion, we here adopt the assumption that the protoplanet magnetosphere interaction is localised in the vicinity of the protoplanet and that the associated characteristic time scale is significantly shorter than the orbital period."
 Accordingly we adopt a local frame. that can be taken to be a non rotating local Cartesian: coordinate: system with: origiun instantaneously at the centre of mass of the protoplanct and moving with a coustaut velocity equal to that of the rotational velocity the star would have if it extended to the ↥⋅⋜∥↕⋜↧↕↕∪↸⊳⋜↧↑↕∪∐∪↕≯↑∐↸∖↻↥⋅∪↑∪↻↕⋜∐∐∖↑⋖↕⋟∪∐↸∖∏↴∖↴↸∖↴∖↴⋜↧↕∪↸⊳⋜↧⊔≯↥⋅⋜∐⊔↸∖ Oo Oo ∐⊼↸∖≺⊔∐↑∐∖∪↥⋅↕∩⊾↕∐⋜↧↕↥⋅∪↑⋜↧↑↕↕∩⊾↕⋟↥⋅⋜∐⊔↸∖⋜⋯≼↧↑↕∐∖∐↕∐⋯⊾↕↸∖↸⊳↑↴∖↴≤≥Oo in the equations of motion the same results are obtained)., Accordingly we adopt a local frame that can be taken to be a non rotating local Cartesian coordinate system with origin instantaneously at the centre of mass of the protoplanet and moving with a constant velocity equal to that of the rotational velocity the star would have if it extended to the radial location of the protoplanet (if one uses a local frame fixed in the original rotating frame and then neglects $\Omega_*$ in the equations of motion the same results are obtained).
 The respouse to the protoplauet is readily calculable ⋅ ⋅ ∪↕⋡↑∐↸∖⋉∖↥⋅⊓∐⋅↴⋝⋪↕↑↕∪∐↴∖↴⋅↽∕∏⋯↴∖↴↕≯∪↥⋅⋪↧∩⊾↸∖∐↸∖↥⋅⋪↕↕⋉∖↥⋅⊓∐⋅↴⋝⋪↧↑↕∪∐ ⋜⋯≼↧↕↸∖⋜∥↧↴∖↴↑∪⋜↧↸⊳∪∐∏↥⋅⋯⋜↧⊓∪∐∪↕↑∐↸∖↸⊳∪∐↸⊳↕∏↴∖↴↕∪∐↑∐⋜↧↑↑∐∖ changes ⋝⋅induced iu⋅ the orbit⋅ are associated∙ with∙ wave cherey fluxes propagating away from the protoplanct., The response to the protoplanet is readily calculable and leads to a confirmation of the conclusion that the changes induced in the orbit are associated with wave energy fluxes propagating away from the protoplanet.
 Furthermore when such waves are absent. such as when the relative∙ speed is∙ below all possible∙ wave speeds. there 1s no cherev change induced on the orbit.," Furthermore when such waves are absent, such as when the relative speed is below all possible wave speeds, there is no energy change induced on the orbit."
" The effect of a perturbing protoplanet of mass M, at position. vector r —ry, ds to produce an Enlerian perturbation to tbe eravitational potential =CAL/σεVeΠΩfywhereΤο bids a : softening :.is paramcter which can be specified to ensure ® remains ήτο,"," The effect of a perturbing protoplanet of mass $M_p$ at position vector ${\bf r}={\bf r}_{orb}$ is to produce an Eulerian perturbation to the gravitational potential $\Phi' = -GM_p/\sqrt{|{\bf r} -{\bf r}_{orb}|^2+b^2},$ where $b$ is a softening parameter which can be specified to ensure $\Phi'$ remains finite."
 An appropriate value is expected to be of the order of the radius of the plauct., An appropriate value is expected to be of the order of the radius of the planet.
 Iu the original rotating frame. Για Lagrangian displacement induced by by this perturbing poteutial satisfies the equation (se |y.," In the original rotating frame, the Lagrangian displacement ${\mbox{\boldmath$ $}}$ induced by by this perturbing potential satisfies the equation (see eg."
" üt ∙ ⋜⋯≺↧↑↕∐∖∫⇀⋜↧∶↴∙⊾↥⋅⋜⋯∶↴⋁↕⋜⋯↖↽⋜∐⋅↕⋜↧↑↕∪∐∪↕↑∐↸∖↕∪↥⋅↸⊳↸∖↻↸∖↥⋅⋯∐↑⋯⋜↧↴∖↴↴∖↴ ↕↴∖↴∫−↰⋖∊⋟∶↕∎∣↖⋚∙∇↕∎∙↖↖↽↕↑∐ Vb""."," Frieman Rotenberg 1960; Lynden-Bell Ostriker 1967) + _* = where the convective derivative operator + and the Lagrangian variation of the force per unit mass is $F(\boldxi) = {\bf f}' + \boldxi\cdot \nabla {\bf f},$ with."
 The Eulerian perturbationslip to the lipinaguetfie feld aud density are even by B — aud p= edot(p£)).. respectively., The Eulerian perturbations to the magnetic field and density are given by ' = and = ) respectively.
 For au isothermal equation of state the Eulerian. pressure perturbationB P!=pc.Di EquationsH (7.— - 44 113) thus provide: a complete svstem for: deteriuinimg⋅⋅↜ the Lagrangian.3 displacementR £.- We- inake the usual local approximation: that the steady state variables: p. aud P as well as v and B cau be aken to be coustaut.," For an isothermal equation of state the Eulerian pressure perturbation $P'=\rho' c^2.$ Equations \ref{Lagp} - \ref{indp}) ) thus provide a complete system for determining the Lagrangian displacement $\boldxi.$ We make the usual local approximation that the steady state variables $\rho,$ $\Phi$ and $P$ as well as ${\bf v}$ and ${\bf B}$ can be taken to be constant."
 In the local frame. ( see Figure: ↽1) he velocity: of:theprotoplanet is: taken to be U and we nay use equations (7. - 11)) with Q.=0.., In the local frame ( see Figure 1) the velocity oftheprotoplanet is taken to be ${\bf U}$ and we may use equations \ref{Lagp} - \ref{indp}) ) with $\Omega_*=0.$.
 Equation (7)) 1ion °gives. —|v-V με⋅ =LOU pewhich. oscther with equations (10)). (11)) aud P4=ο enable the Lagrangian. displacement. induced. bv the rturbiug potential 9 to be caleulated.," Equation \ref{Lagp}) ) then gives + )^2 = which together with equations \ref{PF}) ), \ref{indp}) ) and $P' = -\rho' c^2$ enable the Lagrangian displacement induced by the perturbing potential $\Phi'$ to be calculated."
 . ⊟≻↥⋅↑∐↕↴∖↴↕⋯⊳⋜↧↕⋯∪≼∐∖↕∙↖↖⇁↸∖↴∖↴↸∖↑↑∐↸∖↴∖↴∪⇈↸∖⋯∐∶↴∙⊾↻⋜∐⋅⋜∐⊔↸∖↑↸∖↥⋅∣↗∶∩∙ ↴∖↴∪↑∐⋜↧↑↕∐↑∐↸∖↕≺⋈⊳⋜↧↕↸⊳∪∪↥⋅≼∐∐⋜↧↸∖↴∖↴∙↑∐↸∖↻↸∖↥⋅⊓∐⋅↴⋝↕∐∶↴∙⊾↻∪↑↸∖∐↑↕⋜↧↕ ⋯⋜↧∙↖↽↴⋝↸∖↖↖↽↥⋅↕⇈↸∖∐≼↕≻∣∶≬↙∣⊽⋀∐∣↗↴↥⋅⋃⊔∙↖↖↽↕↑∐↥⋅≼∐∖∐∪⊓∐∶↴⋁↑∐↸∖," For this local model, we set the softening parameter $b=0,$ so that in the local coordinates, the perturbing potential may be written $\Phi' = -GM_p/|{\bf r} - {\bf U}t|,$ with ${\bf r}$ denoting the position vector."
 ↻∪↴∖↴↕↑↕∪∐↖↽↸∖↸⊳↑∪↥⋅∙⊺∪↴∖↴∪↕↖↽↸∖≺⊔⊔↖↖⇁↸∖↑⋜∐↘↽↸∖⊟∏∐⋅↕↸∖↥⋅⊓⋅⋜⋯↴∖↴↕⋟∪↥⋅⊔↴∖↴, To solve \ref{Lagploc}) ) we take Fourier transforms of the perturbations.
" ⋅↴↽‘ MES a]‘ ≺∣∏⋜⋯↑↕↑↖↽∣≥∙ Q= Qtek, with the Fourier transform | r))Qtr)d""r"," Thus for a general perturbation quantity $Q,$ Q= ) d^3 with the Fourier transform )Q ) d^3"
" ⋅↴↽‘ MES a]‘ ≺∣∏⋜⋯↑↕↑↖↽∣≥∙ Q= Qtek, with the Fourier transform | r))Qtr)d""r."," Thus for a general perturbation quantity $Q,$ Q= ) d^3 with the Fourier transform )Q ) d^3"
" ⋅↴↽‘ MES a]‘ ≺∣∏⋜⋯↑↕↑↖↽∣≥∙ Q= Qtek, with the Fourier transform | r))Qtr)d""r.5"," Thus for a general perturbation quantity $Q,$ Q= ) d^3 with the Fourier transform )Q ) d^3"
as follows: This defines a swath through paraicter space with the sae shape as the dotted line in Figure 1.,as follows: This defines a swath through parameter space with the same shape as the dotted line in Figure 1.
 For the three vemainine planets (IID 119026b. IID LS9733h. WD 209158b). phase variation measurements help break the degencracy: We create a two-dimensional normalized probability distribution function. (PDF) for cach planet. then add these together to create the global PDF shown in Fiere L.," For the three remaining planets (HD 149026b, HD 189733b, HD 209458b), phase variation measurements help break the degeneracy: We create a two-dimensional normalized probability distribution function (PDF) for each planet, then add these together to create the global PDF shown in Figure \ref{pdf_albedo_epsilon}."
 This is a democratic wav of represcuting the data. since each planct’s distribution coutributes equally.," This is a democratic way of representing the data, since each planet's distribution contributes equally."
 m Figures 5 and 6 we show the distribution functions for he albedo and circulation of the 21 planets in our sale. obtained by imareinalizine the elobal PDF of Fiere | over either Ap or 5 ," In Figures \ref{all_R50_N0_both_circulation_hist} and \ref{all_R50_N0_both_A_hist} we show the distribution functions for the albedo and circulation of the 24 planets in our sample, obtained by marginalizing the global PDF of Figure \ref{pdf_albedo_epsilon} over either $A_{B}$ or $\varepsilon$ ."
Tie solid line in Figure 5.Γ shows no evidence of binnodality in heat redistribution cficiency. although ther eds a wide rangeof behaviors.," The solid line in Figure \ref{all_R50_N0_both_circulation_hist} shows no evidence of bi-modality in heat redistribution efficiency, although there is a wide rangeof behaviors."
 The dashed line in Figure 5. shows the z-distributiou if one requires the albedo to be low. Ap<0.1.," The dashed line in Figure \ref{all_R50_N0_both_circulation_hist} shows the $\varepsilon$ -distribution if one requires the albedo to be low, $A_{B}<0.1$."
 There are then mauv hieli-recirculation planets. since advection is the only wav to epress the day-side temperature iu the absence of albedo.," There are then many high-recirculation planets, since advection is the only way to depress the day-side temperature in the absence of albedo."
 Interestingly. the dashed line show teutative evidence of two separate vealss Wa 21: short-perioc ejaut planets have uniformly low albedos. t101 there appear to be two modes of heat recirculation efficiency;," Interestingly, the dashed line show tentative evidence of two separate peaks in $\varepsilon$: if short-period giant planets have uniformly low albedos, then there appear to be two modes of heat recirculation efficiency."
 We revisit this idea below., We revisit this idea below.
 Figure 6 shows tha planets idji this sampe are consistent with a low Doud alloecο., Figure \ref{all_R50_N0_both_A_hist} shows that planets in this sample are consistent with a low Bond albedo.
 Note tha this constraint is based entirely on near- and imnidàuπάτος. observations. aud is thus iudepeudeu from the claims of low albedo based on searches for reflected light (Roweetal.2008.audreferences therein)..," Note that this constraint is based entirely on near- and mid-infrared observations, and is thus independent from the claims of low albedo based on searches for reflected light \citep[][and references therein]{Rowe_2008}."
 Furthermore. lis is a coustraint on the Boud albedo. rather than the albedo in anv linited wavelength rauge.," Furthermore, this is a constraint on the Bond albedo, rather than the albedo in any limited wavelength range."
 Tn Figure 7 we plot the dimensionless dav-ide effective temperature. T4/T9. against the maxima expected dav-side teiiperature. Tey.," In Figure \ref{beta_vs_T0} we plot the dimensionless day-side effective temperature, $T_{\rm d}/T_{0}$, against the maximum expected day-side temperature, $T_{\varepsilon=0}$."
 The cvan asterisks in Figure 7 show the four hot Jupiters without temperature inversions. while most of the remaiming plaucts have inversions (Ixnutsouotal.2010).," The cyan asterisks in Figure \ref{beta_vs_T0} show the four hot Jupiters without temperature inversions, while most of the remaining planets have inversions \citep{Knutson_2010}."
. The presence or absence of an inversion docs not appear to affect the cficieucy of davnight heat recirculation., The presence or absence of an inversion does not appear to affect the efficiency of day–night heat recirculation.
 Plauets should Bie below the solid red line in Figure 7.. which denotes T.4=(2/3)τι.," Planets should lie below the solid red line in Figure \ref{beta_vs_T0}, which denotes $T_{\varepsilon=0}=(2/3)^{1/4}T_{0}$."
 Of the 21 plaucts in our sample. only one (Cl 136b) has ai day-sice effective temperature sienificautly above the T.9lint.," Of the 24 planets in our sample, only one (Gl 436b) has a day-side effective temperature significantly above the $T_{\varepsilon=0}$."
 This planet is bv far the coolest im our sample. it is," This planet is by far the coolest in our sample, it is"
Of the & line ratios we present. 6 are cousisteut with the NDR 1unodel at the lower deusities imferred by our radiative frausfer modeling (between 105 and 107 cii 2),"Of the 8 line ratios we present, 6 are consistent with the XDR model at the lower densities inferred by our radiative transfer modeling (between $10^4$ and $10^5$ $^{-3}$ )."
 We note that the models do not eo below IE cni1. so for those molecules which favor sohitious at this density. the parzuneters should be considered upper limits.," We note that the models do not go below $10^4$ $^{-1}$, so for those molecules which favor solutions at this density, the parameters should be considered upper limits."
 The remaining 2 ratios are either consistent with NDR //IINC +23) or PDR //TICN 2)) at higher deusities., The remaining 2 ratios are either consistent with XDR /HNC ) or PDR /HCN ) at higher densities.
 The fact that !the two ratios in disagreement with the others coutain nunav indicate something about the plivsical state of this particular molecule: for example. it may be traciue shock-excited material.," The fact that the two ratios in disagreement with the others contain may indicate something about the physical state of this particular molecule; for example, it may be tracing shock-excited material."
" Carcia-Burilloetal.(2010) have found previous evidence of shocks in NCC 1068 using SiO. IINCO, aud methanol (CIT4OID."," \citet{GarciaBurillo:2010} have found previous evidence of shocks in NGC 1068 using SiO, HNCO, and methanol $_3$ OH)."
 Our couclusious are similar to the aforemieutioned paper: we fud some evidence of NDR iufiueuce. but also some possible iuflueuce of shocks.," Our conclusions are similar to the aforementioned paper; we find some evidence of XDR influence, but also some possible influence of shocks."
 The discrepancy of nuuav also leud support to the oxvecn depletion scenario. because in order to match the NDR iodcels at lower deusities both of these ratios would need to be sienificautly higher than we measure.," The discrepancy of may also lend support to the oxygen depletion scenario, because in order to match the XDR models at lower densities, both of these ratios would need to be significantly higher than we measure."
 If ITCN aud TNC are beiug selectively cuhanced overICO!.. that could diive down these ratios. and lence explain why we see lower values than those predicted by the NDR uodels at deusities of ~105 3.," If HCN and HNC are being selectively enhanced over, that could drive down these ratios, and hence explain why we see lower values than those predicted by the XDR models at densities of $\sim 10^{4.5}$ $^{-3}$."
 TE is in fact also tracing eas undergoing a different excitation mechanisin. we again consider the possibility that it should not be included iu the radiative transfer analysis (as in 8??)).," If is in fact also tracing gas undergoing a different excitation mechanism, we again consider the possibility that it should not be included in the radiative transfer analysis (as in \ref{sec:sys}) )."
" When we exclude this molecule. the deusitv peaks at 10°"" iustead of 10? 7."," When we exclude this molecule, the density peaks at $10^{5.0}$ instead of $10^{4.5}$ $^{-3}$."
 Even with the shehtly larger density. there is no additional support for the PDR model. eiven the arguuents already presented that remain the same even when lis exchided from the analysis.," Even with the slightly larger density, there is no additional support for the PDR model, given the arguments already presented that remain the same even when is excluded from the analysis."
 Hosvever. this is now a second. independent line of evidence (in addition to the discussion in 822)) that suggests that nuuav be tracing a different eas conponeut than other tracers.," However, this is now a second, independent line of evidence (in addition to the discussion in \ref{sec:sys}) ) that suggests that may be tracing a different gas component than other tracers."
 This analysis adds support to previous studies that have demonstrated the NDR nature of the CND of NGC 1068., This analysis adds support to previous studies that have demonstrated the XDR nature of the CND of NGC 1068.
 For example. Uxseroetal.(2001) found support for the NDR 1uodel based on relative molecular abunudanuces derived foin LVG. simulations CICN/CO. CS/CO. .CN/TICN).," For example, \citet{Usero:2004} found support for the XDR model based on relative molecular abundances derived from LVG simulations (HCN/CO, CS/CO, ,CN/HCN)."
 Pévez-Beaupuitsctal.(2009) cae to the same conchiion based onu lue ratios of wwith lower J hnes aud Πο35. aud by the N(CNJANCIOCN) — 1 cohuun deusitv ratio.," \citet{Perez:2009} came to the same conclusion based on line ratios of with lower J lines and HCN, and by the N(CN)/N(HCN) $\sim$ 1-4 column density ratio."
 These studies are in addition to the Carcia-Burilloet conclusions in the previous paragraph., These studies are in addition to the \citet{GarciaBurillo:2010} conclusions in the previous paragraph.
 Also. Maloney(1997). demonstrated that the N-rav iradation in the CND is adequate for creating an NDR.," Also, \citet{Maloney:1997p2554} demonstrated that the X-ray irradation in the CND is adequate for creating an XDR."
 The agreement of the majority of our line intensity ratios with the NDR scenario at densities consistent with our radiative trauster modeling results provides stroug support for an NDR aud also demonstrates the utility of Z-Spec's broad bandpass aud self-cousisteut line calibration across the baud., The agreement of the majority of our line intensity ratios with the XDR scenario at densities consistent with our radiative transfer modeling results provides strong support for an XDR and also demonstrates the utility of Z-Spec's broad bandpass and self-consistent line calibration across the band.
 We used our derived total eas mass (105 ME.) to estimate the timescale for accretion outo the ceutral black hole. based on the accretion rate inferred from the ACN huuinosity.," We used our derived total gas mass $10^{8}$ ) to estimate the timescale for accretion onto the central black hole, based on the accretion rate inferred from the AGN luminosity."
 We use the simple equation where gds the cficicucy., We use the simple equation where $\eta$ is the efficiency.
 Using an accretion luminosity of 3.1«104+ (Rigbyetal.2009) (based ou [OTV] as a calibrator). the accretion rate is 0.2 + asstuning ceficicney.," Using an accretion luminosity of $3.1 \times 10^{11}$ \citep{Rigby:2009} (based on [OIV] as a calibrator), the accretion rate is 0.2 $^{-1}$ assuming efficiency."
", Estimating a timescale for accretion by AMevpf£M vields ~ 500 Myr.", Estimating a timescale for accretion by $M_{CND}/\dot M$ yields $\sim$ 500 Myr.
 This is oulv a few times he 200-300 Alyy estimated age of the nuclear stellar core (the innermost 70 pc)., This is only a few times the 200-300 Myr estimated age of the nuclear stellar core (the innermost 70 pc).
 By this estimate. the most recent star formation activity has ouly taken place Or a portion of the ACN’s depletion timescale.," By this estimate, the most recent star formation activity has only taken place for a portion of the AGN's depletion timescale."
 We rote that the timescale of accretion is proportional o efficiency. which is not well kuown. but most ikelv we are underestimating the timescale rather han overestimating. considering that the efficiency is ikelv between (Schartimannctal.2010) mt uot significantle lower than1054.," We note that the timescale of accretion is proportional to efficiency, which is not well known, but most likely we are underestimating the timescale rather than overestimating, considering that the efficiency is likely between \citep{Schartmann:2010} but not significantly lower than."
. There are also uucertainties in the total gas mass. leaving it difficult ο precisely compare the timescales.," There are also uncertainties in the total gas mass, leaving it difficult to precisely compare the timescales."
" Daviesctal.(2007) found a correlation between he age of nuclear starburst and ACN luminosity: of heir sample of 8 galaxies. NGC 1068 had one of the oldest starbursts and also is one of the most hnuniuous ACN, accreting at a significant fraction of its Eddineton tuuinosity."," \citet{Davies:2007} found a correlation between the age of nuclear starburst and AGN luminosity; of their sample of 8 galaxies, NGC 1068 had one of the oldest starbursts and also is one of the most luminous AGN, accreting at a significant fraction of its Eddington luminosity."
 It is interesting that the CND hosts a sjenificaut amount of dense eas that could be used or star formation. vet no significant star formation is currently happening.," It is interesting that the CND hosts a significant amount of dense gas that could be used for star formation, yet no significant star formation is currently happening."
 Daviesetal.(2007). theorize that here is a delav between a starburst aud ACN activity and that the starburst activity can lelp power the AGN., \citet{Davies:2007} theorize that there is a delay between a starburst and AGN activity and that the starburst activity can help power the AGN.
" ""Though the nuclear starburst began 200-300 Myr ago aud las since effectively euded. its infiuence is still significant )ecause it likely coutributed to making NGC 1068 such a strong ACN. which now is the primary influence of the dense molecular gas as we saw in the NDR analysis of he previous section."," Though the nuclear starburst began 200-300 Myr ago and has since effectively ended, its influence is still significant because it likely contributed to making NGC 1068 such a strong AGN, which now is the primary influence of the dense molecular gas as we saw in the XDR analysis of the previous section."
 We have preseuted a new broadband spectrum of NGC 1068 obtained with Z-Spec. complementing the current literature with a set of conmion-calibration ucasurciueuts of CO. its isotopes. and other hieh-deusitv tracing molecules.," We have presented a new broadband spectrum of NGC 1068 obtained with Z-Spec, complementing the current literature with a set of common-calibration measurements of CO, its isotopes, and other high-density tracing molecules."
 After extending the modcling xocedure of Wardetal.(2003) to include multiple species as in Navloretal.(2010). we use radiative rauster erids to investigate the physical properties of he CND.," After extending the modeling procedure of \citet{Ward:2003} to include multiple species as in \citet{Naylor:2010}, we use radiative transfer grids to investigate the physical properties of the CND."
 Iu our preferred ποσο]. we correct for a snall contribution frou the starburst ring for but asstune that all other molecular lines are tracing he CND only.," In our preferred model, we correct for a small contribution from the starburst ring for but assume that all other molecular lines are tracing the CND only."
 Radiative transfer analysis cannot substantially constrain the temperature of the CND. mut we argue there is evidence for a warm component. so future models should reflect ligher gas temperatures.," Radiative transfer analysis cannot substantially constrain the temperature of the CND, but we argue there is evidence for a warm component, so future models should reflect higher gas temperatures."
" ""Though the temperature aud density are degenerate. we are able to constrain thei product (pressure) well at"," Though the temperature and density are degenerate, we are able to constrain their product (pressure) well, at"
"The mass spectrum of dark-matter (DAL) halos collapsing from an initial gaussian-perturbecl density field. has long been an object of study: the formalism. most. widely: used today is that of (LOT4).. who derived a number density of the form where p(MOdAI is the number of halos in the mass range Al to AL|dM. p is the mean mass density. 9, is the critical (linear) density for collapse. and @ is the rms perturbation.","The mass spectrum of dark-matter (DM) halos collapsing from an initial gaussian-perturbed density field has long been an object of study; the formalism most widely used today is that of , who derived a number density of the form where $n(M)dM$ is the number of halos in the mass range $M$ to $M+dM$, $\bar{\rho}$ is the mean mass density, $\delta_c$ is the critical (linear) density for collapse, and $\sigma$ is the rms perturbation."
" LE we consider small scales. where σ is both ὃ, and approximately independent of mass. then the Press-Schechter form reduces to2001)."," If we consider small scales, where $\sigma$ is both $\gg\delta_c$ and approximately independent of mass, then the Press-Schechter form reduces to."
. The luminosity. function. of galaxies is commonly parametrized in the Sehechter form1976).. Le. à power-law of slope à at the low-Iuminosity. end with an exponential cutoll at a characteristic luminosity. Li.," The luminosity function of galaxies is commonly parametrized in the Schechter form, i.e., a power-law of slope $\alpha$ at the low-luminosity end with an exponential cutoff at a characteristic luminosity $L_*$."
 Wo star formation traced the πονάς clark matter distribution exactly (i.c. à Constant mass-to-light ratio). we would expect to observe à=2 at the faint end.," If star formation traced the underlying dark matter distribution exactly (i.e. a constant mass-to-light ratio), we would expect to observe $\alpha\simeq -2$ at the faint end."
 Llowever. ealaxy surveys such as the SDSS have consistently shown a dillerent slope. in the range of 19SsaZo0.9 (see2001.. Fable 2 of2005... anc 2008)).," However, galaxy surveys such as the SDSS have consistently shown a different slope, in the range of $-1.2\la\alpha\la -0.9$ (see, Table 2 of, and )."
 Phe extent to which galaxies or dark-matter halos fail to follow the clustering distribution of linear perturbation theory is known generally as “bias”. anc has been an active area of phenomenological and theoretica research (see. c.g. the introduction to and 866 of for an overview): mos relevant here is the relative bias between galaxies and DA halos. i.c. the extent to which galaxies fail to follow the uncderlving DM.," The extent to which galaxies or dark-matter halos fail to follow the clustering distribution of linear perturbation theory is known generally as “bias”, and has been an active area of phenomenological and theoretical research (see, e.g., the introduction to and 6 of for an overview): most relevant here is the relative bias between galaxies and DM halos, i.e. the extent to which galaxies fail to follow the underlying DM."
" At the low-mass end. the cliscrepancy has come to be called the ""missing satellites problem”1999).. since ive Press-Schechter theory predicts ~300 satellite galaxies for the Local Croup. while only ~40 are observed. although this number continues to increase as data [rom e.g. SDSS are analyzed."," At the low-mass end, the discrepancy has come to be called the “missing satellites problem”, since ve Press-Schechter theory predicts $\sim 300$ satellite galaxies for the Local Group, while only $\sim 40$ are observed —although this number continues to increase as data from e.g. SDSS are analyzed."
 One mieht expect that this relative bias would be a concern for those performing galaxy simulations. ancl indeed DAl-only simulations which assign galaxies to DAL halos after the faet," One might expect that this relative bias would be a concern for those performing galaxy simulations, and indeed DM-only simulations which assign galaxies to DM halos after the fact"
"We assunie (he standard inertial range magnetic turbulence power spectrum which is uncorrelated at different wavenumber vectors: Thus Eq.(17)) reduces to whose time-dependence is entirely. contained in Since SJ,(1)= 0. where (C) stands for imaginary part.we may consider AK./): The orthogonal scale ο be estimated as |D;/0;Dj]. Uhus Eq.(4)) states For WO« 1. it holds Τσ4,(WV) for n>1: moreover. AK.7)e1/n lor large n.","We assume the standard inertial range magnetic turbulence power spectrum which is uncorrelated at different wavenumber vectors: Thus \ref{expl}) ) reduces to whose time-dependence is entirely contained in Since ${\Im }J_n (W) = 0$ , where ${\Im }(\cdot)$ stands for imaginary part,we may consider ${\Re} R({\bf k}, t)$ : The orthogonal scale $1/k_\perp$ can be estimated as $|{B_i}/{\partial_j B_i}|$, thus \ref{first_order}) ) states For $W \ll1$ , it holds $J_0(W) \gg J_n(W)$ for $n \geq 1$; moreover, ${\Re} R({\bf k}, t) \sim 1/n$ for large $n$."
 We may therefore approximate (he sum in Eq.(19)) as its teri with »= 0., We may therefore approximate the sum in \ref{Fmu}) ) as its term with $n=0$ .
" Therefore Eq.(19)) vields. using E«q.(21)). four terms of type: with indexes (r.q.f.p)=(3.3.2.2). (3.2.2.3). (2.3.3.2). (2.2.3.3) lor dp, and Gregflop)=(3.3.1.1). 3.1.1.3). (1.3.3.1). (1.1.3.3) lor dp,4."," Therefore \ref{Fmu}) ) yields, using \ref{ReR}) ), four terms of type: with indexes $(r,q,l,p) = (3,3,2,2)$, $(3,2,2,3)$, $(2,3,3,2)$, $(2,2,3,3)$ for $d_{D_{XX}}$ and $(r,q,l,p) = (3,3,1,1)$, $(3,1,1,3)$, $(1,3,3,1)$, $(1,1,3,3)$ for $d_{D_{YY}}$."
 Eq. (23)), Eq.\ref{Fmu2}) )
 represents the eeneral termcontributing to the first-order(ransverse diit coellicient of a particle in a static first-order perturbed magnetic field., represents the general termcontributing to the first-ordertransverse drift coefficient of a particle in a static first-order perturbed magnetic field.
hat very few positions have a chi-squared ratio comparable o the simulations’ maxima.,that very few positions have a chi-squared ratio comparable to the simulations' maxima.
 Further. those that do. group ogether towards the galactic centre.," Further, those that do, group together towards the galactic centre."
 This is a very dilferent xtern from that observed. for the inter-£. bispectrum asvmmeltrv. implving that anv asymmetry here is unrelated.," This is a very different pattern from that observed for the $\ell$ bispectrum asymmetry, implying that any asymmetry here is unrelated."
 The grouping towards the galactic centre does raise some concern that the mask or foreground. contamination may have a subtle effect on the single-£ bispectrum., The grouping towards the galactic centre does raise some concern that the mask or foreground contamination may have a subtle effect on the $\ell$ bispectrum.
 We briellv comment on the three-point correlation function. CA). and its asvmmetry. reported. by Eriksen& (2004a).," We briefly comment on the three-point correlation function, $C^{(3)}(\theta)$, and its asymmetry reported by \cite{erik1}."
. We refer to this work for a direct definition. (in position space) of ἐπΓΑ)., We refer to this work for a direct definition (in position space) of $C^{(3)}(\theta)$.
" Phe C7(01. is related. to our delinition of the bispeetrum D uiu. by Lt is computationally advantageous to caleulate C""(8) in Fourier space — through this summation — than cirecthy over all points in the sky."," The $C^{(3)}(\theta)$, is related to our definition of the bispectrum ${\hat B_{\ell_1\ell_2\ell_3}}$ , by It is computationally advantageous to calculate $C^{(3)}(\theta)$ in Fourier space – through this summation – than directly over all points in the sky."
 Initially. we follow IZriksen&al(2004a) and choose not to include the dipole and quadrupole terms. because of their well known anomalous behaviour: thereforewe perform the summation in equation. (S8)) for (3.," Initially, we follow \cite{erik1} and choose not to include the dipole and quadrupole terms, because of their well known anomalous behaviour; thereforewe perform the summation in equation \ref{3pt}) ) for $\ell\geq$ 3."
 In Fig., In Fig.
 9. we plot the results from the Northern and Southern hemispheres (scaled. by their respective areas). with 1.2 and 3 sigma error bars from simulations.," \ref{plot3pt} we plot the results from the Northern and Southern hemispheres (scaled by their respective areas), with 1,2 and 3 sigma error bars from simulations."
 As can be appreciated from these plots. we have confirmed the finclings of Eriksen&al(2004a):: a featureless C(8) in the Northern ecliptic hemisphere.," As can be appreciated from these plots, we have confirmed the findings of \cite{erik1}: a featureless $C^{(3)}(\theta)$ in the Northern ecliptic hemisphere."
 Lt is unlikely. however. that this result is simply related to the bispectrum asvmmetrv reported in Section ?77.. as we shall now argue.," It is unlikely, however, that this result is simply related to the bispectrum asymmetry reported in Section \ref{asyminter}, as we shall now argue."
 The three-point correlation function is a function inreal space. while the bispectrum ancl power spectrum are infourier space.," The three-point correlation function is a function in space, while the bispectrum and power spectrum are in space."
 Le is often not intuitively clear how a cature in one will manifest itself in the other., It is often not intuitively clear how a feature in one will manifest itself in the other.
 The δουν) values decrease very quickly with increasing ἐν so it is the earliest. values that contribute the most. ancl determine he overall shape of C**(0).," The $|{X}_{\ell_1\ell_2\ell_3}|$ values decrease very quickly with increasing $\ell$ , so it is the earliest values that contribute the most, and determine the overall shape of $C^{(3)}(\theta)$."
 Including {Z8 terms returns C(8) values approximately an order of magnitude. less han when inclucing / 2 terms., Including $\ell\geq$ 8 terms returns $C^{(3)}(\theta)$ values approximately an order of magnitude less than when including $\ell\geq$ 2 terms.
 Vherefore. an ellect that ooks to propagate through a large range in real space. could be from a small number of values at this low-f end in Fourier space.," Therefore, an effect that looks to propagate through a large range in real space, could be from a small number of values at this $\ell$ end in Fourier space."
 Taking this into account could decrease he significance of such a detection of asymmetry., Taking this into account could decrease the significance of such a detection of asymmetry.
 In light of current uncertainties about. foreground. contamination in the low-/ multipoles (Slosar&Seljak2004:Dielewicz.Gorski&Bancday2004:Hansen.2004) we investigate if the asvmmetry survives when we exclude low-£ contributions.," In light of current uncertainties about foreground contamination in the $\ell$ multipoles \citep{slos, bgb, hbg} we investigate if the asymmetry survives when we exclude $\ell$ contributions."
 We repeat the calculation of C?(0) on the Northern and Southern hemispheres. and. cach time we increase the minimum /£ value allowed. to contribute. Ze. start the summation in equation (8)) from an increasing minimum { value.," We repeat the calculation of $C^{(3)}(\theta)$ on the Northern and Southern hemispheres, and each time we increase the minimum $\ell$ value allowed to contribute, , start the summation in equation \ref{3pt}) ) from an increasing minimum $\ell$ value."
 We do the same for our simulations., We do the same for our simulations.
 For each computation. we perform a reduced: chi-squared analysis over the 8 range 0.70]. and in Table 5. we record the fraction of ~500 simulations that [ind a lower X7 value than the Northern and Southern WALADP results.," For each computation, we perform a reduced chi-squared analysis over the $\theta$ range [0,70], and in Table \ref{Tab3pt} we record the fraction of $\sim$ 500 simulations that find a lower $\chi^2$ value than the Northern and Southern WMAP results."
 Our simulations also use the ecliptic poles., Our simulations also use the ecliptic poles.
 We stress that we are not interested in the level of significance of this anomaly (we have not accounted. Lor selection. cllects) but how the significance changes as we exclude low-f terms.," We stress that we are not interested in the level of significance of this anomaly (we have not accounted for selection effects), but how the significance changes as we exclude $\ell$ terms."
" We find the asymmetry does not diminish as we exclude the lower-f contributions. implving an asvmmetrythat is a widespread. feature in£ space,"," We find the asymmetry does not diminish as we exclude the $\ell$ contributions, implying an asymmetrythat is a widespread feature in$\ell$ space."
 We cmphasise that this asvnimetry ds. dilferentto that explored. above with £2 and J. as here we are not independent of the power spectrum.," We emphasise that this asymmetry is differentto that explored above with $I^3_\ell$ and $J^3_\ell$ , as here we are not independent of the power spectrum."
approximately 3 times solar was needed (N47=—10 from a solar abundance model). and a Kerr space-time geometry was preferred over a Schwarzschild one (rin = 3—L. ro= 8.3 gravitational radii).,"approximately 3 times solar was needed $\Delta \chi^2
= -10$ from a solar abundance model), and a Kerr space-time geometry was preferred over a Schwarzschild one $r_{\mathrm{in}}$ = 3–4, $r_{\mathrm{out}}$ = 8.3 gravitational radii)."
 However. a model with a reflection fraction of unity can only give an adequate fit to the data (47.=258/180 dof) between 2 anc 10 keV. and it cannot fully account for the sharp spectral drop a ~7 keV. Allowing the reflection fraction to increase to 9 resultec in a steeper continuum (LY=1.68 from 1.54) and an improved tit (47=201/179dof: see Fig 2).," However, a model with a reflection fraction of unity can only give an adequate fit to the data $\chi^2=258/180~ dof$ ) between 2 and 10 keV, and it cannot fully account for the sharp spectral drop at $\sim$ 7 keV. Allowing the reflection fraction to increase to 9 resulted in a steeper continuum $\Gamma=1.68$ from 1.54) and an improved fit $\chi^2 = 201/179 ~ dof$; see Fig \ref{fig:fits}) )."
 The high value of 7? can be decreased by increasing the iron abundance of the reflector., The high value of $R$ can be decreased by increasing the iron abundance of the reflector.
 The .ionisation parameter .is £κ=750 erg ems l., The ionisation parameter is $\xi \approx 750$ erg cm $^{-1}$.
 At this. leve of ionisation the reflection-dominated model predicts a strong ike Fe Ka line with an equivalent width of 1.8 keV. This mode clearly is not as good at describing the data as the partial covering model., At this level of ionisation the reflection-dominated model predicts a strong He-like Fe $\alpha$ line with an equivalent width of 1.8 keV. This model clearly is not as good at describing the data as the partial covering model.
 It is extremely difficult for ionised reflection to account for he depth of the drop at ~7 keV without invoking a very extreme Fe abundance and/or reflection fraction., It is extremely difficult for ionised reflection to account for the depth of the drop at $\sim$ 7 keV without invoking a very extreme Fe abundance and/or reflection fraction.
" The blue wing of a very strong relativistically broadened and highly redshifted iron Κά line can produce a drop at 7 keV. A power- plus line (rest energy 6.4 keV) provides a good fit (47=182/180 dof) with an inner dise radius ry,«2r, and an underlying power-law slope of [=2.08.0/15.", The blue wing of a very strong relativistically broadened and highly redshifted iron $\alpha$ line can produce a drop at 7 keV. A power-law plus line (rest energy 6.4 keV) provides a good fit $\chi^{2} = 182 / 180 ~ dof$ ) with an inner disc radius $r_{\rm{in}} < 2 r_{\rm{g}}$ and an underlying power-law slope of $\Gamma=2.08_{-0.13}^{+0.25}$.
" However. the very high equivalent width (AV=5 keV) required to explain the size of the 7 keV drop is difficult to justify,"," However, the very high equivalent width $EW=5$ keV) required to explain the size of the 7 keV drop is difficult to justify."
 Alternatively. a reflection spectrum from cold material will contain a neutral iron edge.," Alternatively, a reflection spectrum from cold material will contain a neutral iron edge."
 Fitting with aRAY model gives a good fit (47=192/182 dof) with a reflection fraction 7»200 (essentially just reflection) and an iron overabundance 7 3., Fitting with a model gives a good fit $\chi^{2}= 192 / 182 ~ dof$ ) with a reflection fraction $\gs200$ (essentially just reflection) and an iron overabundance $>3$ .
 The dominance of the reflected compared to the primary emission. and the weakness of the associated iron Ka line appear fatal for this interpretation.," The dominance of the reflected compared to the primary emission, and the weakness of the associated iron $\alpha$ line appear fatal for this interpretation."
 Fig 3 shows the 0.1—10 keV EPIC pn light curve in 200 s bins (The MOS light curves are essentially identical to the pn light curve)., Fig \ref{fig:light_curve} shows the 0.1–10 keV EPIC pn light curve in 200 s bins (The MOS light curves are essentially identical to the pn light curve).
 The source again showed strong (factor of —4 change during the observation) and rapid variability. as previously seen by (Leighly 19992).," The source again showed strong (factor of $\sim$ 4 change during the observation) and rapid variability, as previously seen by (Leighly 1999a)."
 The relatively low apparent luminosity of the source during this observation. ου10251072 erg 1 means that the rapid variability translates to a fairly modest rate of change of luminosity.," The relatively low apparent luminosity of the source during this observation, $L_{0.2-10}
\approx 8 \times 10^{42}$ erg $^{-1}$, means that the rapid variability translates to a fairly modest rate of change of luminosity."
" For example. the rapid rise toward the end of the observation corresponds to a luminosity increase of AL/Atm2.4.107"" erg 7."," For example, the rapid rise toward the end of the observation corresponds to a luminosity increase of $\Delta L/\Delta t
\approx 2.4 \times 10^{39}$ erg $^{-2}$."
 The corresponding radiative efficiency. using the argument of Fabian (1979). is only #)=0.2 per cent.," The corresponding radiative efficiency, using the argument of Fabian (1979), is only $\eta \gs
0.2$ per cent."
 Light curves in various energy bands were analysed to search for spectral variability., Light curves in various energy bands were analysed to search for spectral variability.
 The variations in each band appear very similar., The variations in each band appear very similar.
 The fractional variability amplitude {ων (see Edelson 2001) was measured in different energy bands (Fig 4) using light curves binned to 1000 s and is consistent with a constant (fractional) variability amplitude., The fractional variability amplitude $F_{\rm{var}}$ (see Edelson 2001) was measured in different energy bands (Fig \ref{fig:rms_spec}) ) using light curves binned to 1000 s and is consistent with a constant (fractional) variability amplitude.
 À cross-correlation analysis using the discrete correlation function (DCF: Edelson Krolik 1988) showed the light curves to be correlated at zero-lag. with no evidence for any time lags.," A cross-correlation analysis using the discrete correlation function (DCF; Edelson Krolik 1988) showed the light curves to be correlated at zero-lag, with no evidence for any time lags."
 Hardness ratios were also examined to search for spectral variability (see Fig 539., Hardness ratios were also examined to search for spectral variability (see Fig \ref{fig:hr}) ).
 These do show spectral variability (the lower two panels) but are uncorrelated with flux., These do show spectral variability (the lower two panels) but are uncorrelated with flux.
" As a final check of (flux-correlated) spectral variability separate pn spectra were extracted from time intervals when the source flux was below the mean (low flux"" spectrum) and bove the mean (“high flux” spectrum).", As a final check of (flux-correlated) spectral variability separate pn spectra were extracted from time intervals when the source flux was below the mean (“low flux” spectrum) and above the mean (“high flux” spectrum).
 The ratio of the two spectra (Fig 61) was consistent with a constant. re-enforcing the claim of no flux-dependent spectral variability.," The ratio of the two spectra (Fig \ref{fig:ratio}) ) was consistent with a constant, re-enforcing the claim of no flux-dependent spectral variability."
 Unfortunately. the limited number of counts above 7 keV (zx140 source counts in the pn) mitigates against a detailed analvsisof the variability of the 7 keV spectral feature.," Unfortunately, the limited number of counts above 7 keV $\approx 140$ source counts in the pn) mitigates against a detailed analysisof the variability of the 7 keV spectral feature."
kinematics.,kinematics.
" The existing data are not sullicient to really constrain the input parameters of the static model. we developed: we can only provide rough estimates of the main parameters such as: the pattern speed Q,=700ct100 tthe CR radius Aeg;=4.9irae the angle between the bar and the line of nodes 65,,,=55°+20°"," The existing data are not sufficient to really constrain the input parameters of the static model we developed: we can only provide rough estimates of the main parameters such as: the pattern speed $\Omega_p = 700 \pm 100$ , the CR radius $R_{CR} = 4.9^{+1.4}_{-0.8}$, the angle between the bar and the line of nodes $\theta_{bar} = 55\degr \pm
20\degr$."
 ]t may be possible to narrow the range of possible values for the pattern speed. of the wave. as its value determines the location of the resonances. but this would. require a full hyerodynamical simulation.," It may be possible to narrow the range of possible values for the pattern speed of the wave, as its value determines the location of the resonances, but this would require a full hydrodynamical simulation."
 A deeper high-resolution image would also help to follow the gascous distribution to &ereater galactocentric distances., A deeper high-resolution image would also help to follow the gaseous distribution to greater galactocentric distances.
As discussed in Section 2.L. our numerical procedure of using a small number of particles cau lead to potential problems when terminating the simulation at a specified time. as some loug-lived particles may stil be active.,"As discussed in Section \ref{subsec:recording}, our numerical procedure of using a small number of particles can lead to potential problems when terminating the simulation at a specified time, as some long-lived particles may still be active."
 In a real debris «isk. however. grain collisions would limit the maximum lifetime of particles.," In a real debris disk, however, grain collisions would limit the maximum lifetime of particles."
 In Figure 6 we show he results of a sinulation of a Jupiter mass planet with ag;LLs AU and ej=0.5 system contalning 500 test paricles with 6=0.1., In Figure \ref{fig:stoptime} we show the results of a simulation of a Jupiter mass planet with $a_{pl} \sim 44.8$ AU and $e_{pl} = 0.5$ system containing 500 test particles with $\beta = 0.1$.
" The est particles are released frou parent bodies wihi 1.4«mdpy/ap<2.0. 0.0<eun0.2 and 0.0.<ip,5.0."," The test particles are released from parent bodies with $1.4<a_{pb}/a_{pl}<2.0$, $0.0<e_{pb}<0.2$ and $0.0\degree<i_{pb}<8.0\degree$."
 The three panels show the cust distribution when the simulation is terminated after 11.7 nillion years. LO million years aud 153 million years. correspoucing ο10%... and of particles reniainiug active.," The three panels show the dust distribution when the simulation is terminated after 11.7 million years, 40 million years and 153 million years, corresponding to, and of particles remaining active."
 It can be seen that tle basic clisk structure is createc rapidly. aud the final few j»articles have a Iunited impact despite having lifetimes tay times tve mecdiau.," It can be seen that the basic disk structure is created rapidly, and the final few particles have a limited impact despite having lifetimes many times the median."
 The lougest-lived particles teud to sharpen’ the distribution aud briig out finer detail. which is hidden in our simlated. observations (telescope FWHM 10 AU).," The longest-lived particles tend to `sharpen' the distribution and bring out finer detail, which is hidden in our simulated observations (telescope FWHM 40 AU)."
 TIus. the inipositio1 of a cut-off lifetime (whether as a numerical convenieuce OF as a substitute for dust destruction through collisious) is of little consequence as long as most (>90%)) particles have already been accreted or ejected.," Thus, the imposition of a cut-off lifetime (whether as a numerical convenience or as a substitute for dust destruction through collisions) is of little consequence as long as most $>$ ) particles have already been accreted or ejected."
 In this section we use the svuthetic disk catalogue to select planetary couliguratious which generate structures similar to the observed debris disk systems Vega. aaud Fomalhaut.," In this section we use the synthetic disk catalogue to select planetary configurations which generate structures similar to the observed debris disk systems Vega, and Fomalhaut."
 We then optimise the planetary aucl parent body parameters. ancl iucrease the resolution. (by increasing the number of test particles) to produce more accurate synthetic observations. which we compare to previous observational aud uumerical studies of these three debris disk systenms.," We then optimise the planetary and parent body parameters, and increase the resolution (by increasing the number of test particles) to produce more accurate synthetic observations, which we compare to previous observational and numerical studies of these three debris disk systems."
 Vega (a Lyrae) was the first star identified with an iulrared excess associated. with a disk of dusty material (Atumannetal.1981) aud has become the prototype lor so-called *Vega-tvpe stars., Vega $\alpha$ Lyrae) was the first star identified with an infrared excess associated with a disk of dusty material \citep{aumann84} and has become the prototype for so-called `Vega-type' stars.
 Su»equent observations wit1 SCUBA (Hollandetal.1998) and he IRAM Plateau de Bure Iuterferoueter (Wilneretal.2002) confirm that Vega is surrounded by a dusty disk. with two proimiuert clumps of emission.," Subsequent observations with SCUBA \citep{holland98}
 and the IRAM Plateau de Bure Interferometer \citep{whk02} confirm that Vega is surrounded by a dusty disk, with two prominent clumps of emission."
 Wiueretal.(2002) proposed that the win lobes of emission seen iutheseiiuages are caused by dust trappecd in MMBs with a massive. ecceuric planet (My=3Mju. ey=0.6 orbiting Vega with a ool 10 Al>.," \citet{whk02} proposed that the twin lobes of emission seen in these images are caused by dust trapped in MMRs with a massive, eccentric planet $M_{pl} = 3 M_{Jup}$, $e_{pl} = 0.6$ ) orbiting Vega with a of 40 AU."
 In their model. the iniial test. particle lougitudes of perilieljon were coustrained to be aligued cosely with the plauet's lougitude of peribeliou (which is reasonable for a high eccentricity," In their model, the initial test particle longitudes of perihelion were constrained to be aligned closely with the planet's longitude of perihelion (which is reasonable for a high eccentricity"
same region as is visible in the 10.75gam image indicates that. globally. stellar emission contributes of the flux in the ISO 26.75 jim band.,"same region as is visible in the $\lambda\,6.75~\mu$ m image indicates that, globally, stellar emission contributes of the flux in the ISO $\lambda \,6.75~\mu$ m band."
 This is within the range estimated for the absolute calibration error on typical ISOCAM observations (£15-20%.. Sect. 3))," This is within the range estimated for the absolute calibration error on typical ISOCAM observations $\pm\,15\,\rightarrow\,20$, Sect. \ref{observations_reductions}) )"
 and therefore is justifiably neglected for the galaxy as a whole., and therefore is justifiably neglected for the galaxy as a whole.
 In the region of the halo itself. no stellar halo is observed (Fig.," In the region of the halo itself, no stellar halo is observed (Fig."
 6. Inset) to the limits of the available 2MASS data., \ref{iso_optical} Inset) to the limits of the available 2MASS data.
 The rms noise of this map places an upper limit of σσ.=0.0002 mJy aresec7 on the stellar contribution to the 26.75 gm band. re. an upper limit. of approximately on the lowest contour shown in Fig.," The rms noise of this map places an upper limit of $\sigma_*\,=\,0.0002$ mJy $^{-2}$ on the stellar contribution to the $\lambda \,6.75~\mu$ m band, i.e. an upper limit of approximately on the lowest contour shown in Fig."
 2aa. Fig.," \ref{iso_map}a a, Fig."
 6 (main) and Fig., \ref{iso_optical} (main) and Fig.
 7 and a lower percentage for the higher contours.," \ref{iso_halpha}
 and a lower percentage for the higher contours."
 In a rectangular region aligned with the major axis centered on the north-east the maximum contribution of the stellar halo is7%., In a rectangular region aligned with the major axis centered on the north-east the maximum contribution of the stellar halo is.
. Since the fraction of stellar emission could vary from location to location. a possible stellar contribution to the halo is discussed further in Sect. 4.3..," Since the fraction of stellar emission could vary from location to location, a possible stellar contribution to the halo is discussed further in Sect. \ref{halo_emission}."
 Fig., Fig.
" 2aa shows strong PAH emission along the disk of NGC 5529, although the entire disk has not been mapped due to the truncated field of view."," \ref{iso_map}a a shows strong PAH emission along the disk of NGC 5529, although the entire disk has not been mapped due to the truncated field of view."
" The emission reaches a maximum at the nucleus and a two-dimensional Gaussian fit to this maximum gives a position of RA = [4% I5"" 34.p + 0.2. DEC = 36° 13’ 36"" + 2"", which agrees with the optical center (Table 5))."," The emission reaches a maximum at the nucleus and a two-dimensional Gaussian fit to this maximum gives a position of RA = $^{\rm h}$ $^{\rm m}$ $^{\rm s}$ $\pm$ $^{\rm s}$, DEC = $^\circ$ $^\prime$ $^{\prime\prime}$ $\pm$ $^{\prime\prime}$, which agrees with the optical center (Table \ref{basic_parameters}) )."
" Several other peaks are also observed along the major axis. a double peak « 40"" from the nucleus along the north-west major axis and a single peak + 50"" to the south-east."," Several other peaks are also observed along the major axis, a double peak $\approx$ $^{\prime\prime}$ from the nucleus along the north-west major axis and a single peak $\approx$ $^{\prime\prime}$ to the south-east."
 This structure is shown more readily in the plot of intensity along the major axis shown in Fig., This structure is shown more readily in the plot of intensity along the major axis shown in Fig.
 6aa. Note that there is no change to the positions and structure described here if a stellar contribution. as outlined in Sect. 4.1..," \ref{slices_fig}a a. Note that there is no change to the positions and structure described here if a stellar contribution, as outlined in Sect. \ref{band_contributions},"
 is subtracted from the map. pixel by pixel.," is subtracted from the map, pixel by pixel."
 Irwin Madden (2006) found that the PAH emission in NGC 5907 followed primarily the molecular gas distribution in that galaxy. with a fainter wing of emission extending farther out into the HI dominated region.," Irwin Madden (2006) found that the PAH emission in NGC 5907 followed primarily the molecular gas distribution in that galaxy, with a fainter wing of emission extending farther out into the HI dominated region."
 To our knowledge. there is no published CO map of NGC 5529 with which to compare the PAH distribution.," To our knowledge, there is no published CO map of NGC 5529 with which to compare the PAH distribution."
 The HI distribution (Kregel et al., The HI distribution (Kregel et al.
 2004b) has a hole at the nucleus and the HI emission rises on either side of the nucleus to maxima that are at positions farther out in radius than the peaks of Fig., 2004b) has a hole at the nucleus and the HI emission rises on either side of the nucleus to maxima that are at positions farther out in radius than the peaks of Fig.
 6aa. Therefore. the HI distribution does not resemble that of the PAHs. to the limits of these measurements.," \ref{slices_fig}a a. Therefore, the HI distribution does not resemble that of the PAHs, to the limits of these measurements."
 We suspect that the bulk of the emission will follow that of the CO. as in NGC 5907. once such measurements are made.," We suspect that the bulk of the emission will follow that of the CO, as in NGC 5907, once such measurements are made."
 Fig., Fig.
 2aa and especially Fig., \ref{iso_map}a a and especially Fig.
 6 show that the PAH emission extends far from the plane of NGC 5529., \ref{iso_optical} show that the PAH emission extends far from the plane of NGC 5529.
 This detection has been possible because of the high sensitivity and dynamic range (Table 7)) of the beam-switehing observations., This detection has been possible because of the high sensitivity and dynamic range (Table \ref{observing_map}) ) of the beam-switching observations.
 The high-latitude emission shows considerable structure. a result that was also seen for NGC 5907.," The high-latitude emission shows considerable structure, a result that was also seen for NGC 5907."
" Indeed. among other features. two discrete are-like structures can be seen on the north-east side of the major axis: one above the nucleus extending to RA = 14"" [5"" 36. DEC = 36° 14 I0"". and one farther to the east extending to RA = 14"" 15"" 39, DEC = 36° 13’ 50""."," Indeed, among other features, two discrete arc-like structures can be seen on the north-east side of the major axis: one above the nucleus extending to RA $\approx$ $^{\rm h}$ $^{\rm m}$ $^{\rm s}$, DEC $\approx$ $^\circ$ $^\prime$ $^{\prime\prime}$ and one farther to the east extending to RA $\approx$ $^{\rm h}$ $^{\rm m}$ $^{\rm s}$ , DEC $\approx$ $^\circ$ $^\prime$ $^{\prime\prime}$."
 Are- features similar to these have been observed in other ISM tracers in other edge-on galaxies (e.g. Lee et al., Arc-like features similar to these have been observed in other ISM tracers in other edge-on galaxies (e.g. Lee et al.
 2001)., 2001).
 At the 3c contour shown in Fig. 6..," At the $3\,\sigma$ contour shown in Fig. \ref{iso_optical},"
 the emission extends to 7=40” (zx8.5 kpc) above the plane.," the emission extends to $z\,\approx\,40^{\prime\prime}$ $z\,\approx\,8.5$ kpc) above the plane."
 Note that there ts o emission (Fig., Note that there is no emission (Fig.
 6 Inset). to the 2c limit of themap. to this height.," \ref{iso_optical} Inset), to the $\sigma$ limit of the, to this height."
" The maximum extent of the stellar emission at the 2c- limit of this map is z=20"" (4.2 kpe)."," The maximum extent of the stellar emission at the $\sigma$ limit of this map is $z\,\approx\,20^{\prime\prime}$ (4.2 kpc)."
 Even higher sensitivity can be achieved by averaging the emission over a broad swath of the major axis (91 i width) d plotting the result as a function of latitude. z. above and below the plane (note that no correction for inclination ts required).," Even higher sensitivity can be achieved by averaging the emission over a broad swath of the major axis $91^{\prime\prime}$ in width) and plotting the result as a function of latitude, $z$, above and below the plane (note that no correction for inclination is required)."
 Not only does this approach improve the sensitivity. but it also allows us to probe whether a broad scale halo exists ralround the galaxy. rather than only discrete extensions.," Not only does this approach improve the sensitivity, but it also allows us to probe whether a broad scale halo exists around the galaxy, rather than only discrete extensions."
 The result is shown in Fig., The result is shown in Fig.
 6bb (main figure)., \ref{slices_fig}b b (main figure).
 Note that the dynamic range of this slice is very high. i.e. 1773/1 (peak/rms).," Note that the dynamic range of this slice is very high, i.e. 1773/1 (peak/rms)."
 Various functions were fitted to these data including an exponential. a Gaussian. a hyperbolic secant. and the square of a hyperbolic secant [that is. so that the range of functions used by others for optical data (Sect. 2))," Various functions were fitted to these data including an exponential, a Gaussian, a hyperbolic secant, and the square of a hyperbolic secant [that is, so that the range of functions used by others for optical data (Sect. \ref{ngc5529}) )"
 are also considered here]., are also considered here].
 The best fit was to a simple Gaussian. the result shown as the grey curve in Fig.," The best fit was to a simple Gaussian, the result shown as the grey curve in Fig."
 6bb (main)., \ref{slices_fig}b b (main).
 The fit is excellent. with the exception of the broad scale wings that are seen in the data but not in the model (see below).," The fit is excellent, with the exception of the broad scale wings that are seen in the data but not in the model (see below)."
" The displayed best fit Gaussian curve. which ts wider than the true distribution because of beam smoothing. has og=4.6"" (979 pe). or a full width at half maximum of FWHM = 10.8” (2.3 kpe)."," The displayed best fit Gaussian curve, which is wider than the true distribution because of beam smoothing, has $\sigma_G\,=\,4.6^{\prime\prime}$ (979 pc), or a full width at half maximum of FWHM = $10.8^{\prime\prime}$ (2.3 kpc)."
" We reproduced this curve by modeling the z distribution of the disk as a Gaussian of FWHM = 8.0"" (1.7 kpe) and then convolving with the Gaussian PSF of the observations (FWHMps,;- = 7.2"". Table 7))."," We reproduced this curve by modeling the $z$ distribution of the disk as a Gaussian of FWHM = $8.0^{\prime\prime}$ (1.7 kpc) and then convolving with the Gaussian PSF of the observations $_{PSF}$ = $7.2^{\prime\prime}$, Table \ref{observing_map}) )."
" Thus. after beam correction. the bulk of the PAH emission can be represented by a disk that declines as aGaussian of FWHM = 8.0"" (1.7 kpe). or cc;=3.4"" (726 pe)."," Thus, after beam correction, the bulk of the PAH emission can be represented by a disk that declines as aGaussian of FWHM = $8.0^{\prime\prime}$ (1.7 kpc), or $\sigma_G\,=\,3.4^{\prime\prime}$ (726 pc)."
" We repeated the same exercise over the same region using the Hadata®.. finding a Gaussian of FWHM = 13.2” (2.8 kpe). οσο=5.6"" (1.2 kpe)."," We repeated the same exercise over the same region using the $\alpha$, finding a Gaussian of FWHM = $13.2^{\prime\prime}$ (2.8 kpc), or $\sigma_G\,=\,5.6^{\prime\prime}$ (1.2 kpc)."
 Thus. the H« disk is =1.6x thicker than the PAH disk.whenthe same regions and criteria are used.," Thus, the $\alpha$ disk is $\approx\,1.6\,\times$ thicker than the PAH disk,whenthe same regions and criteria are used."
 The faint wings of PAH emission in the very high sensitivity plot of Fig., The faint wings of PAH emission in the very high sensitivity plot of Fig.
 6bb (main) reveal that. as suggested by the appearance of Figs.," \ref{slices_fig}b b (main) reveal that, as suggested by the appearance of Figs."
 2aa and Fig. 6.. ," \ref{iso_map}a a and Fig. \ref{iso_optical}, ,"
a broad seale halo of, a broad scale halo of
Cluster lumuinositw functions (CLEs) and colour distributions are the most important diagnostics in the study. of globular and compact star cluster populations in nearby galaxies.,Cluster luminosity functions (CLFs) and colour distributions are the most important diagnostics in the study of globular and compact star cluster populations in nearby galaxies.
 For the old globular cluster (€€) systems in. e.g.. the Galaxy. M31. AIST. ancl old. elliptical galaxies. the CLE shape is well-established: it is roughly. Gaussian. with the peak or turnover magnitude at Ac7.4 mag and a Gaussian FWLIAL of ~3 mag (Whitmore et al.," For the old globular cluster (GC) systems in, e.g., the Galaxy, M31, M87, and old elliptical galaxies, the CLF shape is well-established: it is roughly Gaussian, with the peak or turnover magnitude at $M_V^0 \simeq -7.4$ mag and a Gaussian FWHM of $\sim 3$ mag (Whitmore et al."
 1995. Larris 1996. 2001. Ashman Zepf 1998. Llarris. Harris MeLaughlin. 1998).," 1995, Harris 1996, 2001, Ashman Zepf 1998, Harris, Harris McLaughlin 1998)."
 The well-stuclied voung star cluster (YSC) population in the LAIC and the Galactic. open, The well-studied young star cluster (YSC) population in the LMC and the Galactic open
analysis that there is no evidence of dust in 0050904 and 0090423.,analysis that there is no evidence of dust in 050904 and 090423.
 We find possible evidence of low extinction in 0080913 consistent with an SMC extinetion curve. however the optical spectrum and NIR photometry alone are consistent with a power-law with no dust extinction.," We find possible evidence of low extinction in 080913 consistent with an SMC extinction curve, however the optical spectrum and NIR photometry alone are consistent with a power-law with no dust extinction."
 In no case did we find evidence for any other extinction curve tthe “SN-origin’ curve of ?))., In no case did we find evidence for any other extinction curve the `SN-origin' curve of \citealt{maiolino}) ).
 Comparing to a much larger spectroscopic sample of GRB extinctions. we find a distinct absence of extinction at high redshifts.," Comparing to a much larger spectroscopic sample of GRB extinctions, we find a distinct absence of extinction at high redshifts."
 While the high redshift sample is very small and some of the effect may be explained by restframe UV-selection bias. this is not the whole picture. as all the z>6 GRBs would have been detected at the typical Ay of lower redshift reddened bursts.," While the high redshift sample is very small and some of the effect may be explained by restframe UV-selection bias, this is not the whole picture, as all the $z>6$ GRBs would have been detected at the typical $A_V$ of lower redshift reddened bursts."
 This hints that there is less dust along sightlines to the highest redshift GRBs. indicating less dust in the early Universe.," This hints that there is less dust along sightlines to the highest redshift GRBs, indicating less dust in the early Universe."
 From these results. we also infer an extremely low dust-to-metals ratio in GRBs at high redshift. suggestive either of efficient dust destruction. or a delay of at least several million years between the formation of metals and the formation and growth of dust.," From these results, we also infer an extremely low dust-to-metals ratio in GRBs at high redshift, suggestive either of efficient dust destruction, or a delay of at least several million years between the formation of metals and the formation and growth of dust."
triaxial torque is {ο spin up the planet again.,triaxial torque is to spin up the planet again.
 This counter-action of the triaxial torque is svmnmetric. in that if the planets spin accelerated during the two complete orbits the torque will rectifv. its rotation.," This counter-action of the triaxial torque is symmetric, in that if the planet's spin accelerated during the two complete orbits the torque will rectify its rotation."
 Thus. anv deviation of spin rate from (he resonance value will be iutomatically corrected by the Wiaxial torque.," Thus, any deviation of spin rate from the resonance value will be automatically corrected by the triaxial torque."
 The planet is trapped in a stable equilibrium., The planet is trapped in a stable equilibrium.
 This is not the case when the initial anglee αἱ periastvon is 23/2., This is not the case when the initial angle at periastron is $\pm\pi/2$.
/ Anv net lage in phase angleex will cause a nonzero (iasial torque in (he same direction. causing the planet to continue (o spin down.," Any net lag in phase angle will cause a nonzero triaxial torque in the same direction, causing the planet to continue to spin down."
 In this configuration of unstable ecuilibrium. the torque assists the tides to swiltlv tun the planet around and Iunge through the resonance.," In this configuration of unstable equilibrium, the torque assists the tides to swiftly turn the planet around and lunge through the resonance."
 The same mechanics work for a circular orbit (e=0)., The same mechanics work for a circular orbit $e=0$ ).
 One may ask then. how a planet initially αἱ 0(0)=0 can be entrapped at all. if al some point afterwards it inevitably. (urns sidewise with respect to the star?," One may ask then, how a planet initially at $\theta(0)=0$ can be entrapped at all, if at some point afterwards it inevitably turns sidewise with respect to the star?"
 The answer is that while theplanet turns through 7/2 wilh respect to the star. the spin rate will change because of the (riaxial torque. and will no longer be equal to (124/2)».," The answer is that while theplanet turns through $\pi/2$ with respect to the star, the spin rate will change because of the triaxial torque, and will no longer be equal to $(1+q/2)n$."
 In other words. the condition of resonance passage is a certain area of the 2D phase Returning to Fig. 5..," In other words, the condition of resonance passage is a certain area of the 2D phase Returning to Fig. \ref{escape.fig},"
 we note that the point of the phase space τρίο)=z/2.(0)—2n) does not actually belong to this area. because the passage through resonance is not immediate.," we note that the point of the phase space $\{\theta(0)=\pi/2,\dot\theta(0)=2\,n\}$ does not actually belong to this area, because the passage through resonance is not immediate."
 Indeed. the planet has to reallign its phase space parameters in (wo upward swings before it traverses the resonance at about. 80 vr alter the start of integration.," Indeed, the planet has to reallign its phase space parameters in two upward swings before it traverses the resonance at about 80 yr after the start of integration."
 In order to map (he phase space oLznaediale resonance passages. we performed several hundred short-term integrations (for 1000 vr) varying the initial (0) with a step of 5-10η and 6(0) with a step of 1-10? rad.," In order to map the phase space of resonance passages, we performed several hundred short-term integrations (for 1000 yr) varying the initial $\dot\theta(0)$ with a step of $5\cdot10^{-5}n$ and $\theta(0)$ with a step of $1\cdot10^{-5}$ rad."
 The mapped section of this space for our Alercury-like planet and (he 2:1 resonance is shown in Fig., The mapped section of this space for our Mercury-like planet and the 2:1 resonance is shown in Fig.
 G as ereen-shaclecl area.," \ref{needle.fig}
 as green-shaded area."
 If the planet al periastron turns up within the green area. 1t immediately. (raverses (he resonance.," If the planet at periastron turns up within the green area, it immediately traverses the resonance."
 This green corridor is extremely narrow in 8 al A/n>2 and oceupies a tiny [fraction of the phase space.," This ""green corridor"" is extremely narrow in $\theta$ at $\dot\theta/n>2$ and occupies a tiny fraction of the phase space."
 The black dots display the periastron positions of the planet., The black dots display the periastron positions of the planet.
 The dots are lined up in live consecutive libration trajectories. each following one probing lower spin rates around 9=7/2.," The dots are lined up in five consecutive libration trajectories, each following one probing lower spin rates around $\theta=\pi/2$."
" The final lap hits the ""needles eve” and falls through the resonance. never to return to it."," The final lap hits the “needle's eye"" and falls through the resonance, never to return to it."
 Note that the last but one trajectory actually reached ϱ=27 at & very close to 7/2. but the Gpping point is al slightly smaller 8.," Note that the last but one trajectory actually reached $\dot\theta=2\,n$ at $\theta$ very close to $\pi/2$, but the tipping point is at slightly smaller $\theta$ ."
 The area of the phase space above the resonant spin rate. through which a Mercury-Iike," The area of the phase space above the resonant spin rate, through which a Mercury-like"
In order to understand. formation and evolution of galaxies. il is necessary to search and study high-redshift galaxies.,"In order to understand formation and evolution of galaxies, it is necessary to search and study high-redshift galaxies."
 Lyman Break Galaxies (LDOs). which are selected bv rest- UV broad band photometry (e.g. Steidel&ILunilon1992:etal. 1995)). make the largest sample of galaxies at 2Z3 among various populations selected. through different methods. and their statistical and individual studies have been made extensively.," Lyman Break Galaxies (LBGs), which are selected by rest-frame UV broad band photometry (e.g., \citealp{ste92,ste95}) ), make the largest sample of galaxies at $z \gtrsim 3$ among various populations selected through different methods, and their statistical and individual studies have been made extensively."
 For instance. based on photometric samples. rest-lrame UV. TJuminosity fictions (UVLES) ol LBGs al z~3—6 are derived (e.g.. Steideletal.1999:Iwata2003.2007:etal.2006:Bouwensοἱ 2007)) and attempts to measure the LF have been made even al higher redshifts (e.g.. Richardetal.2006.2008:Stark2007:BouwensStan-wavetal.2008b:Oeschοἱ 2008)).," For instance, based on photometric samples, rest-frame UV luminosity functions (UVLFs) of LBGs at $z \sim 3-6$ are derived (e.g., \citealp{ste99,iwa03,iwa07,leh03,ouc04a,bec06,saw06,yos06,bou07}) ) and attempts to measure the LF have been made even at higher redshifts (e.g., \citealp{ric06,ric08,star07,bou08,sta08b,oes08}) )."
 These studies are revealing the cosmic star formation history: the cosmic star formation rate density rises [rom ze5—6 (02 —2-3 and turns to decline toward z~0 (e.g.. Hopkins&Beacom 2006)).," These studies are revealing the cosmic star formation history; the cosmic star formation rate density rises from $z \sim 5-6$ to $z \sim 2-3$ and turns to decline toward $z \sim 0$ (e.g., \citealp{hop06}) )."
 ILowever two different evolutions of UVLFs are claimed., However two different evolutions of UVLFs are claimed.
 One is that [rom z~5 to ~3 the number density of UV faint galaxies increases while (hat of bright galaxies remains almost constant (SawickiIwataοἱal. 2007)).," One is that from $z \sim 5$ to $\sim3$ the number density of UV faint galaxies increases while that of bright galaxies remains almost constant \citealp{saw06,iwa07}) )."
 The other is that while (he number density of UV faint galaxies remains constant. Chat of bright galaxies increase Yoshidaetal.2006:Bouwens 2007)).," The other is that while the number density of UV faint galaxies remains constant, that of bright galaxies increase \citealp{yos06,bou07}) )."
 The inuuber densitv of UVLE in bright part is Κον (o understaid the galaxy. evolution al these recdshilts., The number density of UVLF in bright part is key to understand the galaxy evolution at these redshifts.
 Follow-up optical spectroscopic surveys have also been made (e.g.. Steidelοἱal.1996a.b. 1999)).," Follow-up optical spectroscopic surveys have also been made (e.g., \citealp{ste96a, ste96b,ste99}) )."
 Shaplevetal.(2003) studied. ~300 spectra of z~3 LDGs ancl classified them into + categories according to their rest-Drune Lya equivalent widths (EWs)., \citet{sha03} studied $\sim 800$ spectra of $z \sim 3$ LBGs and classified them into 4 categories according to their rest-frame $\alpha$ equivalent widths (EWs).
 They made a composite spectrum of each category. and found that LDGs with smaller Lya EW tend to show larger EWs of low-ionization interstellar (LIS) absorption lines. larger velocity difference between Lya ancl LIS absorption. aud redder rest-frame UV continua.," They made a composite spectrum of each category and found that LBGs with smaller $\alpha$ EW tend to show larger EWs of low-ionization interstellar (LIS) absorption lines, larger velocity difference between $\alpha$ and LIS absorption, and redder rest-frame UV continua."
 Follow-up spectroscopic observations have been made for LBGs at zZ5 (e.g.. Lehnert&Bremer2003:Stamwvayοἱal.2003.2004.2007:AndoetDow-IIvgelund2005. 2007)).," Follow-up spectroscopic observations have been made for LBGs at $z \gtrsim 5$ (e.g., \citealp{leh03,sta03,sta04,sta07,and04,and07,dow05,dow07}) )."
 As redshift increases. the targets become fainter and the characteristic spectral features move into the wavelength region where night skv emissions are severe. (hus detailed spectroscopic studies of zZ5 LDGs are still not easy.," As redshift increases, the targets become fainter and the characteristic spectral features move into the wavelength region where night sky emissions are severe, thus detailed spectroscopic studies of $z \gtrsim 5$ LBGs are still not easy."
 Hence the sample size of spectroscopically identified, Hence the sample size of spectroscopically identified
account for line blending and population synthesis to define (he dominant contribution to the stellar luminosity.,account for line blending and population synthesis to define the dominant contribution to the stellar luminosity.
 In SB galaxies (he stellar IR continuum. usually dominates over the possible eas and dust emission., In SB galaxies the stellar IR continuum usually dominates over the possible gas and dust emission.
 This represents a major. conceptual simplification in population ancl spectral sviithesis techniques. making possible and easier (he interpretation of integrated spectra from distant stellar clusters and galaxies.," This represents a major, conceptual simplification in population and spectral synthesis techniques, making possible and easier the interpretation of integrated spectra from distant stellar clusters and galaxies."
" since the near IR continuum of SB galaxies is dominated by luminous RSCGs (see2000.lora review).. their integrated spectrum: can be modeled with an equivalent. average star. whose stellar parameters (temperature T,;j. gravitylogg and microturbulence velocity £) mainly depend on the stellar age aud metallicity."," Since the near IR continuum of SB galaxies is dominated by luminous RSGs \citep[see e.g.][for a review]{oo00}, their integrated spectrum can be modeled with an equivalent, average star, whose stellar parameters (temperature $_{eff}$, gravity and microturbulence velocity $\xi$ ) mainly depend on the stellar age and metallicity."
 Both observations and evolutionary models (seee.g.Weller1999:Origliaetal.erence(herein) suggest that RSGs of ages between 6 ancl 100 Myr. ancl metallicities between 1/10 and Solar are characterized bv low gravities (log g«1.0). low temperatures (< 4000 IX) and relatively high microturbulence velocity (£73 km/s).," Both observations and evolutionary models \citep[see e.g.][ and reference therein]{k99,o99,ori03}
 suggest that RSGs of ages between $\simeq $ 6 and 100 Myr and metallicities between 1/10 and Solar are characterized by low gravities (log $<$ 1.0), low temperatures $\le$ 4000 K) and relatively high microturbulence velocity $\xi\ge$ 3 km/s)."
 A variation in the adopted stellar parameters for the average RSC population by AT .g200 Ix. log e=+0.5 and .N£x0.5 kins ! implies a <40.1 dex change in the abundances estimated [rom atomic lines. ancl x+0.2 dex in those estimated from the molecular lines. which are more sensitive to stellar parameters.," A variation in the adopted stellar parameters for the average RSG population by $\Delta $ $_{\rm eff}$$\pm$ 200 K, $\Delta $ log $\pm$ 0.5 and $\Delta \xi$$\mp$ 0.5 km $^{-1}$ implies a $\le \pm 0.1$ dex change in the abundances estimated from atomic lines, and $\le \pm 0.2$ dex in those estimated from the molecular lines, which are more sensitive to stellar parameters."
 L.8j longslit spectra (see Fig. 1)), $\mu$ longslit spectra (see Fig. \ref{IR}) )
 at R=2.500 of the nuclear region of M32 were obtained with the Near IR. Camera Spectrometer (NICS) mounted al (he Italian. Telescopio Nazionale Galileo (ING) on December 2002.," at R=2,500 of the nuclear region of M82 were obtained with the Near IR Camera Spectrometer (NICS) mounted at the Italian Telescopio Nazionale Galileo (TNG) on December 2002."
 The spectra were sky-subtracted. flat-fielded and corrected for atmospheric absorption using an O-s(ar specirunm as reference.," The spectra were sky-subtracted, flat-fielded and corrected for atmospheric absorption using an O-star spectrum as reference."
" Thev were wavelength calibrated by using a Ne-Ar lamp ancl the monocimensional spectra were extracted by stumming over the central 0.5""x3.0"". corresponding to an aperture projected on the source of 9x52 pe at the distance of M32."," They were wavelength calibrated by using a Ne-Ar lamp and the monodimensional spectra were extracted by summing over the central $0.5\arcsec\times 3.0\arcsec$, corresponding to an aperture projected on the source of $9\times 52$ pc at the distance of M82."
 The spectra have been normalized to the continuum. which was determined applving a low-pass smoothing filler to each spectrum.," The spectra have been normalized to the continuum, which was determined applying a low-pass smoothing filter to each spectrum."
 Total integration times of 72 and 32 min in ihe J and IL bands. respectively. were used. »ovidinee a final signalex to noise ratio 240.," Total integration times of 72 and 32 min in the J and H bands, respectively, were used, providing a final signal to noise ratio $\ge$40."
 Dv measuring (he absorption line broadening a stellar velocity dispersion σ 105220 kin/s has been derived. in perfect agreement with the estimate by (1993) and in the tvpical range of values measured in other massive SB galaxies 1999).," By measuring the absorption line broadening a stellar velocity dispersion $\sigma \simeq 105 \pm$ 20 km/s has been derived, in perfect agreement with the estimate by \citet{glt93}
 and in the typical range of values measured in other massive SB galaxies \citep{oli99}."
. A ον of svnthetic spectra of red. supergiant stars for different input. atmospheric parameters and abundances have been computed. using an updated (Origlia version of the code described in Origliaetal. (1993)..," A grid of synthetic spectra of red supergiant stars for different input atmospheric parameters and abundances have been computed, using an updated \citep{o02,o03}
 version of the code described in \citet{o93}. ."
 Briefly. the code uses the," Briefly, the code uses the"
of the molecular gas is too cold to emit in the near-IR.,of the molecular gas is too cold to emit in the near-IR.
" The mass of warm gas seen in the near-IR is several orders of magnitude smaller than that observed in the mid-IR rotational lines by?,, and also much smaller than the virial mass given in Sect. 3.3.."," The mass of warm gas seen in the near-IR is several orders of magnitude smaller than that observed in the mid-IR rotational lines by\citet{brandl09}, and also much smaller than the virial mass given in Sect. \ref{sec:cs}."
" However, even though the mass is small, the energies are very large as we will see in Sect. ??.."," However, even though the mass is small, the energies are very large as we will see in Sect. \ref{sec:lum}. ."
 The ratio Ro_1/1-9= H5» 2-1 S(1)/H5 1-0 S(1) is often used to discriminate between UV heating of H» in photodissociation regions (PDRs) and shock heating., The ratio $R_{2-1/1-0}\equiv$ $_2$ $-$ 1 $_2$ $-$ 0 S(1) is often used to discriminate between UV heating of $_{2}$ in photodissociation regions (PDRs) and shock heating.
" The observed values of the extended emission and compact molecular source are 0.2 and 0.1, respectively."," The observed values of the extended emission and compact molecular source are $0.2$ and $0.1$, respectively."
 We compare these values with the Meudon PDR code described in ?.., We compare these values with the Meudon PDR code described in \citet{lepetit06}. .
 Their Fig., Their Fig.
 27 shows that PDRs typically have R-1/1-0=0.5--0.6., 27 shows that PDRs typically have $R_{2-1/1-0}=0.5-0.6$.
" Values of R-1/1-0 in the range 0.1 to 0.2 are only reached for very high densities and strong radiation fields (n=10° cm-3, y= 10°)."," Values of $R_{2-1/1-0}$ in the range 0.1 to 0.2 are only reached for very high densities and strong radiation fields $n\gtrsim10^{5}$ $^{-3}$, $\chi\gtrsim10^{5}$ )."
" Such conditions do exist in massive star-forming regions, but only close to massive stars."," Such conditions do exist in massive star-forming regions, but only close to massive stars."
" This solution does not apply here because the compact ssource is far away from the SSC, and has no strong ionized line emission indicating the absence of massive young stars."," This solution does not apply here because the compact source is far away from the SSC, and has no strong ionized line emission indicating the absence of massive young stars."
 For the compact source PDR emission is also ruled out by the ratio Hy 1-0 S(0)/H2 1—0 S(1) (seeFig.28of?)..," For the compact source PDR emission is also ruled out by the ratio $_2$ $-$ 0 $_2$ $-$ 0 S(1) \citep[see Fig.~28
of][]{lepetit06}."
 The extended eemission-line region extends over the whole field-of-view and does not peak at the position of the SSC., The extended emission-line region extends over the whole field-of-view and does not peak at the position of the SSC.
 We therefore discard UV heating and favor shocks as main heating mechanism for the extended emission as well., We therefore discard UV heating and favor shocks as main heating mechanism for the extended emission as well.
" Using the models of ?,, we find that J- and C-shocks match the observed values over a wide range of densities (10*—10° cm~?) and shock velocities (15-50 km s!)"," Using the models of \citet{kristensenphd}, we find that J- and C-shocks match the observed values over a wide range of densities $10^{4} - 10^{6}$ $^{-3}$ ) and shock velocities $15-50$ km $^{-1}$ )."
" The ratio Rsqy;ygry= Hz 1-0 S(1)/Bry is an additional, more empirical diagnostics to identify the origin of the H5 excitation."," The ratio $R_\mathrm{S(1)/Br\gamma}\equiv$ $_2$ $-$ 0 $\gamma$ is an additional, more empirical diagnostics to identify the origin of the $_{2}$ excitation."
 These diagnostics provide additional evidence of shock excitation of H» in our data., These diagnostics provide additional evidence of shock excitation of $_{2}$ in our data.
 We compare our results for the Antennae overlap region with data presented by ?.., We compare our results for the Antennae overlap region with data presented by \citet{puxley90}.
" Their data include 44 star-forming regions in 30 galaxies, where, typically, the Hz 1--0 S(1) line is weaker than the Bry line (only 4 regions have a ratio higher than 1)."," Their data include 44 star-forming regions in 30 galaxies, where, typically, the $_2$ $-$ 0 S(1) line is weaker than the $\gamma$ line (only 4 regions have a ratio higher than 1)."
 Their flux ratios range between 0.1 and 1.5 with a mean of 0.5 and a dispersion of 0.3., Their flux ratios range between $0.1$ and $1.5$ with a mean of $0.5$ and a dispersion of $0.3$.
 A few sources such as the merger NGC 6240 show a much higher Rgv1)/pry value., A few sources such as the merger NGC 6240 show a much higher $R_\mathrm{S(1)/Br\gamma}$ value.
" ? present models for a number of different scenarios and conclude that the line emission is most likely from Hu/PDR gas in massive star forming regions, for most of their sources."," \citet{puxley90}
 present models for a number of different scenarios and conclude that the line emission is most likely from /PDR gas in massive star forming regions, for most of their sources."
" With our data, we obtain Rsq1)/pry=0.72 and >100 for the extended emission and the compact source, respectively."," With our data, we obtain $R_\mathrm{S(1)/Br\gamma}=0.72$ and $>100$ for the extended emission and the compact source, respectively."
" The compact source shows a very high ratio, which is far from the typical range in star-forming galaxies."," The compact source shows a very high ratio, which is far from the typical range in star-forming galaxies."
 The ratio of the extended emission is within the range of observed and modeled values by ?.., The ratio of the extended emission is within the range of observed and modeled values by \citet{puxley90}.
" However, this is not true for all of the extended gas."," However, this is not true for all of the extended gas."
" In the region South of the bright llineemission (Fig. 7)),"," In the region South of the bright lineemission (Fig. \ref{fig:cont}) ),"
 we measure aratio of 3., we measure aratio of 3.
 This value suggests that at least parts of the extended H» emission is excited by shocks., This value suggests that at least parts of the extended $_2$ emission is excited by shocks.
" Before discussing the Hz emission as a coolant of the molecular gas in Sect. ??,,"," Before discussing the $_2$ emission as a coolant of the molecular gas in Sect. \ref{sec:large}, ,"
 we need to estimate the H5 , we need to estimate the $_2$ 
Our choice of the noise model for the convereence field relies on previous work.,Our choice of the noise model for the convergence field relies on previous work.
 Van Waerbeke ct al (1999) have shown that the convergence field cau be accurately recoustructed from observed ellipticity data in he absence of systematic errors., Van Waerbeke et al (1999) have shown that the convergence field can be accurately reconstructed from observed ellipticity data in the absence of systematic errors.
 Both the reconstruction schemes of Iadser Squires (1993) and the Maxiuiuna-Likehhood algorithm of Dartelinaun et al (1996) recover he couvergeuce field with adequate accuracy for fields of order a degree ou a side., Both the reconstruction schemes of Kaiser Squires (1993) and the Maximum-Likelihood algorithm of Bartelmann et al (1996) recover the convergence field with adequate accuracy for fields of order a degree on a side.
 Van Waerbeke (1999) ther showed that the noise properties of peaks cau be analytically described using Catussian statistics (Bardecu ct al 1986: Bond Efstathiou 1987). aud the weak. lensing approxiuation.," Van Waerbeke (1999) further showed that the noise properties of peaks can be analytically described using Gaussian statistics (Bardeen et al 1986; Bond Efstathiou 1987), and the weak lensing approximation."
 Figure d. shows a test of the analytical model of Vau Waerbeke (1999) for peaks duc to noise. aud checks the accuracy of the poak distribution iu the reconstructed. &.," Figure \ref{fignoise1} shows a test of the analytical model of Van Waerbeke (1999) for peaks due to noise, and checks the accuracy of the peak distribution in the reconstructed $\kappa$."
 The dot-dashed curve shows the histogram of peaks from this analytical noise model: the double peaked shape is due to peaks with positive curvature aud troughs with negative curvature., The dot-dashed curve shows the histogram of peaks from this analytical noise model; the double peaked shape is due to peaks with positive curvature and troughs with negative curvature.
 Almost overlapping with the analytical curve is the measured histoeraums of peaks in a field with pure noise., Almost overlapping with the analytical curve is the measured histograms of peaks in a field with pure noise.
 The distribution of peak heights measured in lnaps of & recoustructed from noisy ellipticity data (dashed lue) is compared with the distribution iu the # maps of the signal plus Gaussian noise with variance given bv equation 1l (solid Lue)., The distribution of peak heights measured in maps of $\kappa$ reconstructed from noisy ellipticity data (dashed line) is compared with the distribution in the $\kappa$ maps of the signal plus Gaussian noise with variance given by equation \ref{noise} (solid line).
 The close agreeineut of the two curves demonstrates the accuracy of the ij recoustruction sclicme., The close agreement of the two curves demonstrates the accuracy of the $\kappa$ reconstruction scheme.
", The success of the reconstruction gives us confidence im working with the couvergence data aud the noise model directly. avoiding the slow aud expensive recoustruction process,"," The success of the reconstruction gives us confidence in working with the convergence data and the noise model directly, avoiding the slow and expensive reconstruction process."
 We have also verified that usimg the mass aperture, We have also verified that using the mass aperture
an are of an apparent diameter of the microflare.,an arc of an apparent diameter of the microflare.
 The density is estimated from the emission ratio analysis in. the previous section., The density is estimated from the emission ratio analysis in the previous section.
 The initial temperature is set to 2x10° Κ. the temperature of maximum fractional ionization forxv.," The initial temperature is set to $2\times10^6$ K, the temperature of maximum fractional ionization for."
 The cooling time is calculated down to 10° K. The estimated cooling times are presented in Table 2.. which is 9 min for QRI and 8 min for CHI.," The cooling time is calculated down to $10^5$ K. The estimated cooling times are presented in Table \ref{table:bp}, which is 9 min for QR1 and 8 min for CH1."
 Judging from the lightcurves. the observed time delay between the peaks in the X-ray and cool emission lines Is 12 min for QRI and 15 min for CHI (Figs.," Judging from the lightcurves, the observed time delay between the peaks in the X-ray and cool emission lines is 12 min for QR1 and 15 min for CH1 (Figs."
 3. and 4))., \ref{fig:lc01} and \ref{fig:lc08}) ).
 This discrepancy between estimated cooling time and observed delay times is comparable to the crude temporal resolution of the rastering spectrometers. which may have missed the real emission peak.," This discrepancy between estimated cooling time and observed delay times is comparable to the crude temporal resolution of the rastering spectrometers, which may have missed the real emission peak."
 The uncertainty of the loop length estimation could cause error in the cooling time calculation. since the loop was not resolved in XRT images.," The uncertainty of the loop length estimation could cause error in the cooling time calculation, since the loop was not resolved in XRT images."
 If the actual loop was longer than our estimation. the cooling time could be longer as it largely depends on £L (Equation 2)).," If the actual loop was longer than our estimation, the cooling time could be longer as it largely depends on $L$ (Equation \ref{eq:cool}) )."
 As an alternative. the coincidence of independent brightenings in the transition region could be an explanation.," As an alternative, the coincidence of independent brightenings in the transition region could be an explanation."
 But considering the fact that the delayed cool emissions are observed in 7 out of 10 events (Fig 5)). the cool emissions are likely to be related to microflares. rather than coincidence of other brightening events in transition region.," But considering the fact that the delayed cool emissions are observed in 7 out of 10 events (Fig \ref{fig:lc_all}) ), the cool emissions are likely to be related to microflares, rather than coincidence of other brightening events in transition region."
 Another possibility is that the simplified loop cooling model might be inadequate to reproduce the observations., Another possibility is that the simplified loop cooling model might be inadequate to reproduce the observations.
 The estimated cooling times of microflares are much shorter than those of large scale flares., The estimated cooling times of microflares are much shorter than those of large scale flares.
" This is explained by the short conductive cooling time scale of small coronal loops expressed as τω=dx1074,ET> where n,. L. and Ty respectively denote density. loop length. and temperature (?).."," This is explained by the short conductive cooling time scale of small coronal loops expressed as $\tau_{cond} = 4\times10^{-10} n_e L^2 T_0^{-5/2}$ where $n_e$ , $L$, and $T_0$ respectively denote density, loop length, and temperature \citep{cargill1995}."
 Shorter loop lengths result in shorter conductive cooling time. while radiative cooling does not depend on the loop length.," Shorter loop lengths result in shorter conductive cooling time, while radiative cooling does not depend on the loop length."
 The temporal evolution of microflares is traced in quiet regions and in polar coronal holes., The temporal evolution of microflares is traced in quiet regions and in polar coronal holes.
 The lighteurves exhibit impulsive Increases In emission over à wide temperature ranging from toxv., The lightcurves exhibit impulsive increases in emission over a wide temperature ranging from to.
 Hot emission peaks in X-ray.xv. and were followed by enhancements in cool emission in or in for 7 out of 10 microflares.," Hot emission peaks in X-ray, and were followed by enhancements in cool emission in or in for 7 out of 10 microflares."
 The delay time between the peaks in X-ray and cool emission was 8 — 22 min., The delay time between the peaks in X-ray and cool emission was 8 – 22 min.
 Such delayed peaks in cool emission are common to large flares (???).. although the time delays vary from tens of minutes to hours.," Such delayed peaks in cool emission are common to large flares \citep{curdt2004,teriaca2006,raftery2009}, although the time delays vary from tens of minutes to hours."
 The lighteurve behavior is interpreted as the cooling of hot plasma produced by impulsive heating at the beginning of the flare., The lightcurve behavior is interpreted as the cooling of hot plasma produced by impulsive heating at the beginning of the flare.
 To test this hypothesis. observed time delays are compared with the simple loop cooling model by ?..," To test this hypothesis, observed time delays are compared with the simple loop cooling model by \citet{cargill1995}."
 Although the temporal resolution 1s not better than 6 min. theoretical cooling times generally agree with observations.," Although the temporal resolution is not better than 6 min, theoretical cooling times generally agree with observations."
 Therefore. lightcurves of microflares can be explained as impulsive heating followed by the cooling of small coronal loops.," Therefore, lightcurves of microflares can be explained as impulsive heating followed by the cooling of small coronal loops."
 The small loop length of microflare compared to that of active regions results in a shorter conductive cooling time. while the radiative cooling time is independent of the loop length.," The small loop length of microflare compared to that of active regions results in a shorter conductive cooling time, while the radiative cooling time is independent of the loop length."
 The short time scale of microflares is largely due to the small length of the loops., The short time scale of microflares is largely due to the small length of the loops.
 ? investigated the cross-correlations between lighteurves of different emission lines and claimed that and peak emission precedes the peak emission by 5 min. although the correlation coefficients. were low.," \citet{benz1999} investigated the cross-correlations between lightcurves of different emission lines and claimed that and peak emission precedes the peak emission by 5 min, although the correlation coefficients were low."
 They infer that the chromosphere is heated to coronal temperatures by small heating events., They infer that the chromosphere is heated to coronal temperatures by small heating events.
 We did not find noticeable peaks in cool emissions prior to the coronal emission peak., We did not find noticeable peaks in cool emissions prior to the coronal emission peak.
 A possible reason is that the temporal resolution of our observations was insufficient. to resolve rapidly changing phenomena., A possible reason is that the temporal resolution of our observations was insufficient to resolve rapidly changing phenomena.
 Nevertheless. enhanced cool emission was found at the time of peak X-ray flux. which could be a signature of chromospheric evaporatior analogous to large flares.," Nevertheless, enhanced cool emission was found at the time of peak X-ray flux, which could be a signature of chromospheric evaporation analogous to large flares."
 ? also reported that the lighteurve lags the lightcurve by 23 s. which was interpreted as coronal plasma cooling.," \citet{benz1999} also reported that the lightcurve lags the lightcurve by 23 s, which was interpreted as coronal plasma cooling."
 It is interesting to note that lighteurves of microflares in quiet regions and im coronal holes have common characteristics., It is interesting to note that lightcurves of microflares in quiet regions and in coronal holes have common characteristics.
 Delayed cool emission in a bright. point associated with a polar coronal jet was reported (?).., Delayed cool emission in a bright point associated with a polar coronal jet was reported \citep{culhane2007a}.
 In our study. the delayed emission was found both in coronal holes and in quiet regions.," In our study, the delayed emission was found both in coronal holes and in quiet regions."
 It suggeststhat the cooling of microflares, It suggeststhat the cooling of microflares
-.6in,-.6in
They developed a theory that a planet with small semimajor axis may not be formed by a type IL migration because of planet-dise interaction. but through an excitation of the eccentricity because of the Kozai mechanism to values close to e=| such that the pericenter is close to the host star.,"They developed a theory that a planet with small semimajor axis may not be formed by a type II migration because of planet-disc interaction, but through an excitation of the eccentricity because of the Kozai mechanism to values close to $e = 1$ such that the pericenter is close to the host star."
 Then the tidal friction may cireularize the planet's orbit and as a consequence the authors predict that close-in-planets may have large inclinations up to retrograde orbits., Then the tidal friction may circularize the planet's orbit and as a consequence the authors predict that close-in-planets may have large inclinations up to retrograde orbits.
 ? treated the action of the Kozai resonance for five multiplanetary systems which are not in Mean Motion Resonances (=MMR). namely v Andromedae. HD 169830. HD 12661. HD 74156 and HD 155358.," \citet{Lib09b} treated the action of the Kozai resonance for five multiplanetary systems which are not in Mean Motion Resonances (=MMR), namely $\upsilon$ Andromedae, HD 169830, HD 12661, HD 74156 and HD 155358."
 They found out by numerical studies in which they varied the unknown inclinations and also the nodal longitudes. that with the exception of HD 155358 all of them could be in the Kozai resonance when the mutual inclination is larger than 45°.," They found out by numerical studies in which they varied the unknown inclinations and also the nodal longitudes, that with the exception of HD 155358 all of them could be in the Kozai resonance when the mutual inclination is larger than $45^{\circ}$."
 In another investigation by ? the role of mutually inclined orbits with respect to the stability of the orbits have been made in cases when the planets are in MMR., In another investigation by \citet{Lib09a} the role of mutually inclined orbits with respect to the stability of the orbits have been made in cases when the planets are in MMR.
 The authors have undertaken a parametric study varying masses and orbital parameters of the planets where they also took care of the migration rate and the rate of eccentricity damping., The authors have undertaken a parametric study varying masses and orbital parameters of the planets where they also took care of the migration rate and the rate of eccentriciy damping.
 It turned out that due to a capture process in the early phases of the system formation besides the 2/1 MMR (which was treated by ?)) also in the 3/1. 4/1 and 5/1 MMR the mutual inclinations may reach values as high as 70° when the eccentricity of one planet is larger than e>0.4.," It turned out that due to a capture process in the early phases of the system formation besides the 2/1 MMR (which was treated by \citealt{ThomLis}) ) also in the 3/1, 4/1 and 5/1 MMR the mutual inclinations may reach values as high as $70^{\circ}$ when the eccentricity of one planet is larger than $e > 0.4$."
 Due to their results they say that strong., Due to their results they say that .
. The case of TrEs-2b was explored by ?: this is an EPS with a planet with a period P=2.4days., The case of TrEs-2b was explored by \citet{Scu10}; this is an EPS with a planet with a period $P = 2.4 \mbox{ days}$.
 Using older transit observations from 2006 by ? and comparing them to the ones of the Kepler mission recently published by ? they could not find significant changes in the orbital parameters since the discovery of that planet.," Using older transit observations from 2006 by \citet{Odo06} and comparing them to the ones of the Kepler mission recently published by \citet{Gil10}
 they could not find significant changes in the orbital parameters since the discovery of that planet."
 In another system. GJ436. the transiting planet has a period of P=2.6days but a surprisingly large eccentricity of 0.15.," In another system, GJ436, the transiting planet has a period of $P = 2.6 \mbox{ days }$ but a surprisingly large eccentricity of $e = 0.15$ ."
 ? — using secular evolution of a two and also a three planet system describing GJ436 — succeeded in understanding this fact: the circularization has very long times scales when other planets on eccentric orbits are present which may even move in almost the same orbital plane as GJ436b., \citet{Bat09} – using secular evolution of a two and also a three planet system describing GJ436 – succeeded in understanding this fact: the circularization has very long times scales when other planets on eccentric orbits are present which may even move in almost the same orbital plane as GJ436b.
 A different topic in astronomy. namely the observation of period variations in eclipsing binaries due to the presence of a third component. offers possibilities of studying the so-called transit time variations (TTV) also in the analysis of transit photometry of extrasolar-planets (?)..," A different topic in astronomy, namely the observation of period variations in eclipsing binaries due to the presence of a third component, offers possibilities of studying the so-called transit time variations (TTV) also in the analysis of transit photometry of extrasolar-planets \citep{Bor10}."
 But also the duration of an eclipse was studied in the context of EB. and which is the consequence of an influence on the inclination of the two components with respect to the line of sights.," But also the duration of an eclipse was studied in the context of EB, and which is the consequence of an influence on the inclination of the two components with respect to the line of sights."
 The subject of this article. the extrasolar planetary system CoRoT-7 is hosting one planet with the very small period of P=20hours. and at least one other planet.," The subject of this article, the extrasolar planetary system CoRoT-7 is hosting one planet with the very small period of $P = 20 \mbox{ hours }$ and at least one other planet."
 We did numerical investigations for different sets of parameters which may cause a visible change in the transit observations especially with respect to the possible change of the inclination of the innermost planet CoRoT-7b., We did numerical investigations for different sets of parameters which may cause a visible change in the transit observations especially with respect to the possible change of the inclination of the innermost planet CoRoT-7b.
 The EPS CoRoT-7 contains one transiting planet CoRoT-7b (?) and one non transiting but confirmed planet CoRoT-7e detected by radial velocity (2)..," The EPS CoRoT-7 contains one transiting planet CoRoT-7b \citep{Leg09}
 and one non transiting but confirmed planet CoRoT-7c detected by radial velocity \citep{Que09}."
 A reanalysis of RV data led to the assumption that there could be a third planet in this system (?).., A reanalysis of RV data led to the assumption that there could be a third planet in this system \citep{Hat10}.
 CoRoT-7 is — up to now -- a unique example for an EPS where the dynamical time scales are very short which is due to the closeness of the planets to the host star., CoRoT-7 is – up to now – a unique example for an EPS where the dynamical time scales are very short which is due to the closeness of the planets to the host star.
 The short periods of 0.85. 3.7 and 9 days for the three planets mean that the system is under very fast development: an integration time of à million years of CoRoT-7 in a numerical experiment is comparable to an integration of 100 million years for the solar system.," The short periods of 0.85, 3.7 and 9 days for the three planets mean that the system is under very fast development: an integration time of a million years of CoRoT-7 in a numerical experiment is comparable to an integration of 100 million years for the solar system."
 Effects caused by mutual perturbations of the planets will be visible on short time scales., Effects caused by mutual perturbations of the planets will be visible on short time scales.
 We list in Table | the orbital elements of the CoRoT-7 planets. showing the results obtained from two different analyses of the data.," We list in Table \ref{tb:corotelements} the orbital elements of the CoRoT-7 planets, showing the results obtained from two different analyses of the data."
 Former estimations showed that the regions where the planets are moving are very stable (e.g. ?). , Former estimations showed that the regions where the planets are moving are very stable \citep[e.g.][]{Hat10}. .
We show that even inclinations of ος~60° have no big influence on the eccentricity of the planets which would make the system unstable (seeFig. 3))., We show that even inclinations of $i_{7cb} \sim 60^{\circ}$ have no big influence on the eccentricity of the planets which would make the system unstable (seeFig. \ref{fig:31abcd}) ).
 Theeccentricity changes are well inside the probable errors 1n. the derived element, Theeccentricity changes are well inside the probable errors in the derived element
images from the Advanced Camera for Surveys on the (HST;??),images from the Advanced Camera for Surveys on the \citep[\textit{HST;} .
" In a complementary approach, The CAmbridge Sloan Survey Of Wide ARcs in the skY (CASSOWARY) targets multiple, blue companions around massive ellipticals in the SDSS photometric catalogue as likely candidates for wide-separation gravitational lens systems (see?).."," In a complementary approach, The CAmbridge Sloan Survey Of Wide ARcs in the skY (CASSOWARY) targets multiple, blue companions around massive ellipticals in the SDSS photometric catalogue as likely candidates for wide-separation gravitational lens systems \citep[see][]{belokurov09}."
" These and other search strategies have led to numerous recent discoveries of strongly lensed galaxies at z1 which have been followed up in detail over a range of >wavelengths (see2222222222222,andreferences therein).."," These and other search strategies have led to numerous recent discoveries of strongly lensed galaxies at $z \gtrsim 1$ which have been followed up in detail over a range of wavelengths \citep[see][and
 references
 therein]{lemoine03,cabanac05,belokurov07,coppin07,swinbank07,cabanac08,stark08,finkelstein09,hainline09,
 quider09,siana09,dessauges09,quider10}."
" However, all of these studies have of necessity focused only on limited wavelength intervals, mostly in the rest-frame UV (redshifted to optical wavelengths) or the rest frame optical (redshifted into the near-infrared)."," However, all of these studies have of necessity focused only on limited wavelength intervals, mostly in the rest-frame UV (redshifted to optical wavelengths) or the rest frame optical (redshifted into the near-infrared)."
" It is only with the recent advent of the X-shooter spectrograph on the Very Large Telescope (VLT) facility of the European Southern Observatory (ESO), that both wavelength ranges can now be recorded at once with good sensitivity."," It is only with the recent advent of the X-shooter spectrograph on the Very Large Telescope (VLT) facility of the European Southern Observatory (ESO), that both wavelength ranges can now be recorded at once with good sensitivity."
" The obvious advantages have been demonstrated in our recent study of 220 (CSWA20,?),, a blue star-forming galaxy at Zem=1.433 lensed into a wide separation (~ aarcsec) Einstein Cross by a massive luminous red galaxy at zabs6=0.741 with velocity dispersion σιοις500 kkm sl."," The obvious advantages have been demonstrated in our recent study of 20 \citep[CSWA\,20,][]{pettini10}, a blue star-forming galaxy at $z_{\rm
 em} = 1.433$ lensed into a wide separation $\sim 6$ arcsec) Einstein Cross by a massive luminous red galaxy at $z_{\rm abs} =
0.741$ with velocity dispersion $\sigma_{\rm lens} \simeq
500$ km $^{-1}$."
 The curent CASSOWARY includes seven spectroscopically confirmed lensed galaxies at z>1., The current CASSOWARY includes seven spectroscopically confirmed lensed galaxies at $z > 1$.
 In this paper we confirm an eighth case with X-shooter observations of 55., In this paper we confirm an eighth case with X-shooter observations of 5.
" As can be seen from Figure 1,, this multiple system consists of two foreground red galaxies one of which (L2) has an SDSS spectrum which shows it to be at zabs=0.3877, and four fainter sources, separated by ~ 1-5aarcsec."," As can be seen from Figure \ref{fig:SDSS_im}, this multiple system consists of two foreground red galaxies one of which (L2) has an SDSS spectrum which shows it to be at $z_{\rm abs} = 0.3877$, and four fainter sources, separated by $\sim 1$ arcsec."
" Three of these images, iml, im2 and im3, are shown by the observations reported here to be gravitationally lensed images of an emission line galaxy at zo;=1.0686, while the fourth image, not covered by our observations, remains to be confirmed spectroscopically."," Three of these images, im1, im2 and im3, are shown by the observations reported here to be gravitationally lensed images of an emission line galaxy at $z_{\rm em} = 1.0686$, while the fourth image, not covered by our observations, remains to be confirmed spectroscopically."
 Coordinates and magnitudes of the different components of 55 are given in Table 1.., Coordinates and magnitudes of the different components of 5 are given in Table \ref{tab:maggies}.
" However, we caution that the SDSS pipeline magnitudes, which are based on exponential profile fits, are likely to be significantly more uncertain than formal errors quoted, because of blending between the different components of this complex system."," However, we caution that the SDSS pipeline magnitudes, which are based on exponential profile fits, are likely to be significantly more uncertain than formal errors quoted, because of blending between the different components of this complex system."
" In particular, while image im1 may appear to be bluer than im2 from the magnitudes listed in Table 1,, we suspect that this is just due to blending with the foreground red galaxy L1, because our X-shooter spectra of all three images, iml, im2, and im3, exhibit the same continuum slope."," In particular, while image im1 may appear to be bluer than im2 from the magnitudes listed in Table \ref{tab:maggies}, we suspect that this is just due to blending with the foreground red galaxy L1, because our X-shooter spectra of all three images, im1, im2, and im3, exhibit the same continuum slope."
Photometric data are available for 553W091 in R. J. H.& K magnitudes. covering approximately 2000 — 9000 In the rest-frame wavelength at the redshift of this galaxy (5=1.552).,"	Photometric data are available for 53W091 in $R$ ,$J$ , $H$, $K$ magnitudes, covering approximately 2000 – 9000 in the rest-frame wavelength at the redshift of this galaxy $z=1.552$ )."
 Photometric data may be less indicative of the age of a population than the UV spectrum. but it is also less affected by the uncertainties in reddening and in metallicity.," 	Photometric data may be less indicative of the age of a population than the UV spectrum, but it is also less affected by the uncertainties in reddening and in metallicity."
 We have performed a weighted 4 test on the four photometric data points., 	We have performed a weighted $\chi^{2}$ test on the four photometric data points.
 The errors in magnitudes were also adopted from Spinrad et al. (, 	The errors in magnitudes were also adopted from Spinrad et al. (
1997).,1997).
 Figure 11 shows the reduced  as a function of age., 	Figure 11 shows the reduced $\chi^{2}$ as a function of age.
 The models with minimum V indicate ages of «2 Gyr unless the majority of stars in 5533W09] are very metal-poor., 	The models with minimum $\chi^{2}$ indicate ages of $< 2$ Gyr unless the majority of stars in 53W091 are very metal-poor.
 Composite models with OS = 0.2 suggest ages of approximately 1.5 Gyr., 	Composite models with OS = 0.2 suggest ages of approximately 1.5 Gyr.
 Unlike the analysis of the UV spectrum ($5.1). inclusion of OS raises the age estimates based on the photometric data substantially. 20 —50%.," 	 	Unlike the analysis of the UV spectrum 5.1), inclusion of OS raises the age estimates based on the photometric data substantially, 20 –."
.. This is because OS has a large impact on the visible — IR flux when normalized to the UV (see Figure 7)., 	This is because OS has a large impact on the visible – IR flux when normalized to the UV (see Figure 7).
 If no OS is adopted. composite models suggest very small ages (< 1.0 Gyr).," 	If no OS is adopted, composite models suggest very small ages $<$ 1.0 Gyr)."
 The effect of inclusion of metallicity mixtures is also substantial., 	The effect of inclusion of metallicity mixtures is also substantial.
 When the mean metallicity of stars in 553W09] is assumed to be twice solar. our models with OS suggest larger ages when metallicity mixtures are included.," 	When the mean metallicity of stars in 53W091 is assumed to be twice solar, our models with OS suggest larger ages when metallicity mixtures are included."
 This is because composite models contain some metal-poor stars which match the A—J color (1.e.. rest-frame UV continuum) at older ages than metal-rich stars do.," 	This is because composite models contain some metal-poor stars which match the $R-J$ color (i.e., rest-frame UV continuum) at older ages than metal-rich stars do."
 When we use composite models with OS. our analysis on the photometric data of 553W09] suggests 1.5 + 0.2 Gyr. which ts significantly smaller than the age estimate of Spinrad et al.," 	When we use composite models with OS, our analysis on the photometric data of 53W091 suggests 1.5 $\pm$ 0.2 Gyr, which is significantly smaller than the age estimate of Spinrad et al."
 but consistent with that of Bruzual Magris., but consistent with that of Bruzual Magris.
 This value is somewhat smaller than our estimate from the UV analysis in the previous section., 	This value is somewhat smaller than our estimate from the UV analysis in the previous section.
 Table 2 shows the list of our age estimates based on the two data sets., 	Table 2 shows the list of our age estimates based on the two data sets.
 They range approximately | — 2 Gyr unless extreme metallicities are assumed., 	They range approximately 1 – 2 Gyr unless extreme metallicities are assumed.
 Figure 12 shows a sample fit., 	Figure 12 shows a sample fit.
 The observed data are shown as a continuous line and filled circles., 	The observed data are shown as a continuous line and filled circles.
 The overplotted models are those with the infall mixture and OS. which are probably more realistic than single abundance models with no OS.," 	The overplotted models are those with the infall mixture and OS, which are probably more realistic than single abundance models with no OS."
 In the UV. only the model with the minimum  (1.9 Gyr) is shown for clarity.," 	In the UV, only the model with the minimum $\chi^{2}$ (1.9 Gyr) is shown for clarity."
 In the longer wavelength regions. multiple models with various ages are plotted over the photometric data.," 	In the longer wavelength regions, multiple models with various ages are plotted over the photometric data."
 This figure demonstrates the difference between the age estimates from the analyses of the UV spectrum and of the photometric data., 	This figure demonstrates the difference between the age estimates from the analyses of the UV spectrum and of the photometric data.
 This difference between the two age estimates may appear substantial to some readers., 	 	This difference between the two age estimates may appear substantial to some readers.
 However. such a difference would not be unnatural if the data have been affected by some reddening. whether it is Galactic or internal.," 	However, such a difference would not be unnatural if the data have been affected by some reddening, whether it is Galactic or internal."
 If we use the extinction curve of Cardelli. Clayton. Mathis (1989). in conjunction. with that of ODonnell (1994). a moderate amount of Galactic reddening. £(B—V)=0.04. reduces the difference between these two age estimates.," 	If we use the extinction curve of Cardelli, Clayton, Mathis (1989), in conjunction with that of O'Donnell (1994), a moderate amount of Galactic reddening, $E(B-V) = 0.04$ reduces the difference between these two age estimates."
 In fact. the," 	In fact, the"
also (he wavevector range of the distribution and its possible anisotropy. or nonuniformity.,also the wavevector range of the distribution and its possible anisotropy or nonuniformity.
 The first step to study such realistic features was presented. [or mildly relativistic shocks in a recent paper (Niemiec&Ostrowski2004)..! where we considered the power-law wave spectra upstream of the shock within a relatively wide range of wavevectors. wilh only static perturbations involved.," The first step to study such realistic features was presented for mildly relativistic shocks in a recent paper \citep{nie04}, where we considered the power-law wave spectra upstream of the shock within a relatively wide range of wavevectors, with only static perturbations involved."
 The perturbations are imposed on the mean field component. inclined al a given angle to the shock normal.," The perturbations are imposed on the mean field component, inclined at a given angle to the shock normal."
 The downstream magnetic field structure was derived by assuming a simple shock compression of the upstream field. (hus preserving the continuity of the perturbed magnetic field across the shock.," The downstream magnetic field structure was derived by assuming a simple shock compression of the upstream field, thus preserving the continuity of the perturbed magnetic field across the shock."
 By choosing different. aaiplitudes. and spectral indices of the turbulence and different mean field inclinations. we were able to study the acceleration processes in a variely of background. condiüons. including both sub- aud superluminal shocks with varving amount of background perturbations.," By choosing different amplitudes and spectral indices of the turbulence and different mean field inclinations, we were able to study the acceleration processes in a variety of background conditions, including both sub- and superluminal shocks with varying amount of background perturbations."
 The model allowed us also (ο investigate the role in the shock acceleration of short. resonance ancl long waves al dillerent particle energies. to better understand (he microplivsies of the acceleration processes.," The model allowed us also to investigate the role — in the shock acceleration — of short, resonance and long waves at different particle energies, to better understand the microphysics of the acceleration processes."
 The study showed that the particle spectra generally cliverge [rom a simple power law: Che exact shape of the spectrum depends on both the ;nplitude and (he wave power spectrum of the magnetic field perturbations., The study showed that the particle spectra generally diverge from a simple power law; the exact shape of the spectrum depends on both the amplitude and the wave power spectrum of the magnetic field perturbations.
 We reported and discussed such leatures as spectrum hardening before the cutoff at high particle energies in oblique subluminal shocks. and the steepening of the spectrum and (he eutoff formation in (he resonance energy range in superluminal shocks.," We reported and discussed such features as spectrum hardening before the cutoff at high particle energies in oblique subluminal shocks, and the steepening of the spectrum and the cutoff formation in the resonance energy range in superluminal shocks."
 In the latter case. the role of (he long-wave magnetic fiekl perturbations proved to be critical lor the generation of the power-law spectral tails.," In the latter case, the role of the long-wave magnetic field perturbations proved to be critical for the generation of the power-law spectral tails."
 For parallel shocks. we showed that the presence of finite-amplitude magnetic field perturbations leads to the formation of locally oblique field configurations at the shock and the respective magnetic field compressions.," For parallel shocks, we showed that the presence of finite-amplitude magnetic field perturbations leads to the formation of locally oblique field configurations at the shock and the respective magnetic field compressions."
 This implies a «qualitative modification of the particle acceleration process by introducing some features present in oblique shocks., This implies a qualitative modification of the particle acceleration process by introducing some features present in oblique shocks.
 As a result. nonmonotonic variations of the particle spectral index with the turbulence amplitude were observed.," As a result, nonmonotonic variations of the particle spectral index with the turbulence amplitude were observed."
 In the present paper we slightly modilv the above approach to study shocks with larger Lorentz factors 5>>1., In the present paper we slightly modify the above approach to study shocks with larger Lorentz factors $\gamma \gg 1$.
 Besides several detailed results. we report significant difficulties to form. the wide-energv. range power-law particle. spectral (ailsH inH quasi-perpendictular," Besides several detailed results, we report significant difficulties to form the wide-energy range power-law particle spectral tails in quasi-perpendicular"
 Besides several detailed results. we report significant difficulties to form. the wide-energv. range power-law particle. spectral (ailsH inH quasi-perpendictularB," Besides several detailed results, we report significant difficulties to form the wide-energy range power-law particle spectral tails in quasi-perpendicular"
the relerence star.,the reference star.
 The properties of the final eross-correlation function then depend on the properties of the various sources (hat contribute to the total light. their degree of similarity. nol only among themselves but with the reference star. and the fractional contribution that each component makes to the total light.," The properties of the final cross-correlation function then depend on the properties of the various sources that contribute to the total light, their degree of similarity, not only among themselves but with the reference star, and the fractional contribution that each component makes to the total light."
 In principle. a threshold value of the measured velocity dispersion Chat corresponds {ο some pre-defined value of f can be determined from simulations. in which a spectrum corresponding to that of the dominant lisht source are. added in various proportions to a spectrum representing (he galaxv body that has been smeared (o simulate the impact of velocity dispersion.," 	In principle, a threshold value of the measured velocity dispersion that corresponds to some pre-defined value of $f$ can be determined from simulations, in which a spectrum corresponding to that of the dominant light source are added in various proportions to a spectrum representing the galaxy body that has been smeared to simulate the impact of velocity dispersion."
 The situation is simplified somewhat in near-inlrared studies of intermediate age or older svstems. as a large fraction of the A—band light comes from the brightest evolved stars.," The situation is simplified somewhat in near-infrared studies of intermediate age or older systems, as a large fraction of the $K-$ band light comes from the brightest evolved stars."
 A consequence is that the spectra of the brightest resolved stars are similar to that of the underlving integrated light. making it is easier {ο match the reference star spectrum with those of the integrated galaxy light and any resolved stars.," A consequence is that the spectra of the brightest resolved stars are similar to that of the underlying integrated light, making it is easier to match the reference star spectrum with those of the integrated galaxy light and any resolved stars."
 A basic assumption is that the resolution elements (hat are dominated by individual stars have an underlving stellar component with a σ (hat surpasses (hat which can be measured Irom the data., 	A basic assumption is that the resolution elements that are dominated by individual stars have an underlying stellar component with a $\sigma$ that surpasses that which can be measured from the data.
 Higher spectral resolutions will be required to detect single sources in fields with lower co., Higher spectral resolutions will be required to detect single sources in fields with lower $\sigma$.
 It will thus prove more challenging to identilv locations that are dominated by light from a single source in erowded enviroments where the stars have modest (at least when compared with the centers of galaxy bulges) random motions. such as globular clusters.," It will thus prove more challenging to identify locations that are dominated by light from a single source in crowded enviroments where the stars have modest (at least when compared with the centers of galaxy bulges) random motions, such as globular clusters."
 The velocity dispersions of the sources are listed in the last column of Table 1., 	The velocity dispersions of the sources are listed in the last column of Table 1.
 The AIOLLI star III. 4884. which was also observed with NIES. was used as (he reference template for these measurements.," The M0III star HR 4884, which was also observed with NIFS, was used as the reference template for these measurements."
 7 was measured in (he 2.1— 2.3) interval. which is where the S/N ralio is highest.," $\sigma$ was measured in the $2.1 - 2.3\mu$ m interval, which is where the S/N ratio is highest."
 The uncertainties in the σ measurements are <+10 km !., The uncertainties in the $\sigma$ measurements are $\leq \pm 10$ km $^{-1}$.
 The spectra of sources 10 and 14 were too ποῖδον to allow velocity. dispersious to be measured., The spectra of sources 10 and 14 were too noisey to allow velocity dispersions to be measured.
 There is a broad range in velocity dispersions in Table 1. ancl half of the objects have c>90 km |.," 	There is a broad range in velocity dispersions in Table 1, and half of the objects have $\sigma > 90$ km $^{-1}$."
 Given that the velocity dispersion of M32. within the central 1.5 arcsec is >100 km +. then the majority of locations in the M32 NLIFS spectra that contain a distinct point source are actually blends where the light originates [rom a mix of stars. with no single star dominating the light output.," Given that the velocity dispersion of M32 within the central 1.5 arcsec is $\geq 100$ km $^{-1}$, then the majority of locations in the M32 NIFS spectra that contain a distinct point source are actually blends where the light originates from a mix of stars, with no single star dominating the light output."
 The appearance of a point source al these locations is simply the result of statistical fIucuations in the stellar content., The appearance of a point source at these locations is simply the result of statistical flucuations in the stellar content.
D? where (pj is the comoving mean density and where Pry is the gas pressure aud e dis the eas internal energy vr tit) volume.,where where $\langle\rho\rangle$ is the comoving mean density and where $p_{\rm th}$ is the gas pressure and $\epsilon$ is the gas internal energy per unit volume.
 We assume the aciabatic index ~=5/3 and the mean molecular weight µ=0.5 iu he equation of state., We assume the adiabatic index $\gamma = 5/3$ and the mean molecular weight $\mu = 0.5$ in the equation of state.
 Tn equation (3). C represcuts he cooling rate per unit volume.," In equation (3), ${\cal C}$ represents the cooling rate per unit volume."
 We use stancard abulated aud publicly available cooling curves. 9) or metallieitv Z=0.3Z.., We use standard tabulated and publicly available cooling curves \citep{sutherland93} for metallicity $Z = 0.3Z_{\odot}$.
 Our snuulatious co not iuchude star formation or ACN feedback aud we set a Hoor in the cooling function at 0.01 keV in physical The anisotropic thermal couduction leat fix Fy is eiven bv where ep is a unit vector poiutiusg iu the direction of the magnetic field and & is the Spitzer-Draginskii conduction coefficient eiven by &(P)=L6G«10Το Pore tam !K +., Our simulations do not include star formation or AGN feedback and we set a floor in the cooling function at 0.01 keV in physical The anisotropic thermal conduction heat flux ${\bf F}_{\rm ph}$ is given by where $\hat{\bf e}_{B}$ is a unit vector pointing in the direction of the magnetic field and $\kappa$ is the Spitzer-Braginskii conduction coefficient given by $\kappa (T) = 4.6\times 10^{-7}T^{5/2}$ erg $^{-1}$ $^{-1}$ $^{-1}$.
 In the above equation aud in all equations below all variables with the “pho subscript denote physical quantities., In the above equation and in all equations below all variables with the “ph” subscript denote physical quantities.
 Following ? we included the effect of conduction saturation whenever the characteristic Ienethscale associate with the temperature eradicut exceeds the mean frec| path. though in the bull of the ICAL this effect is no senificaut.," Following \citet{cowie77} we included the effect of conduction saturation whenever the characteristic lengthscale associated with the temperature gradient exceeds the mean free path, though in the bulk of the ICM this effect is not significant."
 We also iuposed an upper Tauit ou conduction suchthat AK.—aορ(ορ). where et=(L)iuiiworfMoire (the extra factor of « comes from) the prefactor im Equation 8).," We also imposed an upper limit on conduction suchthat $K\equiv a^{-6}\kappa /(c_{v}\rho)$, where $c_{v}^{-1} = (\gamma -1)\mu m_{\rm prot}/k_{\rm boltz}$ (the extra factor of $a$ comes from the prefactor in Equation 8)."
 We set AcRua5s qUT7ensgs +., We set $K<K_{\max} = 5\times 10^{32}$ $^{2}$ $^{-1}$.
 This ceiling was introduced in order to prevent extremely small diffusive timesteps., This ceiling was introduced in order to prevent extremely small diffusive timesteps.
" We found that changing A,,44 to larger values did not affect our results significantly,", We found that changing $K_{\rm max}$ to larger values did not affect our results significantly.
 Only a relatively lnnmited volune far from the cluster center was subject to this upper We also note that the effective couduction may be lamited in the outer regions of the cluster for plivsical reasous (2)., Only a relatively limited volume far from the cluster center was subject to this upper We also note that the effective conduction may be limited in the outer regions of the cluster for physical reasons \citep{medvedev07}.
 Moreover. around or bevoud the virial radius. the clectrou-proton energy equilibratiou nuescales start to exceed the dvuamical or buovaucv inescale.," Moreover, around or beyond the virial radius, the electron-proton energy equilibration timescales start to exceed the dynamical or buoyancy timescale."
 Iu this region. the MIID approximation starts ο break down aud the AITT erowth rates are reduced.," In this region, the MHD approximation starts to break down and the MTI growth rates are reduced."
 However. at sinaller radii. the equilibration timescale due o Coulomb collisions are comparable to or shorter than he buovancy timescale. some thermal coupling between electrons aud protous begins to be possible. aud the AIIID analysis applies at least approximately.," However, at smaller radii, the equilibration timescale due to Coulomb collisions are comparable to or shorter than the buoyancy timescale, some thermal coupling between electrons and protons begins to be possible, and the MHD analysis applies at least approximately."
 Finally. here are independent observational arguiucuts for fast equilibration in the ICAL," Finally, there are independent observational arguments for fast equilibration in the ICM."
 Specifically. ? detected a sharp electron tempcrature jump in the post shock reeion in the Bullet Cluster aud used this observation o sugeests that the electron-proton equilibration is uuch shorter than the collisional timescale.," Specifically, \citet{markevitch07} detected a sharp electron temperature jump in the post shock region in the Bullet Cluster and used this observation to suggests that the electron-proton equilibration is much shorter than the collisional timescale."
 The inferred equilibration timescale iu the postslocl ICM is at least 5 times shorter than the Coulomb equilibration time aud nay even be consistent with being instantancous., The inferred equilibration timescale in the postshock ICM is at least 5 times shorter than the Coulomb equilibration time and may even be consistent with being instantaneous.
 While he issue of equilibration is clearly an open research opic. here we work under the svuiplifvine assuniptiou hat equilibration is iustantancous aud that the MIID approxinationu is appropriate. and then investigate the consequences of these The most stringent Blnitation on the timestep df iu the simulation was due to the diffusive process with dt~(Av)?/W.," While the issue of equilibration is clearly an open research topic, here we work under the symplifying assumption that equilibration is instantaneous and that the MHD approximation is appropriate, and then investigate the consequences of these The most stringent limitation on the timestep $dt$ in the simulation was due to the diffusive process with $dt\sim (\Delta x)^{2}/K$."
 The maxima resolution that we could afford computationally was 7 levels of refinement for blocks consisting of 16 zones on a side., The maximum resolution that we could afford computationally was 7 levels of refinement for blocks consisting of 16 zones on a side.
 That is. the effective resolution was ~315 tkpe.," That is, the effective resolution was $\sim 31h^{-1}$ kpc."
 The strong limit iuposed on the timestep can be avoided by cuploving iuplicit iutegratioun methods., The strong limit imposed on the timestep can be avoided by employing implicit integration methods.
 We are now mipleiieutiug such inethods in theFLASII code (Lee et al.," We are now implementing such methods in the code (Lee et al.,"
 iu prep.)., in prep.).
 This approach will siguificautly speed up the conrputations and will allow us to perform simulations for a range of cluster masses (Ruszkowski et al..," This approach will significantly speed up the computations and will allow us to perform simulations for a range of cluster masses (Ruszkowski et al.,"
 in Equations 1 through 8 were obtained by starting from the MIID equations aud applving cosmological expansion transformation for the spatial eracdicuts aud the following variable transformations where 6 is the cosnological expansiou factor., in Equations 1 through 8 were obtained by starting from the MHD equations and applying cosmological expansion transformation for the spatial gradients and the following variable transformations where $a$ is the cosmological expansion factor.
 The velocity vector v= x. where x=elr is the code position vector and r is the physical position vector.," The velocity vector ${\bf v}=\dot{{\bf x}}$ , where ${\bf x}=a^{-1}{\bf r}$ is the code position vector and ${\bf r}$ is the physical position vector."
find still somewhat higher densities for the hot gas. up to 10 em.,"find still somewhat higher densities for the hot gas, up to $10^7$ $^{-3}$."
 The hot ensemble mass is close to the 170 M. derived from early CO 7-6 observations by Jaffeetal.(1989)., The hot ensemble mass is close to the 170 $M_\odot$ derived from early CO 7-6 observations by \citet{Jaffe1989}.
. The total mass of the PDR ensembles falls between the mass limits derived by Jakobetal.(2007) from dust observations and from line radiative transfer fits.," The total mass of the PDR ensembles falls between the mass limits derived by \citet{Jakob}
 from dust observations and from line radiative transfer fits."
 While the existing ground-based observations provide a very good constraint on the properties of the extended cool gas. and the ISO lines show the total amount of hot gas. it is only the set of new HIFI data that puts the hot and cold distributions well apart from each other in terms of the temperature structure.," While the existing ground-based observations provide a very good constraint on the properties of the extended cool gas, and the ISO lines show the total amount of hot gas, it is only the set of new HIFI data that puts the hot and cold distributions well apart from each other in terms of the temperature structure."
 While Jakobetal.(2007) obtained cooling curves with single peaks. the new data for the 10-9 lines of the CO isotopes and the HCO™ transitions force the fit to a bimodal distribution of excitation conditions.," While \citet{Jakob} obtained cooling curves with single peaks, the new data for the 10-9 lines of the CO isotopes and the $^+$ transitions force the fit to a bimodal distribution of excitation conditions."
 When we exclude the data from the model fit. we obtain a parameter set that shows a UV field that is lower by a factor ten for the hot ensemble. i.e.. that would imply molecular clumps farther away from the central cluster.," When we exclude the data from the model fit, we obtain a parameter set that shows a UV field that is lower by a factor ten for the hot ensemble, i.e., that would imply molecular clumps farther away from the central cluster."
 For the cold ensemble. the fitted UV field 1s also somewhat lower. while all other parameters remain similar to those from the full fit.," For the cold ensemble, the fitted UV field is also somewhat lower, while all other parameters remain similar to those from the full fit."
" Only with theHersche/ data. we therefore obtain a parameter set that is consistent with the source geometry,"," Only with the data, we therefore obtain a parameter set that is consistent with the source geometry."
 As the two-ensemble PDR model is able to fit all of the observed lines. we find no evidence for a shock heating of the dense gas.," As the two-ensemble PDR model is able to fit all of the observed lines, we find no evidence for a shock heating of the dense gas."
 This is in agreement with the analysis of Laneetal.(1990).. explicitely excluding a shock origin of the fine-structure lines. but seems to be in contradiction with the analysis of the line profiles in Sect.," This is in agreement with the analysis of \citet{Lane}, explicitely excluding a shock origin of the fine-structure lines, but seems to be in contradiction with the analysis of the line profiles in Sect."
 3. that shows excited outflow material., \ref{sect_profiles} that shows excited outflow material.
 We conclude that the material visible in the blue line wing. characterizing the blister outflow. is contained in dense clumps that are accelerated by the outflow. but that are chemically and energetically fully dominated by the UV field and not by the associated shock.," We conclude that the material visible in the blue line wing, characterizing the blister outflow, is contained in dense clumps that are accelerated by the outflow, but that are chemically and energetically fully dominated by the UV field and not by the associated shock."
than Class II sources (CTTSs).,than Class II sources (CTTSs).
 Since flat-spectrum protostars are believed to be evolutionary precursors of CTTSs. then this finding may indicate that certain pliysical conditions characteristic oL protostellar evolution (e.g.. high accretion rates) may result in their higher rotation rates.," Since flat-spectrum protostars are believed to be evolutionary precursors of CTTSs, then this finding may indicate that certain physical conditions characteristic of protostellar evolution (e.g., high accretion rates) may result in their higher rotation rates."
 Lt is therefore interesting to ask whether tlie less evolved Class [sources might rotate even more rapidly than the flat spectrum sources., It is therefore interesting to ask whether the less evolved Class I sources might rotate even more rapidly than the flat spectrum sources.
 Indeed. some models of protostar development prediet that such objects should be rotating near breakup (Shu1991).," Indeed, some models of protostar development predict that such objects should be rotating near breakup \citep{S91}."
. Recently. strong hard X-ray. flares have been observed with the ASCA satellite from two of ——e Class I protostars in our sample — IRS 13 aud WL 6 (Montinerleetal.)).," Recently, strong hard X-ray flares have been observed with the ASCA satellite from two of the Class I protostars in our sample – IRS 43 and WL 6 \citeauthor{MGTK00}) )."
 These remarkable ¢bservatious reveal relatively rapid periodicities in the X-ray. emission [rom these sources. enabling he derivation of their photometric rotation rates.," These remarkable observations reveal relatively rapid periodicities in the X-ray emission from these sources, enabling the derivation of their photometric rotation rates."
 Montimerleetal. find that WL 6 is rotating with period of about 3 days (e sin £c10 kins + for a 0.5 M. star on the birthline). comparable to he rotation rates of the flat-spectrum sources observed here aud in Paper II.," \citeauthor{MGTK00} find that WL 6 is rotating with a period of about 3 days $v$ sin $i \simeq 40$ km $^{-1}$ for a 0.5 $_{\sun}$ star on the birthline), comparable to the rotation rates of the flat-spectrum sources observed here and in Paper II."
 This source also has weak outflow (Sekimotoetal.1997).. undetectecl millimeter emission from its envelope ontinerle199 [).. aud au IR euergy distribution which cau be moclelecd as a highly extiuguished lat spectrum YSO (Montinerleet al.)).," This source also has a weak outflow \citep{STUKTH97}, undetected millimeter emission from its envelope \citep{AM94}, and an IR energy distribution which can be modeled as a highly extinguished flat spectrum YSO \citeauthor{MGTK00}) )."
 All of these properties iudicate that WL 6 may indeed be very slinilar to the flat spectrum YSOs for which we lave detected absorption lines aud have found o be rotating more rapidly than CTTSs (Class Il YSOx)., All of these properties indicate that WL 6 may indeed be very similar to the flat spectrum YSOs for which we have detected absorption lines and have found to be rotating more rapidly than CTTSs (Class II YSOs).
 Montimerleοἱal. also argue that the central star of IRS 13 has a 20 hi rotation period. the observed period of its X-ray. variability.," \citeauthor{MGTK00} also argue that the central star of IRS 43 has a 20 h rotation period, the observed period of its X-ray variability."
 This requires that its 1uass be greater tliau or equal to 1.5 AL. in order for it to be rotating below breakup velocity if it is ou tlie birthline (Montinerleet al.))., This requires that its mass be greater than or equal to 1.8 $_{\odot}$ in order for it to be rotating below breakup velocity if it is on the birthline \citeauthor{MGTK00}) ).
 Thus the mass of IRS 13 is constrained to lie iu the range 1.5 — 2.2 M. by its X-ray. emission aud bolometric luminosity (see § 1.2) if its rotation period is indeed equal to its X-ray. variability period of 20 lh. IRS 13 does have a more steep iuid-IB. energy clistributiou (clearly Class I) than WL 6. aud it also has spatially-resolved (r  3000 AU) millimeter emission with a derived envelope mass of approximately 0.1 NM... (André& Moutmerle)).," Thus the mass of IRS 43 is constrained to lie in the range 1.8 – 2.2 $_{\odot}$ by its X-ray emission and bolometric luminosity (see $\S 4.2$ ) if its rotation period is indeed equal to its X-ray variability period of 20 h. IRS 43 does have a more steep mid-IR energy distribution (clearly Class I) than WL 6, and it also has spatially-resolved (r $\simeq$ 3000 AU) millimeter emission with a derived envelope mass of approximately 0.1 $_{\odot}$ \citeauthor{AM94}) )."
 Therefore it is likely to be in an earlier evolutionary state than WL 6., Therefore it is likely to be in an earlier evolutionary state than WL 6.
 The slow rotation velocities of C'TTSs liave been explained by augular iuoumentum regulation ol these stars by magnetic coupling to their disks., The slow rotation velocities of CTTSs have been explained by angular momentum regulation of these stars by magnetic coupling to their disks.
 Edwardsetal.(1993) found that late-type T Tawi stars (TTSs) with laree H—A IR excesses (CTTSs) bac slow rotation periods. P > | d. They also fouud that TTS with simall A—4 IR excesses had a broad range of periods. including a significant number with P « Idays.," \citet{ESHSHHAMPG93} found that late-type T Tauri stars (TTSs) with large $H-K$ IR excesses (CTTSs) had slow rotation periods, P $>$ 4 d. They also found that TTS with small $H-K$ IR excesses had a broad range of periods, including a significant number with P $<$ 4 days."
 Edwardsetal. interpreted this correlation to arise because the magnetic fields of the CTTSs were were coupled to their disks. providing stellar angular momentuui regulation and therefore loiD>ο stellar rotation periods.," \citeauthor{ESHSHHAMPG93} interpreted this correlation to arise because the magnetic fields of the CTTSs were were coupled to their disks, providing stellar angular momentum regulation and therefore long stellar rotation periods."
 The low-excess TTSs had already cissipated their disks ancl so were uot subject to this regulation mechanism., The low-excess TTSs had already dissipated their disks and so were not subject to this regulation mechanism.
 More recent studies of larger TTS samples have both disputed that the correlation between IR excess aud rotation period exists (Stassunetal.1999). aud have provided evidence that it exists but is weak (Herbstetal.2000)., More recent studies of larger TTS samples have both disputed that the correlation between IR excess and rotation period exists \citep{SMMV99} and have provided evidence that it exists but is weak \citep{HRHC00}.
. Our studies of YSO rotation (this paper and Paper II) have shown that flat-spectruum protostars rotate significantly more rapidly than Class II YSOs or CTTSs. suggestingMOD that rotation velocities," Our studies of YSO rotation (this paper and Paper II) have shown that flat-spectrum protostars rotate significantly more rapidly than Class II YSOs or CTTSs, suggesting that rotation velocities"
Gamma Rav Bursts (GRBs). the most. powerful cosmic explosions. are very important as probes of the high-redshift universe.,"Gamma Ray Bursts (GRBs), the most powerful cosmic explosions, are very important as probes of the high-redshift universe."
 This is clue to several factors: (4) GRBs occur at cosmological distances. the most distant so far being at 2—4.5 (GARB 000131. Andersen ct al.," This is due to several factors: (i) GRBs occur at cosmological distances, the most distant so far being at z=4.5 (GRB 000131, Andersen et al."
 2000). (i) they can be extremely bright and therefore could be seen (though for a very short. period. of time) up to redshift 20 (e.g. Wijers et al.," 2000), (ii) they can be extremely bright and therefore could be seen (though for a very short period of time) up to redshift 20 (e.g. Wijers et al."
 1998: Lamb Reichart 2000). (ii) the burst. and its afterglow illuminate the material situated between the GRB and the observer. providing an opportunity to study ealaxies that otherwise would be too faint to be observed spectroscopicallv.," 1998; Lamb Reichart 2000), (iii) the burst and its afterglow illuminate the material situated between the GRB and the observer, providing an opportunity to study galaxies that otherwise would be too faint to be observed spectroscopically."
 This is particularly relevant for the host galaxies of GRBs where. in analogy to quasar (QSO) absorption line studies. the strengths of the absorption lines can be used to measure important properties. such as their gas content. dynamical structure. anc metallicity.," This is particularly relevant for the host galaxies of GRBs where, in analogy to quasar (QSO) absorption line studies, the strengths of the absorption lines can be used to measure important properties, such as their gas content, dynamical structure, and metallicity."
" In addition. since they are transient. phenomena. they do not cause the large scale ""proximity elfect of quasars. ie. they do not modify their environment at [large distances. although they are expected to do so in their immediate. vicinity (see section 5)). (iv) the currently most. popular physical explanation of this phenomenon is the. so-called. fireball model (ef"," In addition, since they are transient phenomena, they do not cause the large scale “proximity effect” of quasars, i.e. they do not modify their environment at large distances, although they are expected to do so in their immediate vicinity (see section \ref{sec:discussion}) ), (iv) the currently most popular physical explanation of this phenomenon is the so-called fireball model (cf."
 Rees Aleszaros L992) which is a relativistic explosion caused. by either the collapse of a massive star (Woosley 1993: Paczvnski 1998) or the merging of two compact objects (Lattimer Schramm 1974)., Rees Meszaros 1992) which is a relativistic explosion caused by either the collapse of a massive star (Woosley 1993; Paczynski 1998) or the merging of two compact objects (Lattimer Schramm 1974).
 In the first case the rate of GRBs in the universe should. be directly linked to the instantaneous star formation rate (SER)., In the first case the rate of GRBs in the universe should be directly linked to the instantaneous star formation rate (SFR).
 This indeed seems to be supported by the observations. since the host galaxies of are undergoing strong star-formation activity (c.g. Fruchter et al.," This indeed seems to be supported by the observations, since the host galaxies of are undergoing strong star-formation activity (e.g. Fruchter et al."
 2001: Vreeswijk et al., 2001; Vreeswijk et al.
 2001)., 2001).
 Pherefore. GRBs," Therefore, GRBs"
from a V-band image of MI2 taken with the ESO 2.2m telescope Wide Field Imager (WFI) on 2002 June 20 for 100 seconds.,from a V-band image of M12 taken with the ESO 2.2m telescope Wide Field Imager (WFI) on 2002 June 20 for 100 seconds.
 The WFI has 2 « 4 CCDs array. each CCD has 8” « 16’ field of view. resulting in a total 33 « 34’ field of view.," The WFI has 2 $\times$ 4 CCDs array, each CCD has $\arcmin$ $\times$ $\arcmin$ field of view, resulting in a total $\arcmin$ $\times$ $\arcmin$ field of view."
 MI2 was on one of the 8 chips (chip 7) with the whole area covering the half-mass radius., M12 was on one of the 8 chips (chip 7) with the whole area covering the half-mass radius.
 We first did the bias and flat-fielc calibrations. and then performed the astrometry on the image extracted from that chip only.," We first did the bias and flat-field calibrations, and then performed the astrometry on the image extracted from that chip only."
 We chose 173 stars on the WFI image that matched the USNO CCD Astrograph Catalog (UCAC2) standards (Zacharias et al., We chose 173 stars on the WFI image that matched the USNO CCD Astrograph Catalog (UCAC2) standards (Zacharias et al.
 2004) and then fitted them with a six-parameter transformation., 2004) and then fitted them with a six-parameter transformation.
 As a result. we obtained a transformatioi solution with residual errors of 07087 in R.A. and 07093 i1 Decl.," As a result, we obtained a transformation solution with residual errors of $\farcs$ 087 in R.A. and $\farcs$ 093 in Decl."
 We then chose 317 stars on the AACS drizzled F625W image and matched them with the WFI image., We then chose 317 stars on the ACS drizzled F625W image and matched them with the WFI image.
 The resultant residual error of the transformatior solution was 07036 in R.A. and 07021 in Decl., The resultant residual error of the transformation solution was $\farcs$ 036 in R.A. and $\farcs$ 021 in Decl.
 At this stage. we have aligned the WFI image and the AACS image together and placed them both on the absolute ICRS frame.," At this stage, we have aligned the WFI image and the ACS image together and placed them both on the absolute ICRS frame."
 For the image. we first improved its absolute astrometry with the Aspect privided by the XX-ray Center.," For the image, we first improved its absolute astrometry with the Aspect privided by the X-ray Center."
 lt provides an absolute astrometry of 076 (909c))., It provides an absolute astrometry of $\farcs$ 6 ).
 The calculated offset is -0705 in R.A. and -0709 in Dec. We found CX2. CX4. CX6. CXII and CXI2 coincided with possible counterparts in the WFI image when we overlaid the confidence error circles of ssources on the WFI image.," The calculated offset is $\farcs$ 05 in R.A. and $\farcs$ 09 in Dec. We found CX2, CX4, CX6, CX11 and CX12 coincided with possible counterparts in the WFI image when we overlaid the confidence error circles of sources on the WFI image."
 CX2 has 3 ooptical counterparts. and in the WFI image the three counterparts are not resolved.," CX2 has 3 optical counterparts, and in the WFI image the three counterparts are not resolved."
 CX4 has 2 optical counterparts on the limage., CX4 has 2 optical counterparts on the image.
 Since we can not be sure which source will be the correct counterpart. we did not take CX2 and CX4 as astrometry references.," Since we can not be sure which source will be the correct counterpart, we did not take CX2 and CX4 as astrometry references."
 In addition. ΝΤ corresponds to a non-point-like source while CX6 corresponds to a point-like source.," In addition, CX11 corresponds to a non-point-like source while CX6 corresponds to a point-like source."
 Although CX6 corresponds to a WFI source just beside a saturated star. the WEI counterpart is bright enough and could be distinguished from the saturated star. and in addition. could be located precisely.," Although CX6 corresponds to a WFI source just beside a saturated star, the WFI counterpart is bright enough and could be distinguished from the saturated star, and in addition, could be located precisely."
 Based on its brightness and position accuracy. we took it as an astrometry reference.," Based on its brightness and position accuracy, we took it as an astrometry reference."
 The optical counterpart of CX12 on the WFI image is also a point-like source and outside the half-mass radius of M12., The optical counterpart of CX12 on the WFI image is also a point-like source and outside the half-mass radius of M12.
 The stellar density outside the half-mass radius is so low that we consider the identification secure., The stellar density outside the half-mass radius is so low that we consider the identification secure.
 As a consequence. we used the optical counterparts of CX6 and CX12 to perform boresight correction on image.," As a consequence, we used the optical counterparts of CX6 and CX12 to perform boresight correction on image."
 The position offset needed to apply on the ssources is 07421 + 07040 in R.A. and -07042 + 07004 in Decl., The position offset needed to apply on the sources is $\farcs$ 421 $\pm$ $\farcs$ 040 in R.A. and $\farcs$ 042 $\pm$ $\farcs$ 004 in Decl.
 After the boresight correction. the WFI and the image are on the same frame.," After the boresight correction, the WFI and the image are on the same frame."
 With the X-ray and optical astrometry. we obtained a final uncertainties of the X-ray source position in R.A. and in Decl.," With the X-ray and optical astrometry, we obtained a final uncertainties of the X-ray source position in R.A. and in Decl.,"
 which are the root of the quadratic sum of the position uncertainty of the X-ray sources (given in Table 1). the uncertainty of the absolute astrometry in the WEI. the uncertainty of the astrometry in the WFI and limage alignment. and the uncertainty of the X-ray boresight correction.," which are the root of the quadratic sum of the position uncertainty of the X-ray sources (given in Table 1), the uncertainty of the absolute astrometry in the WFI, the uncertainty of the astrometry in the WFI and image alignment, and the uncertainty of the X-ray boresight correction."
 The final Io position error for all X-ray sources, The final $\sigma$ position error for all X-ray sources
" ""wfesur/cced.html) on the 2.5 m Isaac Newton Telescope on La Palma.",$^{\sim}$ wfcsur/ccd.html) on the 2.5 m Isaac Newton Telescope on La Palma.
 AX total of 113 fields in the Virgo Cluster were selected (sce Lig., A total of 113 fields in the Virgo Cluster were selected (see Fig.
 l1) comprising 24.9 deg? taking into account overlaps between the images. which were typically 1: chip or 1/4 o£ the field of view.," 1), comprising 24.9 $^2$ taking into account overlaps between the images, which were typically 1 chip or 1/4 of the field of view."
 ALL fields were imaged through a D (ellective wavelength Ay=0.44 jum) filter for 750 s under photometric conditions., All fields were imaged through a $B$ (effective wavelength $\lambda_0 = 0.44$ $\mu$ m) filter for 750 s under photometric conditions.
 A total of 31 deg? from the WES imaging survey of the North Galactic Cap (NGC). taken in photometric conditions with a £ filter. were also studied. these being used as olfset fields to estimate the background contamination in the Virgo sample.," A total of 31 $^2$ from the WFS imaging survey of the North Galactic Cap (NGC), taken in photometric conditions with a $B$ filter, were also studied, these being used as offset fields to estimate the background contamination in the Virgo sample."
 The exposure times in D (150 s) were the same for the Virgo and NGC data., The exposure times in $B$ (750 s) were the same for the Virgo and NGC data.
 The data were pre-processecl αμα reduced via he WES pipeline wlesur pipeline.html: Irwin. Lewis 2001) and photometric calibration obtained from observations of several (5 10 »er night) stanclare stars., The data were pre-processed and reduced via the WFS pipeline $^{\sim}$ wfcsur/ pipeline.html; Irwin Lewis 2001) and photometric calibration obtained from observations of several (5 – 10 per night) standard stars.
 Ehe photometric zero points were always accurate to 2 uncertainties in the zero point are not a major source of uncertainty in the galaxy magnitudes hat we derive., The photometric zero points were always accurate to 2; uncertainties in the zero point are not a major source of uncertainty in the galaxy magnitudes that we derive.
 The median seeing was 1.95 areseconds for he Virgo data and 1.26 arcseconds for the NGC data (since we will be searching for galaxies with large sizes and low surface-brightnesses. the relatively poor seeing of the Virgo data is not a major concern).," The median seeing was 1.95 arcseconds for the Virgo data and 1.26 arcseconds for the NGC data (since we will be searching for galaxies with large sizes and low surface-brightnesses, the relatively poor seeing of the Virgo data is not a major concern)."
 For the Virgo data. the median 3m point source limiting magnitude was 2=25.," For the Virgo data, the median $\sigma$ point source limiting magnitude was $B=25$."
 Of the 113 fields we studied. 9 of the fields were also imaged through a Z (Ay=0.90 pim) filter for either. 6005 or 1200 s under photometric conditions.," Of the 113 fields we studied, 9 of the fields were also imaged through a $Z$ $\lambda_0 = 0.90$ $\mu$ m) filter for either 600s or 1200 s under photometric conditions."
 This data was used to compute BoZ colours for a small number of galaxies in the D saniple., This data was used to compute $B-Z$ colours for a small number of galaxies in the $B$ sample.
 Ehe median seeing in the Z dataset was 1.53 arcseconcs ENLM., The median seeing in the $Z$ dataset was 1.53 arcseconds FWHM.
 Initially an image detection algorithim developed by Irwin (1985. 1906) was run on the images.," Initially an image detection algorithim developed by Irwin (1985, 1996) was run on the images."
 A source was cdoelined to be 5 contiguous pixels with a Εαν signal-to-noise ratio of 1.5 above the sky., A source was defined to be 5 contiguous pixels with a flux signal-to-noise ratio of 1.5 above the sky.
 Generated. parameters include. the position. intensity. ancl shape of cach detected object.," Generated parameters include the position, intensity and shape of each detected object."
 We used. only. the position information., We used only the position information.
 Aperture magnitudes were subsequentiy calculated. for all the sources using the IRAPAPPHOT task with a local sky value measured for each object (tvpically the mode within an annulus of inner and outer radii 0.5 and. 1.0 arcminutes centered on the object)., Aperture magnitudes were subsequently calculated for all the sources using the IRAF task with a local sky value measured for each object (typically the mode within an annulus of inner and outer radii 0.5 and 1.0 arcminutes centered on the object).
" In our dataset we had about 10"" galaxies. (about 2500/chip) with 6 magnitudes within a 6 aresecond aperture brighter than (6)=23 (here. as throughout the paper we use the notation D(r) to mean the magnitude within an aperture of radius r areseconds)."," In our dataset we had about $10^6$ galaxies (about 2500/chip) with $B$ magnitudes within a 6 arcsecond aperture brighter than $B(6) = 23$ (here, as throughout the paper we use the notation $B(r)$ to mean the magnitude within an aperture of radius $r$ arcseconds)."
 In this section we describe how we identified which ones are likely to be Virgo Cluster members and assess how much confidence we should have in these identifications., In this section we describe how we identified which ones are likely to be Virgo Cluster members and assess how much confidence we should have in these identifications.
 Alany of the brightest. galaxies (tvpically those with BiG)< 18) had spectroscopic velocity measurements available., Many of the brightest galaxies (typically those with $B(6) < 18$ ) had spectroscopic velocity measurements available.
 For such objects we can identify with 100 confidence which are members and which are not., For such objects we can identify with 100 confidence which are members and which are not.
 We follow Anegeeli et al. (, We follow Binggeli et al. (
1985) in assigning cluster membership to all ealaxies with heliocentric velocities between TOO km s and 2700 kin,1985) in assigning cluster membership to all galaxies with heliocentric velocities between $-$ 700 km $^{-1}$ and 2700 km $^{-1}$.
 For the other galaxies. the situation is more complex.," For the other galaxies, the situation is more complex."
 As outlined in Section 1. we expect dwarf galaxies in the cluster to have lower surface brightnesses (and consequentIys more extended. light. profiles) than background. galaxies. of the same apparent magnitudes.," As outlined in Section 1, we expect dwarf galaxies in the cluster to have lower surface brightnesses (and consequently more extended light profiles) than background galaxies of the same apparent magnitudes."
 Morphologically we expect them. to be either smooth (dwarf elliptical/spheroidal galaxies) or lumpy (chwarf irregulars) butοί to have appreciable spiral structure or a distinctive bulge|disk morphology. both of which are characteristic of luminous giant galaxies (Binney Morrifield 1907) which would here be background galaxies.," Morphologically we expect them to be either smooth (dwarf elliptical/spheroidal galaxies) or lumpy (dwarf irregulars) but to have appreciable spiral structure or a distinctive bulge+disk morphology, both of which are characteristic of luminous giant galaxies (Binney Merrifield 1997) which would here be background galaxies."
 We used. both surface-brightness ancl morphology. to assess menibership as follows., We used both surface-brightness and morphology to assess membership as follows.
 Firstly we compared the light concentrations of galaxies in the Virgo and background samples., Firstly we compared the light concentrations of galaxies in the Virgo and background samples.
 For each galaxy. we computed (following Trentham et al., For each galaxy we computed (following Trentham et al.
 2001a) an inner concentration parameter (ICT) and an outer concentration parameter (OCT). defined as The results are presented in Figure 2.," 2001a) an inner concentration parameter (ICP) and an outer concentration parameter (OCP), defined as The results are presented in Figure 2."
 There is à clear excess of objects with D(4.4)>15 and LCP ον1.8 in the Virgo fields that are not seen in the background. fields., There is a clear excess of objects with $B(4.4) > 18$ and ICP $\sim - 1.3$ in the Virgo fields that are not seen in the background fields.
 Έμπορος are presumably Virgo cluster, These are presumably Virgo cluster
distribution in spirals. but also it questions about its universality.,"distribution in spirals, but also it questions about its universality."
 Llowever. while some of this disagreement may reflect an incllicieney of the URCy to reproduce the RCs. the actual performance of the URC is better than it is claimed.," However, while some of this disagreement may reflect an inefficiency of the $_0$ to reproduce the RCs, the actual performance of the URC is better than it is claimed."
 In fact. spurious cata vs predictions disagreements are created in performing this test ancl precisely when they insert in y the values of Lge and By. allectec by (occasionally large) observational errors.," In fact, spurious data vs predictions disagreements are created in performing this test and precisely when they insert in $_0$ the values of $L_B$ and $R_D$, affected by (occasionally large) observational errors."
 By taking into account this ellect the URC) success rate reaches SOC ane moro., By taking into account this effect the $_0$ success rate reaches $80\%$ and more.
 Willick (1998) found. by studying a Large sample of RCs. a radial variation of the scatter of the inverse RPE defined above and he interpreted it as an evidence against the URC.," Willick (1998) found, by studying a large sample of RCs, a radial variation of the scatter of the inverse RTF defined above and he interpreted it as an evidence against the URC."
 Let us show that this argument is incorrect ancl that the properties of the KLE support the URC paracigm., Let us show that this argument is incorrect and that the properties of the RTF support the URC paradigm.
" The increase/decrease of the scatter found is very small (Wallick. 1998): the scatter ranges [rom 0.065 dex (at 24 ) to 0.080 dex (at 0.5/25 and at 3425) and it implies. if totally intrinsic. à prediction error in fogVGr) of (0.0870.0652)""=0.04dew."," The increase/decrease of the scatter found is very small (Willick, 1998): the scatter ranges from 0.065 dex (at $2 R_D$ ) to 0.080 dex (at $0.5 R_D$ and at $3 R_D$ ) and it implies, if totally intrinsic, a prediction error in $log\ V(x)$ of $(0.08^2-0.065^2)^{0.5} = 0.04\ dex$."
 Moreover. some of the scatter increase/decrease is due to the larger random observational errors present in the outermost measurements: in fact. à refined analysis of the issue (Yegorova et al 2007) finds a smaller. predicting error for three large sample of spirals.," Moreover, some of the scatter increase/decrease is due to the larger random observational errors present in the outermost measurements; in fact, a refined analysis of the issue (Yegorova et al 2007) finds a smaller predicting error for three large sample of spirals."
 Pherefore. from the RTE we have that. in the region considered. the luminosity.sfafisiieally predicts the circular velocity at any radius and in any galaxy within an error of 54-10%. a quantity much smaller than the variations of the latter galaxy and. galaxies.," Therefore, from the RTF we have that, in the region considered, the luminosity predicts the circular velocity at any radius and in any galaxy within an error of -, a quantity much smaller than the variations of the latter galaxy and galaxies."
 We thank L. Danese for helpful discussions. and the referee for useful comments.," We thank L. Danese for helpful discussions, and the referee for useful comments."
 ανν. Lis. et al.," Baldry, I.K., et al."
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rate at which they produce outbursts (Fie. 9)).,"rate at which they produce outbursts (Fig. \ref{fig:numsgrs}) ),"
 aud fiud that for a likely recurrence rate of 0.01 yr1 there would be 13!D in the Galaxy.," and find that for a likely recurrence rate of 0.01 $yr^{-1}$, there would be $13^{+18}_{-7}$ in the Galaxy."
 Towever. these populations may be dwarfed by the nuuber of transicut ANPs similar to NTE Jis10197 (seealsoTbrahimetal.2001).," However, these populations may be dwarfed by the number of transient AXPs similar to XTE J1810–197 \citep[see also][]{ibr04}."
. No new examples were detected in archival aand oobservatious. so up to 510 quiescent AXPs could be prescut in the Galaxy.," No new examples were detected in archival and observations, so up to 540 quiescent AXPs could be present in the Galaxy."
 We also considered the fact that two trausieut ANPs have been discovered in outburst by aandATE., We also considered the fact that two transient AXPs have been discovered in outburst by and.
 The fact thatASCA..RNTE..Chandru.. and ttogether have surveved of the Calaxy each vear for the last 11 vears allows us to constrain the nuuber of transicuts as a function of them recurrence time (Fie. 10)).," The fact that, and together have surveyed of the Galaxy each year for the last 11 years allows us to constrain the number of transients as a function of their recurrence time (Fig. \ref{fig:numtaxps}) )."
 For a recurrence time of once every 10 vears. this calculation suggests there are between 80 aud 580 transicut ΑΝΙΕ with confidence.," For a recurrence time of once every 10 years, this calculation suggests there are between 80 and 580 transient AXPs, with confidence."
 Assuniug a ifetime of LO! vr for persistent maenetars (IKouveliotouotal.1999:Gaensleret 1999).. their birth rate is at cast of that of radio pulsus. and could equal the radio pulsar birth rate (c.g...Lorimerctal.2006:FaucherCieuere&Ikaspi 2006).," Assuming a lifetime of $10^4$ yr for persistent magnetars \citep{kou99,ggv99}, their birth rate is at least of that of radio pulsars, and could equal the radio pulsar birth rate \citep[e.g.,][]{lor06,fgk06}."
. The birth rate of transient uagnetars could exceed those of radio pulsars. unless hei lifetimes are significantly louger than 105 vears.," The birth rate of transient magnetars could exceed those of radio pulsars, unless their lifetimes are significantly longer than $10^4$ years."
 The iain uneertainty in the muuber of Calactic uaenetars is introduced by the unkuown activity level of transient cxamples., The main uncertainty in the number of Galactic magnetars is introduced by the unknown activity level of transient examples.
" Further monitoring of knowu transient ANDPs would coustrain their duty cveles and reduce this uncertaimty, so long as the kuown exaniples can be taken as representative of all maguctars."," Further monitoring of known transient AXPs would constrain their duty cycles and reduce this uncertainty, so long as the known examples can be taken as representative of all magnetars."
 Ultimately. though. a wide-field monitoring campaign is needed to ideutify more transient ANDPs in outburst.," Ultimately, though, a wide-field monitoring campaign is needed to identify more transient AXPs in outburst."
" This can be accomplished in principle iu the N-ray baud. although the sensitivity of current and planned wide-field monitors (2«410HV1, 0.510 is only sufficient to barely detect à 107 nuünagnetar to a distance of 1 kpc from Earth."," This can be accomplished in principle in the X-ray band, although the sensitivity of current and planned wide-field monitors \citep[$\approx$$2\times10^{-11}$, 0.5–10 is only sufficient to barely detect a $10^{35}$ magnetar to a distance of 4 kpc from Earth."
 Ou the other haud. the detectious of NTE JISI0197 etal.2006) aud LE 1517.05108 (C'anüloetal.2007) in the radio with characteristic intensities of 2100 uty kpe? suggests that future radio surveys with the Low-Frequency Array. the Mileura Wide-field Array. or the Square Wilometer Array (e.¢..Cordes.Lazio.&MeLanegllin2001) could provide the best constraiuts ou the numbers of maguetars iu the Galaxy.," On the other hand, the detections of XTE J1810–197 \citep{cam06} and 1E 1547.0–5408 \citep{cam07} in the radio with characteristic intensities of $\ga$ 100 mJy $^{2}$ suggests that future radio surveys with the Low-Frequency Array, the Mileura Wide-field Array, or the Square Kilometer Array \citep[e.g.,][]{clm04} could provide the best constraints on the numbers of magnetars in the Galaxy."
 We thaux E. Cotthelf for sharing the results of his observatious of AN J185:.30128., We thank E. Gotthelf for sharing the results of his observations of AX J1853.3–0128.
" This work made use of data obtained from the High Enerey Astroplivsics Science Archive Research Center (TEASARC). provided by NASA's Goddard Space Flight (οπίου,"," This work made use of data obtained from the High Energy Astrophysics Science Archive Research Center (HEASARC), provided by NASA's Goddard Space Flight Center."
resulting structures. typically of the order of a few arcmin radius.,"resulting structures, typically of the order of a few arcmin radius."
 Earlier observations could reveal rings with simaller radi. but this also requires a good angular resolution. as obtainable c.g. withChandra?. to diseutaugle the cluission of the expanding halo from that of the ceutral afterelow source.," Earlier observations could reveal rings with smaller radii, but this also requires a good angular resolution, as obtainable e.g. with, to disentangle the emission of the expanding halo from that of the central afterglow source."
 It is evident from the two CRB alos discussed above tha the derivation of quantitative information on the uuscattered X-ray cuuission is prescuthy subject to laree systematic uncertainties due to the assunptions required in the halo modeling., It is evident from the two GRB halos discussed above that the derivation of quantitative information on the unscattered X-ray emission is presently subject to large systematic uncertainties due to the assumptions required in the halo modelling.
 Iu this respect. the observation with and/or Swift-NRT.. of dust halos around CRBs for which bright prom Nouv (or carly afterglow) eiiission is detected by Swift would be extremely useful. since it could eive the chance to compare ie Xrav properties inferred from the halo with those actually observed.," In this respect, the observation with and/or Swift-XRT, of dust halos around GRBs for which bright prompt X-ray (or early afterglow) emission is detected by Swift would be extremely useful, since it could give the chance to compare the X–ray properties inferred from the halo with those actually observed."
 The dust halo reported around GRB 050721 (Romanoetal. 20053) has already demonstrated 10 Swift capabilities to detect such features;, The dust halo reported around GRB 050724 \cite{romano}) ) has already demonstrated the Swift capabilities to detect such features.
 A systematic application of the method presented here to all GRB ollow-up observations would possibly cularee the sample wough the detection of fainter halos. not inuuediatelv seen iu the traditional Xτας images.," A systematic application of the method presented here to all GRB follow-up observations would possibly enlarge the sample through the detection of fainter halos, not immediately seen in the traditional X–ray images."
with the absorbing gas covering only a fraction of the polarized source image and requires that the position angle of polarization of the continuum changes across the source image.,with the absorbing gas covering only a fraction of the polarized source image and requires that the position angle of polarization of the continuum changes across the source image.
 The fourth panel shows the percentage circular polarizatioαυ] VY)1., The fourth panel shows the percentage circular polarization $V(v)/I$.
 The absorption line exhibits a dip in fractional Stokes V of 0.03%., The absorption line exhibits a dip in fractional Stokes $V$ of $0.03\%$.
 However. we believe this apparent circular polarization is meaningless.," However, we believe this apparent circular polarization is meaningless."
 The line ts 5% linearly polarized. and some of the total intensity and also the linear polarizatio20 will leak into Stokes V.," The line is $5\%$ linearly polarized, and some of the total intensity and also the linear polarization will leak into Stokes $V$."
 The observed Stokes V intensity of th line is only about 1% of its linearly-polarized intensity., The observed Stokes $V$ intensity of the line is only about $1\%$ of its linearly-polarized intensity.
 ErrorsX in the Stokes parameters at these levels are not unexpected. because of the small day-to-day changes in the Mueller matrix elements.," Errors in the Stokes parameters at these levels are not unexpected, because of the small day-to-day changes in the Mueller matrix elements."
 We now show that the GBT polarization spectra provide new and unique information about the kinematic. and spatial structure of the gas that gives rise to 21 cm absorption toward 3C 286., We now show that the GBT polarization spectra provide new and unique information about the kinematic and spatial structure of the gas that gives rise to 21 cm absorption toward 3C 286.
 In particular we note the following model-independent facts: (1) The difference between the velocity centroids of the Stokes 7 and fractional polarization spectra is inconsistent with a single-cloud origin for the absorption. (, In particular we note the following model-independent facts: (1) The difference between the velocity centroids of the Stokes $I$ and fractional polarization spectra is inconsistent with a single-cloud origin for the absorption. (
2) The Gaussian shape of the absorption profile is plausibly explained through the cental limit theorem by the superposition of many small clouds. (,2) The Gaussian shape of the absorption profile is plausibly explained through the cental limit theorem by the superposition of many small clouds. (
3) The position-angle spectrum reinforces the conclusions drawn from the fractional polarization. but requires detailed modeling as described below.,"3) The position-angle spectrum reinforces the conclusions drawn from the fractional polarization, but requires detailed modeling as described below."
 Besides studying the spatial and kinematic structure of the absorbing gas. we also Investigate its two-phase structure: Le.. whether the gas is a cold neutral medium (CNM) with temperature 7 ~ 100 K. a warm neutral medium (WNM) with T~ 8000 K (Wolfire 1995). a thermally unstable phase with temperature between these two extremes. or some combination of any of these.," Besides studying the spatial and kinematic structure of the absorbing gas, we also investigate its two-phase structure; i.e., whether the gas is a cold neutral medium (CNM) with temperature $T$ $\sim$ 100 K, a warm neutral medium (WNM) with $T~\sim$ 8000 K (Wolfire 1995), a thermally unstable phase with temperature between these two extremes, or some combination of any of these."
 Furthermore. we consider clues to the identity of the galaxy hosting the absorbing gas.," Furthermore, we consider clues to the identity of the galaxy hosting the absorbing gas."
 The 21 em absorption likely forms against a source with the core-jet configuration shown in Fig. (, The 21 cm absorption likely forms against a source with the core-jet configuration shown in Fig. (
Wilkinson 1979).,Wilkinson 1979).
 This VLBI map. obtained at 609 MHz. indicates that 3C 286 Is an asymmetric source consisting of a compact core and an extended jet.," This VLBI map, obtained at 609 MHz, indicates that 3C 286 is an asymmetric source consisting of a compact core and an extended jet."
" The core. which covers a solid angle. Ovo10«5 milli-aresec? (mas?) with major-axis position angle equals 42°. contributes a flux density $,.(609)=3.5 Jy. while in the case of the jet Oi =40\ 15 mas’. PA.=42°. and $,(609)214. Jy."," The core, which covers a solid angle, $\Omega_{\rm core} {\approx}10{\times}5$ $^{2}$ $^{2}$ ) with major-axis position angle equals $^{\circ}$ contributes a flux density $S_{\nu}(609)$ =3.5 Jy, while in the case of the jet $\Omega_{\rm jet}$ $\times$ 15 $^{2}$, $^{\circ}$, and $S_{\nu}(609)$ =14 Jy."
 Since a similar core-jet structure is detected at 329 MHz. 1667 MHz (Simon 1980). 5 GHz (Jiang 1996; Cotton 1997). and 15 GHz (Kellermann 1997. priv.," Since a similar core-jet structure is detected at 329 MHz, 1667 MHz (Simon 1980), 5 GHz (Jiang 1996; Cotton 1997), and 15 GHz (Kellermann 1997, priv."
 comm.).," comm.),"
 it is safe to assume that such a structure Is present at the absorption frequency of 839.4 MHz., it is safe to assume that such a structure is present at the absorption frequency of 839.4 MHz.
 Adopting the spectral indices measured by Simon (1980) Of Ocore= —0.2910.15 and αι-0.55. we find that SY 2:3.2 Jy and S z:11.7 Jy at i zz 840 MHz.," Adopting the spectral indices measured by Simon (1980) of ${\alpha_{\rm core}}=-0.29{\pm}$ 0.15 and $\alpha_{\rm jet}=-0.55$, we find that $S_{\nu}^{\rm core}$${\approx}$ 3.2 Jy and $S_{\nu}^{\rm jet}$ 11.7 Jy at $\nu$ $\approx$ 840 MHz."
 While the polarization. structure of 3C 286 hàs not been detected at low frequencies. VLBI measurements. of all four Stokes parameters have been obtained at 5 GHz (Hang 1996; Cotton 1997; Cotton 2010 [priv.," While the polarization structure of 3C 286 has not been detected at low frequencies, VLBI measurements of all four Stokes parameters have been obtained at 5 GHz (Jiang 1996; Cotton 1997; Cotton 2010 [priv."
 comm.]|., comm.]).
 These data show that unlike most quasar VLBI sources. the fractional polarizations. te.. fraction of total surface brightness in V/U7-Q. where UÜ and Q are the Stokes parameters. of the core and jet are comparable.," These data show that unlike most quasar VLBI sources, the fractional polarizations, i.e., fraction of total surface brightness in $\sqrt{U^{2}+Q^{2}}$, where $U$ and $Q$ are the Stokes parameters, of the core and jet are comparable."
 More specifically. the ratio of jet to core polarized surface brightnesses is about 0.4.," More specifically, the ratio of jet to core polarized surface brightnesses is about 0.4."
 Nevertheless the ratio of polarized flux densities PA/pro zz 5 owing to the much larger solid angle subtended by the jet. where the polarized flux density P=4Q2.," Nevertheless the ratio of polarized flux densities $P^{jet}/P^{core}$ $\approx$ 5 owing to the much larger solid angle subtended by the jet, where the polarized flux density $P=\sqrt{U^{2}+Q^{2}}$."
 Because of the steeper spectral index of the jet. /U2it 1s reasonable to assume that this ratio increases with decreasing frequency.," Because of the steeper spectral index of the jet, it is reasonable to assume that this ratio increases with decreasing frequency."
 As a result. we shall ignore the core source when computing absorption in Stokes U and Q parameters. or equivalently in the spectrum of polarized flux density P (note. P=pl) and polarization position angle \=0.5« arctan(U/Q).," As a result, we shall ignore the core source when computing absorption in Stokes $U$ and $Q$ parameters, or equivalently in the spectrum of polarized flux density $P$ (note, $P=pI$ ) and polarization position angle $\chi=0.5{\times}$ $U/Q$ )."
 We check the self-consistency of this assumption in 5 7., We check the self-consistency of this assumption in $\S$ 7.
 We now describe the procedures used to model the Stokes parameter spectra., We now describe the procedures used to model the Stokes parameter spectra.
 To aecount for the absorption spectrum we first. adopt a simplified model. illustrated in. Fig.3," To account for the absorption spectrum we first adopt a simplified model, illustrated in Fig.,"
. comprised of one uniform H I ‘cloud’ toward the jet (cloud 1) and another uniform H I ‘cloud’ toward the core (cloud 2)., comprised of one uniform H I `cloud' toward the jet (cloud 1) and another uniform H I `cloud' toward the core (cloud 2).
 A single cloud model is ruled out by past VLBI observations of the 839.4 MHz absorption feature that indicated an inhomogeneous structure of the absorbing gas., A single cloud model is ruled out by past VLBI observations of the 839.4 MHz absorption feature that indicated an inhomogeneous structure of the absorbing gas.
 The velocity difference between the maximum depths of the phase-shift and fringe amplitude spectra detected by a two-element VLBI experiment (Wolfe 1976) indicates a velocity difference between the H I towards the core and Jet sources., The velocity difference between the maximum depths of the phase-shift and fringe amplitude spectra detected by a two-element VLBI experiment (Wolfe 1976) indicates a velocity difference between the H I towards the core and jet sources.
" As a result the relative change in Stokes / parameter at velocity vis given by where the number of clouds 5,22. C, and C» are the area covering factors of the gas toward the jet and core. fj and f> are the fractions of Stokes 7 flux density in the jet and the core sources.and τιν) and 720"") are the 21 em optical depthsaveraged across the areas subtended by jet and core at the absorber."," As a result the relative change in Stokes $I$ parameter at velocity $v$ is given by where the number of clouds $n_{c}$ =2, $C_{1}$ and $C_{2}$ are the area covering factors of the gas toward the jet and core, $f_{1}$ and $f_{2}$ are the fractions of Stokes $I$ flux density in the jet and the core sources,and ${\tau_{1}(v)}$ and ${\tau_{2}(v)}$ are the 21 cm optical depthsaveraged across the areas subtended by jet and core at the absorber."
 Note. / in the denominator of eq. (," Note, $I$ in the denominator of eq. ("
4) is the,4) is the
"cold matter. The discovery of broadening due to gravitational redshift and Doppler shifts of the iron K,, fluorescence lines in Seyfert X-ray spectra suggests that at least part of the reprocessing takes place in the inner parts of the accretion disk € Tanakaetal. 1995.. Nandraetal. 1997.. Reynolds&Begelman 19971).","cold matter, The discovery of broadening due to gravitational redshift and Doppler shifts of the iron $_{\alpha}$ fluorescence lines in Seyfert X-ray spectra suggests that at least part of the reprocessing takes place in the inner parts of the accretion disk ( \ncite{Tanaka}, \ncite{Nandra97}, \ncite{Reynolds}) )."
" However. the detection of narrow components to the Fe K,, line with (eg Weaveretal. 19973) and recently (Yaqoobetal.2001... Kaspietal. 20023) and (Reevesetal.2001) suggests reprocessing of X-rays in matter ata arger distance from the black 10le. eg in the BLR or in the molecuar torus."," However, the detection of narrow components to the Fe $_\alpha$ line with (eg \ncite{Weaver97}) ) and recently \ncite{Yaqoob}, \ncite{Kaspi}) ) and \cite{Reeves}
 suggests reprocessing of X-rays in matter at a larger distance from the black hole, eg in the BLR or in the molecular torus."
 The analysis of X-ray variabiliy is a powerful tool for the investigation of the inner regions of 1ο AGN central engine. since he time-scales of the variability ¢an give an indication of the geometrical sizes of the regions inv'olved.," The analysis of X-ray variability is a powerful tool for the investigation of the inner regions of the AGN central engine, since the time-scales of the variability can give an indication of the geometrical sizes of the regions involved."
 Yaqoobetal.(1996) reported a rapid (<3-107 s) res;xonse of the Fe fluorescence ine flux to the continuum flux in NGC 7314., \scite{Yaqoob96} reported a rapid $<3\cdot 10^4$ s) response of the Fe fluorescence line flux to the continuum flux in NGC 7314.
 From the resulting maximum distance of the reflecting material from the primary ray source they derive an upper limit on the mass of the black 10le., From the resulting maximum distance of the reflecting material from the primary X-ray source they derive an upper limit on the mass of the black hole.
 However. tdhe analysis of fluorescence line variability in a number of other Seyfert galaxies have led to confusing and partly contradictory results.," However, tdhe analysis of fluorescence line variability in a number of other Seyfert galaxies have led to confusing and partly contradictory results."
 While in some sources a positive correlation of line and continuum flux has been observed. averaged over long timescales. (eg NGC 5506. Lamer.Uttley&M'Hardy 2000). other studies found no evidence for a close relationship between continuum and fluorescence line (Weaver.Gelbord&|Yagoob 2001.. Vaughan&Edelson 20019.," While in some sources a positive correlation of line and continuum flux has been observed, averaged over long timescales, (eg NGC 5506, \ncite{Lamer2000}) ), other studies found no evidence for a close relationship between continuum and fluorescence line \ncite{Weaver2001}, \ncite{Vaughan}) )."
 Flux-dependent variability of the continuum spectral shape itself can be used to constrain models of the Comptonising accretion disk corona teg Haardt.Maraschi&Ghisellini 1997)})., Flux-dependent variability of the continuum spectral shape itself can be used to constrain models of the Comptonising accretion disk corona (eg \ncite{Haardt}) ).
 In general. the continuum slope appears to be correlated with flux teg Leighlyetal.1996:: M'Hardy.Papadakis&Uttley1998 Lamer.Uttley&M'Hardy 20000). as predicted by most simple Comptonisation models.," In general, the continuum slope appears to be correlated with flux (eg \ncite{Leighly}; \ncite{mch98} ; \ncite{Lamer2000}) ), as predicted by most simple Comptonisation models."
 However. it is important to determine whether the degree of slope variability. and the form of the correlation with flux (eg does the continuum slope saturate at some maximum value?)," However, it is important to determine whether the degree of slope variability, and the form of the correlation with flux (eg does the continuum slope saturate at some maximum value?)"
 are in agreement with existing models., are in agreement with existing models.
 NGC 4051 is a nearby (720.0023) low luminosity Seyfert | galaxy. whieh is among the most variable AGN in the X-ray band (Green.MHardy&Lehto1993).," NGC 4051 is a nearby (z=0.0023) low luminosity Seyfert 1 galaxy, which is among the most variable AGN in the X-ray band \cite{Green}."
. We have been monitoring NGC 4051 with RXTE since 1996., We have been monitoring NGC 4051 with RXTE since 1996.
 Since our monitoring commenced. the 2-10 keV flux from the object has varied by a factor of ~ 100.," Since our monitoring commenced, the 2-10 keV flux from the object has varied by a factor of $\sim$ 100."
 In 1998 NGC 4051 entered a state of extremely low X-ray flux that lasted for ~150 days (Guainazzietal.1998.. Uttleyetal. 1999)).," In 1998 NGC 4051 entered a state of extremely low X-ray flux that lasted for $\sim 150$ days \ncite{Guainazzi}, \ncite{Uttley99}) )."
 Here we report on the spectral changes that accompanied. the dramatic flux variability., Here we report on the spectral changes that accompanied the dramatic flux variability.
 In section 2. we describe the observations and the reduction of the ARXTE data., In section \ref{obs} we describe the observations and the reduction of the data.
 Results from observations in May 1998. during the extreme low state. are re-analysed in a consistent manner with the other observations and are discussed in section 3..," Results from observations in May 1998, during the extreme low state, are re-analysed in a consistent manner with the other observations and are discussed in section \ref{lowstate}."
 In section + we present the analysis of the continuum and fluorescence line variations observed during the long term monitoring and during anAXTE long look in December 1996., In section \ref{specvar} we present the analysis of the continuum and fluorescence line variations observed during the long term monitoring and during an long look in December 1996.
 We have been monitoring NGC 4051 with RATE since May 1996., We have been monitoring NGC 4051 with RXTE since May 1996.
 The observations were separated by a range of time intervals., The observations were separated by a range of time intervals.
 For the first 6 months the observations took place weekly with periods of twice-daily observations in May 1996 and daily observations in October 1996., For the first 6 months the observations took place weekly with periods of twice-daily observations in May 1996 and daily observations in October 1996.
 Since November 1996 the source has been observed at fortnightly intervals., Since November 1996 the source has been observed at fortnightly intervals.
 observed NGC 4051 with the Proportional Counter Array (PCA) and the High Energy ray Timing Experiment (HEXTE)., observed NGC 4051 with the Proportional Counter Array (PCA) and the High Energy X-ray Timing Experiment (HEXTE).
 The PCA (Zhangetal.1993) consists of 5 Xenon-filled Proportional Counter Units (PCUs). sensitive to X-ray energies from 2-60 keV. The maximum effective area of the PCA is 6500 eni.," The PCA \cite{Zhang} consists of 5 Xenon-filled Proportional Counter Units (PCUs), sensitive to X-ray energies from 2-60 keV. The maximum effective area of the PCA is 6500 $^2$."
 As he signal to noise ratios of he higher energy HEXTE spectra on NGC 4051 are low. we only present the results from the PCA.," As the signal to noise ratios of the higher energy HEXTE spectra on NGC 4051 are low, we only present the results from the PCA."
 For technical reasons the igh voltage settings of the PCA uniς have been adjusted severa imes during their lifetime., For technical reasons the high voltage settings of the PCA units have been adjusted several times during their lifetime.
 In this paper we only present data aken during the PCA gain epoch 3. which lasted from April 1996 o March 1999.," In this paper we only present data taken during the PCA gain epoch 3, which lasted from April 1996 to March 1999."
 The total exposure time of the 130 individua monitoring observations during gain epoch 3 is 85 Ksec., The total exposure time of the 130 individual monitoring observations during gain epoch 3 is 85 ksec.
 Apar rom the monitoring observations performed two long-look observations of NGC 4051 in December 1996 with a total source time of 36.4 ksec within 3 days and in May 1998 for 84.4 Ksec ton-source) within2 days., Apart from the monitoring observations performed two long-look observations of NGC 4051 in December 1996 with a total on-source time of 36.4 ksec within 3 days and in May 1998 for 84.4 ksec (on-source) within2 days.
 We have used for the reduction of the PCA and HEXTE data., We have used for the reduction of the PCA and HEXTE data.
" PCA “good times"" have been selected from he Standard 2. mode data sets using the following criteria: target elevation =—10. pointing offset «0.01%. time since SAA yassage >30 min. standard threshold for electron contamination."," PCA “good times” have been selected from the Standard 2 mode data sets using the following criteria: target elevation $> 10^{\circ}$, pointing offset $<0.01^{\circ}$, time since SAA passage $> 30$ min, standard threshold for electron contamination."
 Occasionally one or more of the 5 Proportional Counter Units (PCUs) of the PCA were switched off during the observations., Occasionally one or more of the 5 Proportional Counter Units (PCUs) of the PCA were switched off during the observations.
 In order to maximise the signal to noise ratio of the spectra. we used the data from all PCUs that were switehed on during each individual pointing and used only PCU layer |.," In order to maximise the signal to noise ratio of the spectra, we used the data from all PCUs that were switched on during each individual pointing and used only PCU layer 1."
 We calculated the background in the PCA with the tool using he L7 model for faint sources. which is suitable for determining he PCA background for energies <24 keV. From a series of 100 blank field observations from the public archive (proposal ID P30801) we determined the uncertainties in the estimated background rates as a function of energy.," We calculated the background in the PCA with the tool using the L7 model for faint sources, which is suitable for determining the PCA background for energies $\leq 24$ keV. From a series of $\sim 100$ blank field observations from the public archive (proposal ID P30801) we determined the uncertainties in the estimated background rates as a function of energy."
 At energies below 24 keV the lo uncertainty of the background is of the order LOletsstheV+ (see Lamer.Uttlley&M°Hardy2000— for details)., At energies below 24 keV the $\sigma$ uncertainty of the background is of the order $0.01 {\rm cts\; s^{-1} keV^{-1}}$ (see \ncite{Lamer2000} for details).
 The PCA response matrices were calculated invididually or each observation usingV2.37. taking into account emporal variation of the detector gain and the changing numbers of detectors used.," The PCA response matrices were calculated invididually for each observation using, taking into account temporal variation of the detector gain and the changing numbers of detectors used."
 Our monitoring with the RATE PCA revealed that during a 150 day period from January 1998 to May1998 NGC 4051 was in a state of unusually low X-ray flux (Uttleyetal.1999. see Fig. 19).," Our monitoring with the RXTE PCA revealed that during a 150 day period from January 1998 to May1998 NGC 4051 was in a state of unusually low X-ray flux \ncite{Uttley99}, see Fig. \ref{lightcurve}) )."
 Near the of endthis period NGC 4051 was observed simultaneously byRXTE (Uttleyetal.1999). and Guainazzietal., Near the end of this period NGC 4051 was observed simultaneously by \cite{Uttley99} and \scite{Guainazzi}.
 (1998).. We derived a 2-10 keV flux of 2.0-10erestem7 from the RXTE PCAspectrum and 1.4-LO12eres1con7 2ofrom the MECS data.," We derived a 2-10 keV flux of $2.0 \cdot 10^{-12} \rm{erg s^{-1} cm^{-2} }$ from the PCAspectrum and $1.4 \cdot 10^{-12} \rm{erg s^{-1}
cm^{-2} }$ from the MECS data."
 Bothobservations showed a very hard X-rayspectrum and unresolved 6.4 keV iron line emission with a very high equivalent width (~1000 eV)., Bothobservations showed a very hard X-rayspectrum and unresolved 6.4 keV iron line emission with a very high equivalent width $\sim 1000$ eV).
 This spectrum was interpreted as a pure reflection component with the central source, This spectrum was interpreted as a pure reflection component with the central source
calculate the deusity aud velocity configurations alone the line of sight exhibiting minim dissipation.,calculate the density and velocity configurations along the line of sight exhibiting minimum dissipation.
" Wigh-quality data of the Thibble Deep Field South (IIDF-S) quasar J2233606 (24,=2.238. Box 17.5) Were obtained durius the conunissioning of the UVES on the VLT 8.211 Kueven telescope at Paranal (Chile) in October 1999."," High-quality data of the Hubble Deep Field South (HDF-S) quasar J2233–606 $z_{\rm em} = 2.238$, B $\simeq 17.5$ ) were obtained during the commissioning of the UVES on the VLT 8.2m Kueyen telescope at Paranal (Chile) in October 1999."
 The resolving powers at which the spectra Were recorded im the spectral ranges AA=305010000 ((J2233606) and AA=3700)5000.6670.—9100 ((Q0000-2620) were Rox15000 and R~ 19000. which correspond o velocity resolutious of FWIIMN ~6.7 kins ! and 26.1 kins . respectively.," The resolving powers at which the spectra were recorded in the spectral ranges $\lambda\lambda = 3050-10000$ (J2233–606) and $\lambda\lambda = 3700-5000, 6670-9100$ (Q0000-2620) were $R \simeq 45000$ and $R \simeq 49000$ , which correspond to velocity resolutions of FWHM $\simeq 6.7$ km $^{-1}$ and $\simeq 6.1$ km $^{-1}$, respectively."
 The data reduction and the xlenti&cation of metal absorptiou-line systems in the J2233606 spectrum are reported in Cristiani D'Odorico (2000)., The data reduction and the identification of metal absorption-line systems in the J2233–606 spectrum are reported in Cristiani D'Odorico (2000).
 Tn our study we also used the J2233606 echelle spectrum CRzz30000. ÀA22753118 ÀJ) obtained with the IIST/STIS (Savaglio 1998).," In our study we also used the J2233–606 echelle spectrum $R \simeq 30000$, $\lambda\lambda~2275-3118$ ) obtained with the HST/STIS (Savaglio 1998)."
 The complete description of our computational procedure is given in Paper E. Here we stummarize briefly basic model asstuuptious and emphasize new details receuthy inchided in the MCT., The complete description of our computational procedure is given in Paper I. Here we summarize briefly basic model assumptions and emphasize new details recently included in the MCI.
 We asstune that all lines observedin a metal system. arise in a continuous absorbing gas slab of a thickuess L (prestunably the outer region of a foreground distaut ealaxv)., We assume that all lines observedin a metal system arise in a continuous absorbing gas slab of a thickness $L$ (presumably the outer region of a foreground distant galaxy).
 The absorber exhibits a fluctuating gas deusity and a iuixture of bulk motions such as iufall aud outflows. tidal Hows ete..," The absorber exhibits a fluctuating gas density and a mixture of bulk motions such as infall and outflows, tidal flows etc.,"
 resulting m a stochastic velocity field., resulting in a stochastic velocity field.
 Aletal abuudauces are assumed to be coustaut within the absorber and gas is supposed to be optically thin for the ionizing UV radiation., Metal abundances are assumed to be constant within the absorber and gas is supposed to be optically thin for the ionizing UV radiation.
 Within the absorbing region the radial velocity ¢(s) id total hydrogen deusitv vy(s} distributions along the —-—iue of sight are the same for all ious., Within the absorbing region the radial velocity $v(s)$ and total hydrogen density $n_{\rm H}(s)$ distributions along the line of sight are the same for all ions.
 Iu the computational procedure these two random fields are represented bv their saumpled values at equally spaced intervals Ac (Ge is dimensionless radial coordinate s/L). ie. bv the vectors peeep} and [m....ny} with & large enough to describe the narrowest components of complex spectral lines.," In the computational procedure these two random fields are represented by their sampled values at equally spaced intervals $\Delta x$ $x$ is dimensionless radial coordinate $s/L$ ), i.e. by the vectors $\{ v_1, \ldots, v_k \}$ and $\{ n_1, \ldots, n_k \}$ with $k$ large enough to describe the narrowest components of complex spectral lines."
 Further we suppose the thermodvuaimic aud ionization equilibiun at cach computational poiut along the sielthue which means that fractional jonizations of different lous are determined exclusively x the gas density and vary from poiut to poiut., Further we suppose the thermodynamic and ionization equilibrium at each computational point along the sightline which means that fractional ionizations of different ions are determined exclusively by the gas density and vary from point to point.
 These fractional ionization variations are just the cause of the observed diversity of profile shapes., These fractional ionization variations are just the cause of the observed diversity of profile shapes.
 To calculate the Kinetic teniperature and fractional ionization of ious the photoionization code CLOUDY (Ferland 1997) was usec., To calculate the kinetic temperature and fractional ionization of ions the photoionization code CLOUDY (Ferland 1997) was used.
 The inputs to CLOUDY are the cimeusiouless ionization parameter C=gyn/ip (yn the nuuber density of photons with energies above 1 Rv). metallicity and the backeround ionizius spectrum for which the Taardt-\ladau spectimm (IEM) was adopted (Ulaardt Madan 1996).," The inputs to CLOUDY are the dimensionless ionization parameter $U = n_{\rm ph}/n_{\rm H}$ $n_{\rm ph}$ – the number density of photons with energies above 1 Ry), metallicity and the background ionizing spectrum for which the Haardt-Madau spectrum (HM) was adopted (Haardt Madau 1996)."
" The number density of the ionizing photons for this spectra equals at 2= to 2.26ςla7 cm? [here e. fi. n and J, are the speed of light. the Planck constant. the frequency of the hydrogen Lyinan οσο, aud the specific iuteusitv (Gu ergs 7 | | +) at frequency p. respectively]."," The number density of the ionizing photons for this spectrum equals at $z = 2$ to $2.26\times10^{-5}$ $^{-3}$ [here $c$, $h$, $\nu_{\rm c}$ and $J_\nu$ are the speed of light, the Planck constant, the frequency of the hydrogen Lyman edge, and the specific intensity (in ergs $^{-2}$ $^{-1}$ $^{-1}$ $^{-1}$ ) at frequency $\nu$, respectively]."
 Fractional ionizatiou curves Y(U) were computed with CLOUDYfor solar abunudauce pattern aud differeut metallicities. and τοι icluded: iu the MCI code. to calculate the optical depths for ions involved in the fitting., Fractional ionization curves $\Upsilon (U)$ were computed with CLOUDYfor solar abundance pattern and different metallicities and then included in the MCI code to calculate the optical depths for ions involved in the fitting.
 If the obtained solution revealed the abundance pattern different. from the solar one. Y(U) were recalculated for this new pattern aud all computations repeated.," If the obtained solution revealed the abundance pattern different from the solar one, $\Upsilon (U)$ were recalculated for this new pattern and all computations repeated."
 It should be noted. however. that differeuces of 0.2:0.3 dex from solar values influence the fractional jiouizations ouly weakly.," It should be noted, however, that differences of $0.2\div0.3$ dex from solar values influence the fractional ionizations only weakly."
 The values of velocity and density at stbsequent coluputational points are considered to be correlated and are described by means of Markovian processes., The values of velocity and density at subsequent computational points are considered to be correlated and are described by means of Markovian processes.
" Iu particular. the velocity is computed as follows: where f.=ΠΔ)02. Ry beg the correlation between the velocity values at poiuts separated bv a distance Aw. ie. Re=(0G|A)ο), ov the velocity dispersion of bulk motions. aud € a random normal variable with zero mean and dispersion Deusity is supposed to have a loe-norimal distribution (characterized by the mean 29 and the second ceutral moment cy of dimensionless variable g — njgpny ) aud is described with help of an auxiliary Markovian process wed: where fy.=Πέ.)07. R, being the correlation between the values of rr at pointsseparated by a distance Aw. 0, the logavitlinic density dispersion. 9,= and € a random normal variable with zero mean and dispersion 9,,—σι,Vlfe. ILwius "," In particular, the velocity is computed as follows: where $f_{\rm v} = R_{\rm v}(\Delta x)/\sigma^2_{\rm v}$, $R_{\rm v}$ being the correlation between the velocity values at points separated by a distance $\Delta x$, i.e. $R_{\rm v} = \langle v(x + \Delta x)\,v(x) \rangle$, $\sigma_{\rm v}$ the velocity dispersion of bulk motions, and $\epsilon$ a random normal variable with zero mean and dispersion Density is supposed to have a log-normal distribution (characterized by the mean $n_0$ and the second central moment $\sigma_{\rm y}$ of dimensionless variable $y$ = $n_{\rm H}$ $n_0$ ) and is described with help of an auxiliary Markovian process $\nu (x)$ : where $f_{\nu} = R_\nu(\Delta x)/\sigma^2_\nu$, $R_\nu$ being the correlation between the values of $\nu$ at pointsseparated by a distance $\Delta x$ , $\sigma_\nu$ the logarithmic density dispersion, $\sigma_\nu = \sqrt{\ln(1 + \sigma^2_{\rm y})}$ and $\epsilon$ a random normal variable with zero mean and dispersion $\sigma_{\epsilon,\nu} = \sigma_\nu\sqrt{1 - f^2_{\nu}}$."
"defined vir). the total hydrogen density can be obtained as The optical depth 7z,,;(ÀA) at waveleneth A for clement ""uw in /th ionizing stage iscalculated from the equation:"," Having defined $\nu(x)$ , the total hydrogen density can be obtained as The optical depth $\tau_{{\rm a},i}(\lambda)$ at wavelength $\lambda$ for element `a' in $i$ th ionizing stage iscalculated from the equation:"
of the merger rate.,of the merger rate.
 Indeed. Mateus(2008) rest match. with the observations is achieved for à mass ratio of 1:10.," Indeed, \citet{mateus2008} best match with the observations is achieved for a mass ratio of 1:10."
 lt is also not entirely clear what is he CAS major mereer ratio sensitivity., It is also not entirely clear what is the CAS major merger ratio sensitivity.
 This is generally assumed. to. be 0.25. but it could. equally. be 0.3. or lower.," This is generally assumed to be 0.25, but it could equally be 0.3, or lower."
 We know this is unlikely. based on comparing CAS and ors for nearby ealaxymergers (LHernandez-Toledoetal.2005: DeProprisςal. 2007)).," We know this is unlikely, based on comparing CAS and pairs for nearby galaxymergers \citealt{hernandez2005}; \citealt{depropris2007}) )."
 For a basic agreement at high redshift or at low stellar masses. the CAS method would have o be sensitive to mergers with mass ratios do wnto 1:10.," For a basic agreement at high redshift or at low stellar masses, the CAS method would have to be sensitive to mergers with mass ratios down to 1:10."
 This is again unlikely to be the major reason for tie difference between mocels and observations. because the variations involvec are too small in comparison with the discrepancies in the results.," This is again unlikely to be the major reason for the difference between models and observations, because the variations involved are too small in comparison with the discrepancies in the results."
 However. it is a further indication tiu the mass raio threshold. to identify major mergers is indeed a rather volatile factor both in the observations and in he simulations.," However, it is a further indication that the mass ratio threshold to identify major mergers is indeed a rather volatile factor both in the observations and in the simulations."
 Finally. the mismatch between the οoervations and simulation results might » due to the inability. of the Millennium. simulation to ooduce enough Major Merecrs.," Finally, the mismatch between the observations and simulation results might be due to the inability of the Millennium simulation to produce enough major mergers."
 This can happen. for examxe. if there are oo [ew galaxies in haloes. or if the predicted merging time-scale is too long.," This can happen, for example, if there are too few galaxies in haloes, or if the predicted merging time-scale is too long."
 The lack of galaxies in haloes can be tested by measuring the bias of massive ὃνgalaxies in the Millennium simulation. which have to be more clustered. than the dark matter.," The lack of galaxies in haloes can be tested by measuring the bias of massive galaxies in the Millennium simulation, which have to be more clustered than the dark matter."
 Zehavietal.(2005). found that the predictions of semi-analytical models agree fairly well with the clustering properties. of nearby galaxies., \citet{zehavi2005} found that the predictions of semi-analytical models agree fairly well with the clustering properties of nearby galaxies.
 However. Coiletal.(2008) found that at high redshift simulated blue galaxies are much less clustered than real blue galaxies. especially at small scales.," However, \citet{coil2008} found that at high redshift simulated blue galaxies are much less clustered than real blue galaxies, especially at small scales."
 As the vast majority of merging galaxies at z21 are blue (Conseliceetal. 2003b).. this is a potential problem of the mocdels. which translates in to an underestimate of the merger history.," As the vast majority of merging galaxies at $z > 1$ are blue \citep{cons2003b}, this is a potential problem of the models, which translates in to an underestimate of the merger history."
 In fact. a lack of a strong bias. or equivalently a lower halo occupation number. implies that massive svstems clo not experience as many merger events as they would if they were more clustered.," In fact, a lack of a strong bias, or equivalently a lower halo occupation number, implies that massive systems do not experience as many merger events as they would if they were more clustered."
 In this Subsection we bricks compare our findings with results from. previous works., In this Subsection we briefly compare our findings with results from previous works.
 Guo&White(2008) examine how galaxies grow within the Millennium model., \citet{guo2008} examine how galaxies grow within the Millennium model.
 They find that the mos massive ealaxies have a formation history dominated. by major mergers.while for less massive galaxies the same is true only at z>IL.," They find that the most massive galaxies have a formation history dominated by major mergers,while for less massive galaxies the same is true only at $z>1$."
 Witzbichler&White(2008) use the mocel of DeLucia&Dlaizot.(2007) to calibrate the relationship between the abundance of close galaxy pairs ancl the rate of ealaxv mergers at high. redshifts., \citet{manfred2008} use the model of \citet{delucia2007} to calibrate the relationship between the abundance of close galaxy pairs and the rate of galaxy mergers at high redshifts.
 They. find that close galaxy pairs merge within a few Car. and that this merging time-scale only weakly increases with decreasing redshift.," They find that close galaxy pairs merge within a few Gyr, and that this merging time-scale only weakly increases with decreasing redshift."
 lxitzbichler&White(2008) also argue that while the use of close pairs for estimating the merger history of galaxies is à reliable tool. the merging times used to calculate the merger rates in observations are at least a factor of two too short with respect to their findings. which translates in to an overestimate of the merger rates by a similar factor.," \citet{manfred2008} also argue that while the use of close pairs for estimating the merger history of galaxies is a reliable tool, the merging times used to calculate the merger rates in observations are at least a factor of two too short with respect to their findings, which translates in to an overestimate of the merger rates by a similar factor."
 Patton&Atlield.(2008) show that this is not the case when an appropriate correction factor is included. in the equations., \citet{patton2008} show that this is not the case when an appropriate correction factor is included in the equations.
 Similar conclusions are reached by Wetzel.etal(2008).. who find that only pairs at very small separations can be considered a reliable proxy for the global merger population.," Similar conclusions are reached by \citet{wetzel2008b}, who find that only pairs at very small separations can be considered a reliable proxy for the global merger population."
 Using galaxy pair fractions. Mateus(2008). finds a similar qualitative behaviour with redshift for the merger rate as we do.," Using galaxy pair fractions, \citet{mateus2008} finds a similar qualitative behaviour with redshift for the merger rate as we do."
 Alateus(2008) also compares the results from the model of DeLucia&Blaizot(2OT) with those from the model of Boweretal.(2006) and finds that the 3oweretal.(2006) model better reproduces the qualitative behaviour of the merger rate with redshift tiin the DeLucia&Blaizot(2007) model. although the predicted: values are still lower than observed.," \citet{mateus2008} also compares the results from the model of \citet{delucia2007} with those from the model of \citet{bower2006} and finds that the \citet{bower2006} model better reproduces the qualitative behaviour of the merger rate with redshift than the \citet{delucia2007} model, although the predicted values are still lower than observed."
 This is interesting. because etal.(2006). use different merger trees than DeLucia&Dlaizot(2007). and Jertoneοἱal.(2007) and have cüfferent prescriptions for merging that might better reproduce the evolution of the observec galaxy population.," This is interesting, because \citet{bower2006} use different merger trees than \citet{delucia2007} and \citet{bertone2007} and have different prescriptions for merging that might better reproduce the evolution of the observed galaxy population."
 Other studies have compared. the observed. galaxy merger properties with the Millennium model (Geneletal.2008: Patton&Atlield2008))., Other studies have compared the observed galaxy merger properties with the Millennium model \citealt{genel2008}; ; \citealt{patton2008}) ).
 Patton&Atheld(2008) compare the number of close galaxy. pairs (with separations smaller than 20 kkpe) in the simulation and in a sample of SDSS galaxies at low redshift; and find goodagreement., \citet{patton2008} compare the number of close galaxy pairs (with separations smaller than 20 kpc) in the simulation and in a sample of SDSS galaxies at low redshift and find goodagreement.
" Patton&At-Ποια(2008) estimate that at least 90. per cent. of major mergers occur between galaxies fainter than L,.", \citet{patton2008} estimate that at least 90 per cent of major mergers occur between galaxies fainter than $L_{\star}$.
 Geneletal.(2008). find that there are not enough major mergers in the Millennium. simulation to transform the population of >=2 galaxies in to the elliptical galaxies observed at z=0. which implies that in the model internal evolution must play a more important role to shape their properties than in the real universe.," \citet{genel2008} find that there are not enough major mergers in the Millennium simulation to transform the population of $z=2$ galaxies in to the elliptical galaxies observed at $z=0$, which implies that in the model internal evolution must play a more important role to shape their properties than in the real universe."
 We have carried out a comparison between galaxy merger history predictions based on the Millennium simulation aud a collection of observed. galaxy merger fractions ancl rates taken from Conseliceetal. (2003b).. Conseliceetal. (2008).. Dlucketal.(2000) ancl Conselicectal. (2009)...," We have carried out a comparison between galaxy merger history predictions based on the Millennium simulation and a collection of observed galaxy merger fractions and rates taken from \citet{cons2003b}, \citet{cons2008}, , \citet{bluck2009} and \citet{cons2009}. ."
 The observed merger rates and fractions are measured using the CAS system and by measuring the number of galaxy pairs ab ll 2., The observed merger rates and fractions are measured using the CAS system and by measuring the number of galaxy pairs at $z>2$ .
 Our main results for the Millennium galaxy formation model of Bertoneetal.(2007) can besumimarised as follows:, Our main results for the Millennium galaxy formation model of \citet{bertone2007} can besummarised as follows:
From the above discussion it is clear why the amplitude of the maxima following the pulse is wach sinaller in the horizontal than iu the vertical direction.,From the above discussion it is clear why the amplitude of the maxima following the pulse is much smaller in the horizontal than in the vertical direction.
 The behavior of the eroup velocity and the distribution of 4 values iu space (sec Figs., The behavior of the group velocity and the distribution of $k-$ values in space (see Figs.
 1 and 2)) show that in the horizouta direction a narrow range of & values near &= Lis disperse over a large portion of plivsical space. while the dispersion is ΙΟ lower in the vertical direction: energy is therefor: much more diluted iu the horizontal than in the vertica direction aud. as a consequence. the maxima have 1c[um lower amplitude.," \ref{fig:vgvsk} and \ref{fig:vg}) ) show that in the horizontal direction a narrow range of $k$ values near $k=1$ is dispersed over a large portion of physical space, while the dispersion is much lower in the vertical direction; energy is therefore much more diluted in the horizontal than in the vertical direction and, as a consequence, the maxima have much lower amplitude."
 This behavior becomes even more evident from the fractional energv contained in a cone of infinitesimal opening augle when its svuuuctry axis points in differen irectious., This behavior becomes even more evident from the fractional energy contained in a cone of infinitesimal opening angle when its symmetry axis points in different directions.
 Fig., Fig.
 1l shows this fractional energy at the time when the front has reached the height of 1577. normalized o the total energv. contained m a vertical cone. agains istance nieasured from the frout of a wave.," \ref{fig:efrac} shows this fractional energy at the time when the front has reached the height of $15 H$, normalized to the total energy contained in a vertical cone, against distance measured from the front of a wave."
 Note tha in the horizontal directiou the energw rises from zero a 16 front to the asviuptotie value im ~0.577., Note that in the horizontal direction the energy rises from zero at the front to the asymptotic value in $\sim 0.5 H$.
 This iniplies at most of the wave euergeyv is contained in a thin laver. Ol: VOL width. following the pulse.," This implies that most of the wave energy is contained in a thin layer, of $\sim 0.5 H$ width, following the pulse."
 But in the vertica direction about of the euergv is contained within a aver with a thickness of ~5/I extending in height from LOM to 15H., But in the vertical direction about of the energy is contained within a layer with a thickness of $\sim 5 H$ extending in height from $10H$ to $15H$.
 The wake therefore coutains a siguificaut Traction of the enerev., The wake therefore contains a significant fraction of the energy.
 Approximately of the enerev is in the pulse itself. i.0.. within a thin laver of ~0.5/7 at the head of the wave.," Approximately of the energy is in the pulse itself, i.e., within a thin layer of $\sim 0.5 H$ at the head of the wave."
 Note also the sharp drop of the enerev with polar anele in Fig. 11.., Note also the sharp drop of the energy with polar angle in Fig. \ref{fig:efrac}.
 Evidently the bulk of the energy goes mto the vertical direction., Evidently the bulk of the energy goes into the vertical direction.
 We have carried out a Linear analysis of the dimenssiounal wave propagation m a stratified. isotlermal atinosphere.," We have carried out a linear analysis of the nal wave propagation in a stratified, isothermal atmosphere."
 The motivation was to study the Πταποια of the generally assumed plane waves., The motivation was to study the limitations of the generally assumed plane waves.
 Although the analysis was based ou linearized equations aud idealized medi our results may suggest the seuse in which 3D aud LD waves differ from one another iu the nonlinear regime., Although the analysis was based on linearized equations and idealized medium our results may suggest the sense in which 3D and 1D waves differ from one another in the nonlinear regime.
 A pulse generated iu the *photosplere’ by a poiut source propagates upward with exponcutial amplification., A pulse generated in the `photosphere' by a `point source' propagates upward with exponential amplification.
 The source region has a Caussian pressure perturbation with a full halfwidth of 0.1«ff. where IT is the xessure scale height. for a diameter of LO kin.," The source region has a Gaussian pressure perturbation with a full half-width of $0.4\times{ H}$, where $H$ is the pressure scale height, for a diameter of 40 km."
 Au initial oerturbation of streneth dép/pzc0.2 becomes nonlinear at a height of LO11., An initial perturbation of strength $\delta p/p \approx 0.2$ becomes nonlinear at a height of $10\times H$.
 So this streueth of the source is near he minima needed for bright poiuts., So this strength of the source is near the minimum needed for bright points.
 At 1 Ahn above the source. which corresponds to he height of formation of the II aud EK lines of Ca II iu th solar chromosphere. the pressure and velocity erturbatiouns reach laree values over a region with a iorizontal extent of 1. Ma.," At 1 Mm above the source, which corresponds to the height of formation of the H and K lines of Ca II in the solar chromosphere, the pressure and velocity perturbations reach large values over a region with a horizontal extent of 1 Mm."
 This is comparable to the size of Ca bright poiuts., This is comparable to the size of Ca bright points.
 The oeutial pulse is followed by a wale in the upward direction at the acoustic cutoff period. approximately three müuutes in the solar atmosphere.," The initial pulse is followed by a wake in the upward direction at the acoustic cutoff period, approximately three minutes in the solar atmosphere."
 There is a wake also in the horizontal direction. with the same period bu much ower almplituce.," There is a wake also in the horizontal direction, with the same period but much lower amplitude."
 The energy of the wave is concentrated iu the vertica direction: One quarter of the upward-propagating energy is contained within a cone with a lalfanuele of 30° about the vertical axis: that cone constitutes only about of the volume ofthe ieuisphlere., The energy of the wave is concentrated in the vertical direction: One quarter of the upward-propagating energy is contained within a cone with a half-angle of $^\circ$ about the vertical axis; that cone constitutes only about of the volume of the hemisphere.
 The euergyis couceutratec also in a narrow laver behind the initial pulse: When wave has reached a height of 1077 above he source. of the euergv on the vertical axis is contained within the first two scale heights behind the apex.," The energyis concentrated also in a narrow layer behind the initial pulse: When the wave has reached a height of $10H$ above the source, of the energy on the vertical axis is contained within the first two scale heights behind the apex."
 The height where ai eiven imaenitude of the perturbation is reached iuereases with he size of the region., The height where a given magnitude of the perturbation is reached increases with the size of the region.
" Assumnius that nonlinear conditious are reached ou the vertical axis at a height of 1077. a cbright point” with a diameter of 1 Mau requires that he wave a 1C edge travel another 1.5/7. which requires au additional 2] s fat a sound speed of 7 kms). and for a diunueter of 2 Alin. the additinal travel distance aud time are 3/7 and 145 s. Thus the ""bright poiut grows from the ceuter outward aud upward."," Assuming that nonlinear conditions are reached on the vertical axis at a height of $10H$, a `bright point' with a diameter of 1 Mm requires that the wave at the edge travel another $1.5H$, which requires an additional 21 s (at a sound speed of 7 km/s), and for a diameter of 2 Mm, the additinal travel distance and time are $3H$ and 45 s. Thus the `bright point' grows from the center outward and upward."
 For the arecr diameter and height. correspoudinely higher lavers m the chromosphere would forma the spectre.," For the larger diameter and height, correspondingly higher layers in the chromosphere would form the spectrum."
 The amplitude of the oscillations dhind the pulse weakens in streugth aud shiriuks in size., The amplitude of the oscillations behind the pulse weakens in strength and shrinks in size.
 As a consequence. the maximal iutensitics associated with brightpoiuts iu the wake are weaker aud smaller than those im the initial pulse.," As a consequence, the maximal intensities associated with brightpoints in the wake are weaker and smaller than those in the initial pulse."
" The maxima in the wake decay with time as c(t)tD ""S7. dependingB on the geometry."," The maxima in the wake decay with time as $v(t)\propto t^{-n/2}$ , depending on the geometry."
 For à plane wave - ina one-dimensional medi. 7= l. the velocity decays as t4?) Git matches the plane wave solution of KRBM: for a line source in a two-dimensional mecimm. the decay is as ft +t: and fora point source in3D. it is as e(f)Xf 32," For a plane wave in a one-dimensional medium, $n=1$ , the velocity decays as $t^{-1/2}$; it matches the plane wave solution of KRBM; for a line source in a two-dimensional medium, the decay is as $t^{-1}$ ; and for a point source in3D, it is as $v(t)\propto t^{-3/2}$ ."
 The, The
true if one takes into account that a strict stability of the mean polarisation appears to be rather unphysical.,true if one takes into account that a strict stability of the mean polarisation appears to be rather unphysical.
 Small fluctuations in the magnetic field configuration and/or a precession of the inner accretion disk seem likely., Small fluctuations in the magnetic field configuration and/or a precession of the inner accretion disk seem likely.
 In this regard. our conclusion that the roughly constant mean polarisation points to a preferred position of the Ser A* system seems reasonable.," In this regard, our conclusion that the roughly constant mean polarisation points to a preferred position of the Sgr A* system seems reasonable."
 The existence of a favored polarization angle allows to investigate the orientation of the Ser Á* system on the sky., The existence of a favored polarization angle allows to investigate the orientation of the Sgr A* system on the sky.
 Here. we want to study it within the orbiting blob model.," Here, we want to study it within the orbiting blob model."
 Meyer et al. (2006a. 2006b))," Meyer et al. \cite{ich2,ich}) )"
 calculated polarimetric lighteurves from an orbiting spot (here we use spot and blob interchangeably) around Ser A* and compared them to observations., calculated polarimetric lightcurves from an orbiting spot (here we use spot and blob interchangeably) around Sgr A* and compared them to observations.
 This showed that this simple model. in which general relativistic effects on the radiation of a somehow confined. locally heated region plays the major role. leads to very good fits of the measurements.," This showed that this simple model, in which general relativistic effects on the radiation of a somehow confined, locally heated region plays the major role, leads to very good fits of the measurements."
 More precisely. in this model the sub-flares are due to a blob on a relativistic orbit around the MBH. while an underlying ring accounts for the broad overall flare.," More precisely, in this model the sub-flares are due to a blob on a relativistic orbit around the MBH, while an underlying ring accounts for the broad overall flare."
 Relativistic effects like beaming. lensing. and change of polarization angle imprint on the emitted intrinsic radiation (e.g. Doveiak et al. 2004:;," Relativistic effects like beaming, lensing, and change of polarization angle imprint on the emitted intrinsic radiation (e.g. Dovciak et al. \cite{dovciak};"
 Connors Stark 1977:: Hollywood Melia 1997:; Bromley et al. 20011:, Connors Stark \cite{connors}; ; Hollywood Melia \cite{hollywood2}; ; Bromley et al. \cite{bromley};
 Falcke et al. 2000::, Falcke et al. \cite{falcke};
 Broderick Loeb 2005.2006:; Schnittman 2005)).," Broderick Loeb \cite{broderick1, broderick2}; Schnittman \cite{schnittman}) )."
 In our model we assume that the variability in the polarization angle and the polarization degree are only due to the relativistic effects., In our model we assume that the variability in the polarization angle and the polarization degree are only due to the relativistic effects.
 As the emitted radiation of Sgr A* is synchrotron radiation (emitted in the disk corona). we assumed two different magnetic field configurations to fit the light curves with our model.," As the emitted radiation of Sgr A* is synchrotron radiation (emitted in the disk corona), we assumed two different magnetic field configurations to fit the light curves with our model."
 The first is analogous to a sunspot and results 1n a constant E-vector perpendicular to the disk., The first is analogous to a sunspot and results in a constant $E$ -vector perpendicular to the disk.
 The second configuration is a global azimuthalmagnetic field that leads to a rotation of the E- along the orbit., The second configuration is a global azimuthalmagnetic field that leads to a rotation of the $E$ -vector along the orbit.
 With this model at hand. observed polarimetric lighteurves can be fitted to investigate the parameters of the Sgr A* system.," With this model at hand, observed polarimetric lightcurves can be fitted to investigate the parameters of the Sgr A* system."
 The inclination angle. the dimensionless spin parameter αν. the brightness excess of the spot with respect to the disk. the initial phase of the spot on the orbit. the orientation of the equatorial plane on the sky. and the polarization degree of the disk and the spot are the free parameters.," The inclination angle, the dimensionless spin parameter $a_\star$, the brightness excess of the spot with respect to the disk, the initial phase of the spot on the orbit, the orientation of the equatorial plane on the sky, and the polarization degree of the disk and the spot are the free parameters."
 In their discussion. Meyer et al. (2006a. 2006b))," In their discussion, Meyer et al. \cite{ich2,ich}) )"
 focused on the viewing angle and the spin parameter of Ser A*., focused on the viewing angle and the spin parameter of Sgr A*.
 In this note we extend the discussion to the position angle of the system. Le. the angle of the equatorial plane normal.," In this note we extend the discussion to the position angle of the system, i.e. the angle of the equatorial plane normal."
 Figure 3 shows the lc-— and 3e —confidence contours in the position angle (8) — inclination angle (/) plane., Figure \ref{conf} shows the $1\sigma-$ and $3\sigma-$ confidence contours in the position angle $\theta$ ) – inclination angle $i$ ) plane.
 The contours are results of the fits to the 2006 data presented in Meyer et al. (2006a))., The contours are results of the fits to the 2006 data presented in Meyer et al. \cite{ich2}) ).
 The magnetic field configuration corresponds to the sunspot case., The magnetic field configuration corresponds to the sunspot case.
 The contours have been calculated such that on each point in the @—- / plane the y--minimum with respect to the other free parameters has been taken., The contours have been calculated such that on each point in the $\theta$ – $i$ plane the $\chi^2$ -minimum with respect to the other free parameters has been taken.
" The only exception ts the dimensionless spin parameter a, Which has been fixed to a,=0.6 throughout the calculations.", The only exception is the dimensionless spin parameter $a_\star$ which has been fixed to $a_\star = 0.6$ throughout the calculations.
 The overall y--minimum lies then at 9= 105°. =55° (marked with the black dot).," The overall $\chi^2$ -minimum lies then at $\theta = 105\degr$ , $i=55\degr$ (marked with the black dot)."
 As the contours show. this mmimum is not unambigous. but it is neverthelessnoteworthythat the formal y--minimum coincides exactly," As the contours show, this minimum is not unambigous, but it is neverthelessnoteworthythat the formal $\chi^2$ -minimum coincides exactly"
were performed and the restoration technique used to correct for atmospheric aberrations.,were performed and the restoration technique used to correct for atmospheric aberrations.
 Ouce the time series are ready to analyze. 23. presents some calculations aud statistics of proper motions ofstructures surrounding the siuuple sunspots.," Once the time series are ready to analyze, \ref{sec:analysis} presents some calculations and statistics of proper motions of structures surrounding the sample sunspots."
 We finally discuss the results aud future work in dL., We finally discuss the results and future work in \ref{sec:dis}.
 The observations were obtained with the (SST.Scharmeretal.200343). on La Γαία between 2003 and 2006.," The observations were obtained with the \citep[SST,][]{SST} on La Palma between 2003 and 2006."
 All observations benefited from the use of the SST adaptive optics system (Scharimereal.900510) hat uiunimized the degrading effects of secius., All observations benefited from the use of the SST adaptive optics system \citep{SSTAO} that minimized the degrading effects of seeing.
 Tage post-processing techniques were applic to increase the homosencitv in the quality of the time series and to enhance the image quality over the whole field-of-view (FOV)., Image post-processing techniques were applied to increase the homogeneity in the quality of the time series and to enhance the image quality over the whole field-of-view (FOV).
 For the 2003 data se we applied Multi-Frzune. Blind Deconvolution (AIFBD) using the iniplemieutation developed by Lófdahlil(2002).., For the 2003 data set we applied Multi-Frame Blind Deconvolution (MFBD) using the implementation developed by \citet{lofdahl02MFBD}.
 The MFBD code was succeeded by the Multi-Objec Multi-Frame Blind Decouvolution code (NOMEDD. VanNoort.RouppevanderVoort&Lófdahlil (2005))}) which clmplovs multiple objects aud phase-diversityv.," The MFBD code was succeeded by the Multi-Object Multi-Frame Blind Deconvolution code (MOMFBD, \citet{MOMFBD}) ) which employs multiple objects and phase-diversity."
 NOMEDD was applied to the data sets after 2003., MOMFBD was applied to the data sets after 2003.
 For all the time series we present im this paper. the seeing conditions were eenerallv very good and sometimes excellent.," For all the time series we present in this paper, the seeing conditions were generally very good and sometimes excellent."
 A large raction of the restored tages in the different time series approach the diffraction lut of the telescope., A large fraction of the restored images in the different time series approach the diffraction limit of the telescope.
 After nuage restoration. we applied standard echniques to the time series including correction or the diurual feld rotation. rigid aliguiment aud cle-stretching to correct for seciugdnduced tage warping.," After image restoration, we applied standard techniques to the time series including correction for the diurnal field rotation, rigid alignment and de-stretching to correct for seeing-induced image warping."
" A panodes filter was also applied to the series (threshold ohase velocity 1 4),", A p-modes filter was also applied to the series (threshold phase velocity 4 $^{-1}$ ).
 Table 1 preseuts details for the different active region argets. Table 2. gives some details ou the restored time series," Table \ref{table1} presents details for the different active region targets, Table \ref{table2} gives some details on the restored time series."
 Below we provide more detailed iuforiuatiou ou he different data scts., Below we provide more detailed information on the different data sets.
 An interference filter centered at the continu at 156.1 nii CEWIIM. 1.1 ni) was used., An interference filter centered at the continuum at 436.4 nm (FWHM 1.1 nm) was used.
 Iu the remaiuder. this filter will be referred to as ουν filter.," In the remainder, this filter will be referred to as “G-cont” filter."
 Every 25 s. the 3 highest contrast nuages were selected aud stored to disk.," Every 25 s, the 3 highest contrast images were selected and stored to disk."
 These 3 images were used for MEDD processiug., These 3 images were used for MFBD processing.
 This sunspot is shown in Figure 2.., This sunspot is shown in Figure \ref{F:2}.
 An interference filter centered at the C-band at 130.5 mm CEWIIM. 1.2 ni) was used., An interference filter centered at the G-band at 430.5 nm (FWHM 1.2 nm) was used.
" In the remaiuder. this filter will be referred to as ""Ce-band filter."," In the remainder, this filter will be referred to as “G-band” filter."
 This observation involved 2 cameras sct-yp as a plasxc-diversity pair: one in focus aud one slightly out of focus., This observation involved 2 cameras set-up as a phase-diversity pair: one in focus and one slightly out of focus.
 20 exposures from cach camera. acquired over the course of 7.5 seconds. were used for MOMEDD processing.," 20 exposures from each camera, acquired over the course of 7.5 seconds, were used for MOMFBD processing."
" During processing of the time series. a few bad quality inages, suffering from too much bluring. were dropped."," During processing of the time series, a few bad quality images, suffering from too much blurring, were dropped."
 The exposure time was 11 ms., The exposure time was 11 ms.
 This sunspot is shown in Figure 3.., This sunspot is shown in Figure \ref{F:3}.
 This time series was observed with the G-cout filter., This time series was observed with the G-cont filter.
 The AIOMFBD processing was done using 10 exposures (acquired diriug 7s) from a plasc-diversitv pair of cameras., The MOMFBD processing was done using 10 exposures (acquired during 7s) from a phase-diversity pair of cameras.
 For the central part of the FOV. centered ou the suuspot. the MOMEDD processing was extended to involve a C-band plase-diversity pair of cameras.," For the central part of the FOV, centered on the sunspot, the MOMFBD processing was extended to involve a G-band phase-diversity pair of cameras."
 After MOMEBD restoration. the best quality image from 3 subsequent images was selected for further processing of the tine series.," After MOMFBD restoration, the best quality image from 3 subsequent images was selected for further processing of the time series."
 The exposure tine was 11.3 ms., The exposure time was 11.3 ms.
 This sunspot is shown in Figure 9.., This sunspot is shown in Figure \ref{F:9}.
 This fine series was processed in a similar ΑΠΟ: as S3., This time series was processed in a similar manner as S3.
 After MOMEDD restorations. the worst of the Huaees were dropped from the time series;," After MOMFBD restorations, the worst of the images were dropped from the time series."
 This suuspot is shown in Figure L., This sunspot is shown in Figure \ref{F:4}.
 For this time series. MFDD processing was applied to LS subsequent C-banud images. acquired. during 10 s. Duages that were too blurred aud when the AO system was not actively compcusating. were dropped from the MEBD sets.," For this time series, MFBD processing was applied to 18 subsequent G-band images, acquired during 10 s. Images that were too blurred and when the AO system was not actively compensating, were dropped from the MFBD sets."
 This sunspot is shown in Figure 5.., This sunspot is shown in Figure \ref{F:5}.
 This time series was observed with an iuterference filter ceutered ou the spectral region between Ca Is aud Ca IT at 395. lum (PWIA 1.0 uni)., This time series was observed with an interference filter centered on the spectral region between Ca K and Ca H at 395.4 nm (FWHM 1.0 nm).
 For the NOMEDBD processing. 20 exposures (acquired daring 8.3 s) frou a phase-diversity pair of cameras were used.," For the MOMFBD processing, 20 exposures (acquired during 8.3 s) from a phase-diversity pair of cameras were used."
 Iu addition. simultaneous exposures frou two other cameras were used for the restoration: one equipped with a uarrow-baud filter ceutered in the Ca II wing at 396.5 nm aud one equipped with a uarrow-hand filter centered on the Ca II core (396.8 um).," In addition, simultaneous exposures from two other cameras were used for the restoration: one equipped with a narrow-band filter centered in the Ca H wing at 396.5 nm and one equipped with a narrow-band filter centered on the Ca H core (396.8 nm)."
 The exposure time was 11 mis., The exposure time was 11 ms.
 This sunspot is shown in Figure 6.., This sunspot is shown in Figure \ref{F:6}.
 This time series was observed with an iuterference filter centered at 630.2 mm (EWIIME 0.8 nmn)., This time series was observed with an interference filter centered at 630.2 nm (FWHM 0.8 nm).
 This is the pre-filter for the Lockheed SOUP. filter., This is the pre-filter for the Lockheed SOUP filter.
" NOMEDD processing was applied to 100 exposures obtained during 19 s The restorations included nuages frou, a plasxe-diversity pair of cameras with the wide band filter. aud narrow band SOUP exposures iu the Fe I 630.2 um spectral line."," MOMFBD processing was applied to 400 exposures obtained during 19 s. The restorations included images from a phase-diversity pair of cameras with the wide band filter, and narrow band SOUP exposures in the Fe I 630.2 nm spectral line."
 This suuspot is shown in Figur 7.., This sunspot is shown in Figure \ref{F:7}.
 The high quality. stability. aud long duration of he restored sunspot series; enable us to follow the dynamics of the plasma around the sunspots.," The high quality, stability and long duration of the restored sunspot series, enable us to follow the dynamics of the plasma around the sunspots."
 We lave computed proper-motion velocity fields (horizontal velocities) cimploving the local correlation tracking (LOT) technique (November&Simon(1988) using he tuplementation of \lolowuy-Tloras&Yi(19913) , We have computed proper-motion velocity fields (horizontal velocities) employing the local correlation tracking (LCT) technique \citet{november1988} using the implementation of \citet{molowny1994}) ).
A Gaussian tracking window of FWIRI 10 suitable or tracking mainly erauules has been used for all the series., A Gaussian tracking window of FWHM $1\farcs0$ suitable for tracking mainly granules has been used for all the series.
 The method produces a sequence of map-pairs each one describing. within a short time interval (tens of 1ulliseconds). the horizoutal velocity components ο and =y along the FOV.," The method produces a sequence of map-pairs each one describing, within a short time interval (tens of milliseconds), the horizontal velocity components $-x$ and $-y$ along the FOV."
 These maps of compoucuts have been averaged over 5 and 10 min periods. aud also over the respective total duration of every sunspot series (nore than £0 1011 in all cases).," These maps of components have been averaged over 5 and 10 min periods, and also over the respective total duration of every sunspot series (more than 40 min in all cases)."
 From the averaged laps of velocity componcuts. the distribution of velocity," From the averaged maps of velocity components, the distribution of velocity"
cannot cause the significant spreads in luminosity (or radius) observed in stellar populations (e.g.???)..,"cannot cause the significant spreads in luminosity (or radius) observed in stellar populations \citep[e.g.][]{tout99,hartmann97,siess99}."
" However, accretion and infall rates during the early, ‘assembly’ phase of star formation are expected to be much larger than those observed for more evolved stars."," However, accretion and infall rates during the early, `assembly' phase of star formation are expected to be much larger than those observed for more evolved stars."
" ? demonstrate that assembly phase accretion can significantly affect the radius of the central star, causing accelerated contraction."," \cite{baraffe09} demonstrate that assembly phase accretion can significantly affect the radius of the central star, causing accelerated contraction."
" Since the star’s radius changes on the thermal timescale, which also governs the contraction of the star towards the main sequence, a star’s radius will carry the imprint of this elevated protostellar accretion for a large fraction of its pre-main-sequence lifetime."," Since the star's radius changes on the thermal timescale, which also governs the contraction of the star towards the main sequence, a star's radius will carry the imprint of this elevated protostellar accretion for a large fraction of its pre-main-sequence lifetime."
 find that apparent age spreads of up to 10 Myr can be produced by accretion rates of ~10~*Mo yyr!.," \cite{baraffe09}
 find that apparent age spreads of up to 10 Myr can be produced by accretion rates of $\sim$ $^{-4}$ $_{\odot}$ $^{-1}$."
 Such extreme accretion rates would cause young stars to be much brighter than observed., Such extreme accretion rates would cause young stars to be much brighter than observed.
" However, ? show thatepisodic accretion events with similar rates can still reproduce the observed luminosity spreads, without causing a luminosity problem."," However, \cite{baraffe09} show that accretion events with similar rates can still reproduce the observed luminosity spreads, without causing a luminosity problem."
 There is growing evidence that accretion onto young stars shows large variations in accretion rate (e.g.?).., There is growing evidence that accretion onto young stars shows large variations in accretion rate \citep[e.g.][]{enoch09}.
" Episodic accretion in not a prerequisite for significant luminosity spreads, but it can simultaneously explain the observed luminosity spreads without causing anomolously high luminosities and is more consistent with our current understanding of star formation and disc evolution (???).."," Episodic accretion in not a prerequisite for significant luminosity spreads, but it can simultaneously explain the observed luminosity spreads without causing anomolously high luminosities and is more consistent with our current understanding of star formation and disc evolution \citep{vorobyov06,vorobyov10,zhu09}."
" Obtaining observational evidence for accretion-induced radius spreads is not a straightforward task, as the accretion history of a young star is not easily determined."," Obtaining observational evidence for accretion-induced radius spreads is not a straightforward task, as the accretion history of a young star is not easily determined."
 One property which may be related to accretion history is the stellar rotation rate., One property which may be related to accretion history is the stellar rotation rate.
" Young stars typically rotate with periods in the range of 1-10 days, which is a fraction of their breakup speed."," Young stars typically rotate with periods in the range of 1–10 days, which is a fraction of their breakup speed."
" Whilst the underlying physical mechanism is not yet clear (see?,foradiscussion),, there is strong observational evidence that the population of stars with discs rotate more slowly than those without disks (e.g.?????).."," Whilst the underlying physical mechanism is not yet clear \citep[see][for a discussion]{matt10}, there is strong observational evidence that the population of stars with discs rotate more slowly than those without disks \citep[e.g.][]{edwards93,herbst00,littlefair05,rebull06,cieza07}."
 The modelling work of ? showed that the period distribution of young stars can be explained if accretion from a disc results in a constant stellar spin period of around 7 days., The modelling work of \cite{rebull04a} showed that the period distribution of young stars can be explained if accretion from a disc results in a constant stellar spin period of around 7 days.
" Thus, the currently accepted view of spin evolution of young stars is that stars which are accreting are braked by the star-disc interaction and rotate slowly."," Thus, the currently accepted view of spin evolution of young stars is that stars which are accreting are braked by the star-disc interaction and rotate slowly."
" Once the disc disperses, or the accretion rate drops below some critical threshold, the star will spin up again as it contracts towards the main-sequence."," Once the disc disperses, or the accretion rate drops below some critical threshold, the star will spin up again as it contracts towards the main-sequence."
 This spin-up is slow; it is a direct result of the contraction of the star and thus proceeds on the thermal timescale., This spin-up is slow; it is a direct result of the contraction of the star and thus proceeds on the thermal timescale.
" Therefore, the present day rotation rate of the star depends upon the accretion history, although the exact relationship will depend on the details of the braking mechanism and be complicated by additional factors such as the stellar magnetic field strength."," Therefore, the present day rotation rate of the star depends upon the accretion history, although the exact relationship will depend on the details of the braking mechanism and be complicated by additional factors such as the stellar magnetic field strength."
" If both luminosity and rotation rates are functions of the accretion history, it is possible that a correlation between rotation and luminosity exists."," If both luminosity and rotation rates are functions of the accretion history, it is possible that a correlation between rotation and luminosity exists."
" Motivated by this argument, we present an analysis of the link between rotation rate and luminosity for four star forming regions; the ONC, NGC 2264, Cep OB3b and NGC 2362."," Motivated by this argument, we present an analysis of the link between rotation rate and luminosity for four star forming regions; the ONC, NGC 2264, Cep OB3b and NGC 2362."
" The ONC, NGC 2264, Cep OB3b and NGC 2362 are all young (€5 Myr) star forming regions in which large numbers of rotation periods have been measured."," The ONC, NGC 2264, Cep OB3b and NGC 2362 are all young $\le 5$ Myr) star forming regions in which large numbers of rotation periods have been measured."
" For the analysis presented here we take rotation periods from ? (ONC), ?? (NGC 2264), ? (NGC 2362) and ? (Cep OB3b)."," For the analysis presented here we take rotation periods from \citealt{herbst02} (ONC), \citealt{lamm04,lamm05} (NGC 2264), \citealt{irwin08} (NGC 2362) and \citealt{littlefair10} (Cep OB3b)."
" We concentrate on the higher mass stars (0.4€M< 1.0M;5) in these clusters, since these stars show the clearest difference in rotation rate between stars with discs and those without (see?,forexample), and show little dependence between rotation rate and stellar mass (e.g?).."," We concentrate on the higher mass stars $0.4 \le {\rm M} <
1.0$ $_{\odot}$ ) in these clusters, since these stars show the clearest difference in rotation rate between stars with discs and those without \cite[see][for example]{cieza07}, and show little dependence between rotation rate and stellar mass \cite[e.g][]{littlefair10}."
" Masses for stars in each association were calculated in a consistent manner, using the extinction and distance-corrected /-band magnitudes and the models of ? - see ? for more details."," Masses for stars in each association were calculated in a consistent manner, using the extinction and distance-corrected $I$ -band magnitudes and the models of \cite{baraffe98} - see \cite{littlefair10} for more details."
" For the ONC, individual extinction values from ? were used, but for the other associations a single extinction value was adopted for all stars."," For the ONC, individual extinction values from \cite{hillenbrand97} were used, but for the other associations a single extinction value was adopted for all stars."
" This is unlikely to be realistic assumption, but fortunately the reddening vector lies approximatelya along the isochrones in a V, V-I diagram and so it is unlikely to affect our analysis."," This is unlikely to be a realistic assumption, but fortunately the reddening vector lies approximately along the isochrones in a V, V-I diagram and so it is unlikely to affect our analysis."
 We divided our data into bright and faint sub-samples using the following method., We divided our data into bright and faint sub-samples using the following method.
" First we calculated the median colour in several magnitude bins, and a first-order polynomial was fitted to define the median magnitude as a function of colour."," First we calculated the median colour in several magnitude bins, and a first-order polynomial was fitted to define the median magnitude as a function of colour."
 Stars with magnitudes brighter than this were assigned to the bright sample and vice-versa., Stars with magnitudes brighter than this were assigned to the bright sub-sample and vice-versa.
 The results are shown in figure 1.., The results are shown in figure \ref{fig:age_sel}.
 We note here that the results in this letter are robust against changes in the method described above., We note here that the results in this letter are robust against changes in the method described above.
" Our results do not change significantly if we fit the median colour in a series of magnitude bins, or the median magnitude in a series of colour bins."," Our results do not change significantly if we fit the median colour in a series of magnitude bins, or the median magnitude in a series of colour bins."
 Nor does the result depend on the order of the polynomial used to fit the median values., Nor does the result depend on the order of the polynomial used to fit the median values.
 The period distributions for our data are shown in figure 2.., The period distributions for our data are shown in figure \ref{fig:pdist_v_age}.
" It is readily apparent to the eye that, within each star forming region, the period distributions of the bright and faint sub-samples differ markedly."," It is readily apparent to the eye that, within each star forming region, the period distributions of the bright and faint sub-samples differ markedly."
" For example, in the ONC, the bright sub-sample lacks the characteristic bi-modal distribution observed in this cluster, showing instead a rising distribution towards short periods."," For example, in the ONC, the bright sub-sample lacks the characteristic bi-modal distribution observed in this cluster, showing instead a rising distribution towards short periods."
" By contrast, the faint sub-sample shows the familiar bi-modal distribution, with a secondary peak of slow rotators at 6—10 days."," By contrast, the faint sub-sample shows the familiar bi-modal distribution, with a secondary peak of slow rotators at 6–10 days."
 We tested the null-hypothesis that the bright and faint sub-samples were drawn from the same parent distribution using a 1- K-S test., We tested the null-hypothesis that the bright and faint sub-samples were drawn from the same parent distribution using a 1-D K-S test.
" The null hypothesis was rejected with 99.8 percent confidence for the ONC, 99.3 percent confidence for NGC 2264, 99 percent confidence for Cep OB3b and 98 percent confidence for NGC 2362."," The null hypothesis was rejected with 99.8 percent confidence for the ONC, 99.3 percent confidence for NGC 2264, 99 percent confidence for Cep OB3b and 98 percent confidence for NGC 2362."
" In all cases, the median period of the faint sub-sample is larger than the median period of the bright sub-sample."," In all cases, the median period of the faint sub-sample is larger than the median period of the bright sub-sample."
 Whilst the period distributions of bright and faint sub-samples overlap, Whilst the period distributions of bright and faint sub-samples overlap
This paper generalizes the results of Paper I by allowing both the point light source and the particle density distribution be functions of arbitrary angular form.,This paper generalizes the results of Paper I by allowing both the point light source and the particle density distribution to be functions of arbitrary angular form.
 As in Paper Ι. we restrict ourselves to optically thin envelopes and single Thomson or Rayleigh scattering mechanisms.," As in Paper I, we restrict ourselves to optically thin envelopes and single Thomson or Rayleigh scattering mechanisms."
 Although the effects of a finite-sized tlluminator. will be important for obtaining quantitatively à more accurate polarization amplitude (e.g.. Ignace. Henson. Carson 2008). the point source approximation Is certainly adequate for exploring polarization trends of generalized source and envelope geometries.," Although the effects of a finite-sized illuminator will be important for obtaining quantitatively a more accurate polarization amplitude (e.g., Ignace, Henson, Carson 2008), the point source approximation is certainly adequate for exploring polarization trends of generalized source and envelope geometries."
 In Sect., In Sect.
 2. we develop the theory required to evaluate the polarization. based on Paper I and the discussion of Simmons (1982).," 2, we develop the theory required to evaluate the polarization, based on Paper I and the discussion of Simmons (1982)."
 Applications to specific cases of astrophysical interest are explored in Sect., Applications to specific cases of astrophysical interest are explored in Sect.
 3. with emphasis given to a treatment of the star and the circumstellar envelope as ellipsoidal.," 3, with emphasis given to a treatment of the star and the circumstellar envelope as ellipsoidal."
 A discussion of these results and the need for ongoing modeling is given in Sect., A discussion of these results and the need for ongoing modeling is given in Sect.
 4., 4.
 As in Paper I. we neglect the effects of finite star depolarization and stellar occultation.," As in Paper I, we neglect the effects of finite star depolarization and stellar occultation."
" Consequently scattering particles in the circumstellar envelope ""see"" a point star.", Consequently scattering particles in the circumstellar envelope “see” a point star.
" However. we model the anisotropy of the illumination with a ""shape"" function to represent directional flux."," However, we model the anisotropy of the illumination with a “shape” function to represent directional flux."
 Under these approximations. our goal is to derive the net polarization from dipole scattering for an uresolved system that allows. for an arbitrary circumstellar envelope geometry anc source geometry.," Under these approximations, our goal is to derive the net polarization from dipole scattering for an unresolved system that allows for an arbitrary circumstellar envelope geometry and source geometry."
 Our approach is to represent the different geometrical factors in terms of expansions in spherical harmonics., Our approach is to represent the different geometrical factors in terms of expansions in spherical harmonics.
 As a result. it is important to be clear about the different coordinate systems being employed.," As a result, it is important to be clear about the different coordinate systems being employed."
 Following the notation of Paper IL. a deseription of the problem requires three coordinate systems. each centered on the star: one for the observer. one for the envelope. and one for the illumination pattern of the star itself.," Following the notation of Paper I, a description of the problem requires three coordinate systems, each centered on the star: one for the observer, one for the envelope, and one for the illumination pattern of the star itself."
 In general. none of these systems are collinear.," In general, none of these systems are collinear."
" The systems are defined as follows: In general the scatterer density is 7G.dy""). and the flux of radiation from the star is FG.9.φ)."," The systems are defined as follows: In general the scatterer density is $n(r,\vartheta',\varphi')$, and the flux of radiation from the star is $F(r,\vartheta,\varphi)$."
 The stellar illumination is taken to be unpolarized., The stellar illumination is taken to be unpolarized.
 In fact. stellar atmospheres may show some low level of intrinsic polarization (Collins 1970): however. the polarization from the circumstellar scattering will be dominated by the Stokes-I component from the source.," In fact, stellar atmospheres may show some low level of intrinsic polarization (Collins 1970); however, the polarization from the circumstellar scattering will be dominated by the Stokes-I component from the source."
 Following Eq. (, Following Eq. (
"1) from Paper I (see also Simmons 1982). the scattered flux and Stokes parameters (δις.Q.U) of the scattered radiation at the Earth (distance D) are given m the form: with the same definitions as in Paper I. with ZZ=ο-iU. forie V-lk =2/2 the wave number. and /, and i» the scattering functions as defined by van de Hulst (1957).","1) from Paper I (see also Simmons 1982), the scattered flux and Stokes parameters $(F_{\rm sc}, Q, U)$ of the scattered radiation at the Earth (distance $D$ ) are given in the form: with the same definitions as in Paper I, with $Z^*=Q-iU$, for $i=\sqrt{-1}$ , $k=2\pi/\lambda$ the wave number, and $i_1$ and $i_2$ the scattering functions as defined by van de Hulst (1957)."
 The circular. polarization V. is assumed zero., The circular polarization $V$ is assumed zero.
 For Thomson (free electrons) or Rayleigh scattering. we have where the value of the cross section factor c is chosen according to whether Thomson scattering (c independent of Κ) or Rayleigh scattering (7o£) is considered.," For Thomson (free electrons) or Rayleigh scattering, we have where the value of the cross section factor $\sigma$ is chosen according to whether Thomson scattering $\sigma$ independent of $k$ ) or Rayleigh scattering $\sigma \propto k^4$ ) is considered."
" Provided that F varies smoothly. it may be expressed in terms of spherical harmonies in the observer frame (see Paper I: where and Ri, represents elements of a rotational matrix (see Paper D)."," Provided that $F$ varies smoothly, it may be expressed in terms of spherical harmonics in the observer frame (see Paper I): where and $R^l_{nm}$ represents elements of a rotational matrix (see Paper I)."
 Similarly. the density distribution of scatterers can also be expanded in terms of spherical harmonics. which becomes where The Y4(8.6) factors in Eqs. (3))," Similarly, the density distribution of scatterers can also be expanded in terms of spherical harmonics, which becomes where The $Y_{l{\rm n}}(\theta,\phi)$ factors in Eqs. \ref{eq:Fexpand}) )"
 and (5)) are as defined in Paper I (c.£..," and \ref{eq:nexpand}) ) are as defined in Paper I (c.f.,"
 Jackson 1975)., Jackson 1975).
" For the product of two spherical harmonies. we can express the multipoles of 56.0.6) and F(.0.0) as: or as In this latter expression. theset of CyM, values are Clebsh- coefficients. arising from the products of two spherical"," For the product of two spherical harmonics, we can express the multipoles of $n(r,\theta,\phi)$ and $F(r,\theta,\phi)$ as: or as In this latter expression, theset of $C^{\rm LM}_{ll'\,{\rm nm'}}$ values are Clebsh-Gordon coefficients, arising from the products of two spherical"
The x-ray luminosity of source D. L4(0.5-2.5 keV) =4.1d0.5x10 for a 1 keV bremsstrahlung spectrum. suggests a quiescent LMXD since CVs generally do not reach this luminosity (Verbunt et 11997).,"The x-ray luminosity of source B, $L_{X}$ (0.5-2.5 keV) $=
4.1\pm0.5\times10^{33}$ for a 1 keV bremsstrahlung spectrum, suggests a quiescent LMXB since CVs generally do not reach this luminosity (Verbunt et 1997)."
 Source D is the only source to show clear x-ray. variability. al the level in the Cramer-Von Mises test (Daniel 1990).," Source B is the only source to show clear x-ray variability, at the level in the Cramer-Von Mises test (Daniel 1990)."
 It is possible that this source is responsible for the Type I bursts seen with ASCA and BeppoSAX (int Zand et 11993. Alukai Simale 2000). since neither telescope could resolve source D from the brisht LMXD. source A. only distant.," It is possible that this source is responsible for the Type I bursts seen with ASCA and BeppoSAX (in't Zand et 1998, Mukai Smale 2000), since neither telescope could resolve source B from the bright LMXB, source A, only distant."
 DMLAO98 identilied star 49 as the bright LMXD counterpart based on U-D excess —0.9. 1.2 mags blueward of the main sequence |MS]). and later (in Deutsch et 22000) added weight to the identification bv discovering 1 magnitude flickering and a possible period of 43.6 minutes.," DMA98 identified star 49 as the bright LMXB counterpart based on U-B excess $U-B=-0.9$ , 1.2 mags blueward of the main sequence [MS]), and later (in Deutsch et 2000) added weight to the identification by discovering $\sim1$ magnitude flickering and a possible period of 43.6 minutes."
 We confirm the variabiliiv. al 24 6(V) and 21 6(4). ancl identify (he same possible optical period. but with an FAP of 0.027. the least significant FAP of the power spectrum peaks for the three optical counterparts.," We confirm the variability, at 24 $\sigma (V)$ and 21 $\sigma (I)$, and identify the same possible optical period, but with an FAP of 0.027, the least significant FAP of the power spectrum peaks for the three optical counterparts."
" If the 43.6-minute period is the real binary. period. a main sequence companion is ruled out since the implied orbital separation for a 2 441. svstem is 3.6x10""10 cm. less than the radius of even an M2 star. implying a double-degenerate svstem."," If the 43.6-minute period is the real binary period, a main sequence companion is ruled out since the implied orbital separation for a 2 $\Msun$ system is $3.6\times10^{10}$ cm, less than the radius of even an M2 star, implying a double-degenerate system."
 However. this identification sulfers a problem in (hat star 49 falls onto the AIS in V—7 (see Figure 2. Table 1).," However, this identification suffers a problem in that star 49 falls onto the MS in $V-I$ (see Figure 2, Table 1)."
 The svslen also lies on the MIS in an instrumental V. vs. B-V. CMD from further archival analvsis., The system also lies on the MS in an instrumental $V$ vs. $B-V$ CMD from further archival analysis.
 The AIS colors suggest that lisht from a main-sequence secondary is dominant in V and J. and that the weak accretion disk provides the U excess (see DMA98).," The MS colors suggest that light from a main-sequence secondary is dominant in $V$ and $I$, and that the weak accretion disk provides the $U$ excess (see DMA98)."
 HLowever. this interpretation cannot explain why both V. and J vary in tandem by one magnitude (Deutsch et 22000. Figure 1) without invoking eclipses. nor does it explain the laree llickering (which generally implies a bright. blue disk).," However, this interpretation cannot explain why both $V$ and $I$ vary in tandem by one magnitude (Deutsch et 2000, Figure 1) without invoking eclipses, nor does it explain the large flickering (which generally implies a bright, blue disk)."
 We considered (he possibility that the identification suffers source confusion (background or foreground stars contributing V. 7 light). but there is no evidence for blending in the IIST image.," We considered the possibility that the identification suffers source confusion (background or foreground stars contributing $V$, $I$ light), but there is no evidence for blending in the HST image."
 We consider it most likely that the 42.6 minute period is spurious. but (he lana variability remains a mystery.," We consider it most likely that the 43.6 minute period is spurious, but the 1-mag variability remains a mystery."
 The suggested counterpart of source C lies only 3° from the cluster center. a—138:35:45.66. 0—-32:59:26.5 (J2000). from Grindlay et 22001 (in preparation).," The suggested counterpart of source C lies only 3"" from the cluster center, $\alpha$ =18:35:45.66, $\delta$ =-32:59:26.5 (J2000), from Grindlay et 2001 (in preparation)."
 This is within one core racius(42... Harris 1996).," This is within one core radius, Harris 1996)."
 Ws V—1 color. Mq and Fy/Εν (see Table 1. Figure 2) suggest a CV. since light [rom field CVs is usually dominated by the hot accretion disk throughout," Its $V-I$ color, $M_{V}$, and $F_{X}/F_{V}$ (see Table 1, Figure 2) suggest a CV, since light from field CVs is usually dominated by the hot accretion disk throughout"
mused varied by arouud | ffor the perturbed models.,mixed varied by around 1 for the perturbed models.
 The velocity distribution of chemical isotopes is shown iu Figure 22.., The velocity distribution of chemical isotopes is shown in Figure \ref{velvsel}. .
 ! Ti was mixed to relatively high velocities iu our solar metallicity stars;, $^{44}$ Ti was mixed to relatively high velocities in our solar metallicity stars.
 The solar stars show aand at high velocitiesof d«10? kin s|., The solar stars show and at high velocitiesof $4\times10^3$ km $^{-1}$.
 These isotopes are also mixed all the way to the core of the solar iietallicitv stars. and some fraction of hem reaches neeative velocities of less than 0.5«10? Jans 4.," These isotopes are also mixed all the way to the core of the solar metallicity stars, and some fraction of them reaches negative velocities of less than $-0.5 \times
10^3$ km $^{-1}$."
 In Model $15À. where the most mixing takes ulace. we see 160 aud ?C iixed out to 2«10° lan loda velocity space.," In Model s15A, where the most mixing takes place, we see $^{16}$ O and $^{12}$ C mixed out to $2 \times 10^3$ km $^{-1}$ in velocity space."
 does uot reach the high velociies observed in L9STA. Model «ΙΑ shows a peal iu the velocity at ~O lan +. with a tail extending to {τς10° lan +.," does not reach the high velocities observed in 1987A. Model s15A shows a peak in the velocity at $\approx 0$ km $^{-1}$, with a tail extending to $0.7 \times 10^3$ km $^{-1}$ ."
" Model s254 shows slightly higher veocities., with the distribution peaking at around θες10° kins+ and reaching out to 0.9«10° kins 1."," Model s25A shows slightly higher velocities, with the distribution peaking at around $0.4 \times 10^3$ km $^{-1}$ and reaching out to $0.9\times 10^3$ km$^{-1}$ ."
 These hierer velocities for are probably due to smaller zuiomus of fallbackin this star., These higher velocities for are probably due to smaller amounts of fallbackin this star.
"Following the collapse the dynamical response of the cluster by mass segregation and expansion leads to the formation of a relaxed stellar cusp inside rj, with ng(r)xr-?8 (og=7/4-5/2) for the massive stars, and np(r)οςr-9/? for the low-mass stars2010).","Following the collapse the dynamical response of the cluster by mass segregation and expansion leads to the formation of a relaxed stellar cusp inside $r_{h}$, with $n_{H}(r)\propto r^{-\alpha_{H}}$ $\alpha_{H}=7/4$ $5/2$ ) for the massive stars, and $n_{L}(r)\propto r^{-3/2}$ for the low-mass stars."
". The cusp extends inward down to the collision radius, rin=max(rz,rca)Τοοιv at which 2-body relaxation ceases to be effective because(/), velocities are so high that only physical collisions can substantially change the stellar orbits1976)."," The cusp extends inward down to the collision radius, $r_{in}=\max(r_{t},r_{\mathrm{coll}})=r_{\mathrm{coll}}\sim(\Mbh/\Ms)\Rs$ , at which 2-body relaxation ceases to be effective because velocities are so high that only physical collisions can substantially change the stellar orbits."
". For a standard MS mass-radius relation this gives so that for the cusp to extend down to a,~22 AU the stars must be more massive than ~3.7Mo."," For a standard MS mass-radius relation this gives so that for the cusp to extend down to $a_{\star}\simeq22$ AU the stars must be more massive than $\sim3.7\, M_{\odot}$."
" One notices that this lower limit coincides with the mass range predicted by pre-IMBH mass segregation, while being low enough to include AGB progenitors."," One notices that this lower limit coincides with the mass range predicted by pre-IMBH mass segregation, while being low enough to include AGB progenitors."
 Equilibrium is established when the flux of gravitational binding energy that is released when stars are destroyed by the IMBH equals the flux carried by the expanded cluster core2007)., Equilibrium is established when the flux of gravitational binding energy that is released when stars are destroyed by the IMBH equals the flux carried by the expanded cluster core.
". Assuming a Plummer initial distribution one obtains the mass of the cusp Mi as In contrast to the radius of influence of a supermassive black hole (SMBH) in an approximately isothermal galactic nucleus, that of an IMBH contains only a very small number of stars, Nj=My(),,» O(10)."," Assuming a Plummer initial distribution one obtains the mass of the cusp $M_{h}$ as In contrast to the radius of influence of a supermassive black hole (SMBH) in an approximately isothermal galactic nucleus, that of an IMBH contains only a very small number of stars, $N_{h}=M_{h}/\left\langle \Ms\right\rangle _{h}\sim{\cal O}(10)$ ."
" Equations (12)) and (14)), together with the system parameters , /M., azz, () and (04,0) fully describe the IMBH cusp structure and properties in this simplified model, and allow calculating the mean number of stars on orbits with a<a, and e>ex, (Nae)= εξ), where an isotropic cusp is assumed."," Equations \ref{e:tr12}) ) and \ref{e:Mh}) ), together with the system parameters $\Mbh$, $\Mbh/M_{c}$, $\alpha_{H},$ $\left\langle \Ms\right\rangle $ and $\left\langle \Ms\right\rangle _{h}$ $\left\langle \Rs\right\rangle _{h}$ ) fully describe the IMBH cusp structure and properties in this simplified model, and allow calculating the mean number of stars on orbits with $a\le a_{\star}$ and $e\ge e_{\star}$, $\left\langle N_{a,e}\right\rangle =N_{h}[(5/4)a_{\star}/r_{h}]^{3-\alpha_{H}}(1-e_{\star}^{2})$ , where an isotropic cusp is assumed."
" The Poisson probability for having at least one star on a donor orbit is then P,=1—exp[—(Na,e)]-"," The Poisson probability for having at least one star on a donor orbit is then $P_{1}=1-\exp[-\left\langle N_{a,e}\right\rangle ]$."
 The cusp mass and its stellar density rise with the IMBH mass (Eq. 14))., The cusp mass and its stellar density rise with the IMBH mass (Eq. \ref{e:Mh}) ).
" As a consequence, Pj increases with , but with it also the rate of destructive stellar collisions, and the drain rate scatterings into the IMBH, 2004))."," As a consequence, $P_{1}$ increases with $\Mbh$, but with it also the rate of destructive stellar collisions, and the drain rate (star-star scatterings into the IMBH, )."
" (star-starOver the relevant range of IMBH masses, P, rises from €107? to >0.1."," Over the relevant range of IMBH masses, $P_{1}$ rises from $\lesssim10^{-3}$ to $\gtrsim0.1$."
" The IMBH mass that maximizes Pj subject to both the collisional and drain constraints lies in the range of fewx10°, with maxPj~fewx 0.01."," The IMBH mass that maximizes $P_{1}$ subject to both the collisional and drain constraints lies in the range of $\mathrm{few}\times10^{3}\,\Mo$, with $\max P_{1}\sim\mathrm{few\times0.01}$ ."
" The observed residual UV fluxes from HLX-1, after subtracting a disk model and correcting for extinction (Ep..y= 0.042), are 4.8x107ergem?s! (νυν=5.2x10??ergs! for isotropic emission at D=95 Mpc) in the NUV (2147A—3467 A), and 4.7x1075ergcm?s! (Lruy=5.1xl09?9ergs-! ) in the FUV (1233À—1821 A) (based on a preliminary analysis of HST data, Farrell et al.,"," The observed residual UV fluxes from HLX-1, after subtracting a disk model and correcting for extinction $E_{B-V}=0.042$ ), are $4.8\times10^{-15}\,\mathrm{erg\, cm^{-2}\, s^{-1}}$ $L_{NUV}=5.2\times10^{39}\,\mathrm{erg\, s^{-1}}$ for isotropic emission at $D=95\,\mathrm{Mpc}$ ) in the NUV $2147\,\text{\AA}\,-\, 3467\,{\text{\AA}}$ ), and $4.7\times10^{-15}\,\mathrm{erg\, cm^{-2}\, s^{-1}}$ $L_{FUV}=5.1\times10^{39}\mathrm{erg\, s^{-1}}$ ) in the FUV $1233\,\text{\AA}\,-\, 1821\,\text{\AA}$ ) (based on a preliminary analysis of HST data, Farrell et al.,"
 in preparation)., in preparation).
 These place constraints on the properties of the hypothesized birth cluster of the IMBH., These place constraints on the properties of the hypothesized birth cluster of the IMBH.
" Stellar population synthesis models together with model atmospheres allow to predict the UV flux as function of the IMF, metallicity and age of the system."," Stellar population synthesis models together with model atmospheres allow to predict the UV flux as function of the IMF, metallicity and age of the system."
" Generally, the older the cluster, the less UV it emits."," Generally, the older the cluster, the less UV it emits."
" An older cluster directly implies a lower mass AGB progenitor, that can spend longer on the MS."," An older cluster directly implies a lower mass AGB progenitor, that can spend longer on the MS."
" Preliminary modeling, both assuming black body spectra, and using detailed stellar atmosphere models (the code, 1998)) indicates that a minimal cluster age of ~(0.3—0.6)Gyr yr is required for the cluster's UV luminosity to fall below the residuals (Assuming a cluster mass of 5x109 and solar "," Preliminary modeling, both assuming black body spectra, and using detailed stellar atmosphere models (the code, ) indicates that a minimal cluster age of $\sim(0.3-0.6)\,\mathrm{Gyr}$ yr is required for the cluster's UV luminosity to fall below the residuals (Assuming a cluster mass of $5\times10^{6}\,\Mo$ and solar metallicity)."
This corresponds to an AGB progenitor ofmetallicity). ~(2.7—3.5).," This corresponds to an AGB progenitor of $\sim(2.7-3.5)\,\Mo$."
" Such longer-lived progenitors are more susceptible to collisional destruction and scattering into the IMBH, but are still probable at the Pj~fewx0.01 level."," Such longer-lived progenitors are more susceptible to collisional destruction and scattering into the IMBH, but are still probable at the $P_{1}\sim\mathrm{few\times0.01}$ level."
" 'The constraints on the AGB progenitor mass could be relaxed somewhat by assuming a higher value of /M., and possibly by assuming a different metallicity."," The constraints on the AGB progenitor mass could be relaxed somewhat by assuming a higher value of $\Mbh/M_{c}$, and possibly by assuming a different metallicity."
 This requires a more systematic study., This requires a more systematic study.
" After mass transfer starts in the AGB phase, 2-body perturbations from other stars can be neglected, since the relaxation time at rj is tj~fewx109 yr, and the timescale to significantly affect the orbital eccentricity, (1—e?)t,, is longer than the ~10* yr maximal lifetime of the donor, for all relevant values of the eccentricity."," After mass transfer starts in the AGB phase, 2-body perturbations from other stars can be neglected, since the relaxation time at $r_h$ is $t_h \sim \mathrm{few} \times 10^6$ yr, and the timescale to significantly affect the orbital eccentricity, $(1-e_\star^2) t_h$, is longer than the $\sim 10^4$ yr maximal lifetime of the donor, for all relevant values of the eccentricity."
" Thus, while an IMBH with ~104 is not excluded in the context of the eccentric donor scenario, observations and theoretical considerations favor a somewhat less massive IMBH with ~fewx10°, still consistent with the constraints derived from the observed accretion emission (see Sec. 1))."," Thus, while an IMBH with $\Ms\sim10^{4}\,\Mo$ is not excluded in the context of the eccentric donor scenario, observations and theoretical considerations favor a somewhat less massive IMBH with $\Mbh\sim\mathrm{few}\times10^{3}\,\Mbh$, still consistent with the constraints derived from the observed accretion emission (see Sec. \ref{s:intro}) )."
" 'To summarize, we used a simplified model, based on results from N-body simulations of runaway merger formation of a seed IMBH, to predict the stellar distribution around a newly formed IMBH."," To summarize, we used a simplified model, based on results from $N$ -body simulations of runaway merger formation of a seed IMBH, to predict the stellar distribution around a newly formed IMBH."
 We estimated the probability of finding an AGB progenitor on an orbit that could explain the long period variability of HLX- in terms of mass transferfrom an eccentric evolved donor., We estimated the probability of finding an AGB progenitor on an orbit that could explain the long period variability of HLX-1 in terms of mass transferfrom an eccentric evolved donor.
 We find that between 1/100 to 1/10 of IMBHs with could have a fewx1 MS star evolve to an AGB <10*while avoiding collisional destruction or being scattered into the IMBH by 2-body encounters.," We find that between $1/100$ to $1/10$ of IMBHs with $\Mbh\mathrm{\lesssim10^{4}\,\Mo}$ could have a $\mathrm{few\times}\,1\,\Mo$ MS star evolve to an AGB while avoiding collisional destruction or being scattered into the IMBH by 2-body encounters."
 An important caveat is that the validity of the model andour, An important caveat is that the validity of the model andour
physical model. this inner accretion flow should be similar to the inner porGon of a much nore extended accretion flow.,"physical model, this inner accretion flow should be similar to the inner portion of a much more extended accretion flow."
 We use outflow boundary conclitions at both the inner and outer boundaries. and (he usual polar boundary conditions at 0= and 9=π.," We use outflow boundary conditions at both the inner and outer boundaries, and the usual polar boundary conditions at $\theta = 0$ and $\theta = \pi$."
 We integrate or 2000€Mο)e12hr. or 8 orbital periods at. μις. (," We integrate for $2000 G \MBH/c^3
\simeq 12\hr$, or 8 orbital periods at $R_{\rm max}$. ("
and similar codes) fail if p or 4 are small in comparison to the kinetic and magnetic energy densities. or the clensity in nearby zones.,"and similar codes) fail if $\rho$ or $u$ are small in comparison to the kinetic and magnetic energy densities, or the density in nearby zones."
" To prevent (his we impose a hare ""floor."" so that pod0!pUUGM)9? and u10PpuucU/CM)513."," To prevent this we impose a hard “floor,” so that $\rho >
10^{-4} \rho_{\rm max} (R/\Rg)^{-3/2}$ and $u > 10^{-6} \rho_{\rm max}
c^2 (R/\Rg)^{-5/2}$."
" To ""observe"" the numerical model we must specily (he observers distance D. aud the inclination / of the black hole spin to the line of sight.", To “observe” the numerical model we must specify the observer's distance $D$ and the inclination $i$ of the black hole spin to the line of sight.
 Because the dynamical simulation is scale-Iree but the radiative transfer calculation is not. we need to specilv the simulation length unit £=CM/c. time unit 7=CM/c*. and mass unit M (equivalently: the mass accretion rate).," Because the dynamical simulation is scale-free but the radiative transfer calculation is not, we need to specify the simulation length unit $\lunit = GM/c^2$, time unit $\tunit = G M/c^3$, and mass unit $\munit$ (equivalently: the mass accretion rate)."
 Since we set (he peak density in the GRMIID model to 1 in simulation units. the peak density is M£ in ces units.," Since we set the peak density in the GRMHD model to $1$ in simulation units, the peak density is $\munit/\lunit^3$ in cgs units."
 The mass unit is set by AL (vhich. appears in the dvnamical model only in the combination GA) because the flow mass is «&Af. and has negligible effect on the gravitational field.," The mass unit is set by $M$ (which appears in the dynamical model only in the combination $GM$ ) because the flow mass is $\ll M$, and has negligible effect on the gravitational field."
 M is therefore a free parameter., $\munit$ is therefore a free parameter.
 To calculate the SED we use the relativistic Monte Carlo schemegrmonty., To calculate the SED we use the relativistic Monte Carlo scheme.
. A detailed description of the aleorithm and tests are eiven in Dolenceetal.(2009)., A detailed description of the algorithm and tests are given in \cite{dolence:2009}.
. The code fully accounts For svnchrotvon emission and absorption. and Compton scattering.," The code fully accounts for synchrotron emission and absorption, and Compton scattering."
" It uses a ""stationary flow” approximation. computing the spectrum through each Gime slice of simulation data as if il were time-independent."," It uses a “stationary flow” approximation, computing the spectrum through each time slice of simulation data as if it were time-independent."
 This is an approximation because the light. crossing time is comparable to (he clvnamical time., This is an approximation because the light crossing time is comparable to the dynamical time.
 It is done because tracking photons (through the time-dependent simulation data is still too computationally expensive., It is done because tracking photons through the time-dependent simulation data is still too computationally expensive.
 We will evaluate the quality of this approximation once we are able to caleulate fullv sell-consistent spectra., We will evaluate the quality of this approximation once we are able to calculate fully self-consistent spectra.
 An average spectrum is formed by averaging over 50 single slice spectra Irom each of 4 different realizations of the simulation (the realizations differ in the random nuniber seed. used to generate initial conditions)., An average spectrum is formed by averaging over $50$ single slice spectra from each of $4$ different realizations of the simulation (the realizations differ in the random number seed used to generate initial conditions).
 To image themodel we use the relativistic rav-tracing codeibothros.. which accounts onlv for svnchrotron emission and absorption. again using a stationary [flow approximation (seeNobleetal.2007).," To image themodel we use the relativistic ray-tracing code, which accounts only for synchrotron emission and absorption, again using a stationary flow approximation \citep[see][]{noble:2007}."
. An average image is created using the same averaging procedureas for the spectra., An average image is created using the same averaging procedureas for the spectra.
 To sum up. the model parameters (aside [rom (hose (hat describe (he initial conditions) are AL. D. i eus TATUS aud ," To sum up, the model parameters (aside from those that describe the initial conditions) are $M$, $D$, $i$, $a_*$, $\Trat$, and $\munit$."
AM and D are set by (he observations of stellar orbits: we briefly explore the consequence of varving them below., $M$ and $D$ are set by the observations of stellar orbits; we briefly explore the consequence of varying them below.
" Vf we will set for each model by requiring that Gime-averaged [lux £,(230GlIz)= 3.4Jv."," $\munit$ we will set for each model by requiring that time-averaged flux $F_\nu(230\GHz) =
3.4\Jy$ ."
 The remaining three [ree parameters are ας.," The remaining three free parameters are $a_*$ ,"
When working with experiueutal CMD data. we need to remove a portion of the sky due to foreground residuals.,"When working with experimental CMB data, we need to remove a portion of the sky due to foreground residuals."
 In this paper. we operate with several masks. most notably the WMÁAP KQs5 sky cut (Coldetal. 2008).. which removes of the sky.," In this paper, we operate with several masks, most notably the WMAP KQ85 sky cut \citep{gold:2008}, which removes of the sky."
 The edges of the mask need special consideration. due to the differentiation of the map.," The edges of the mask need special consideration, due to the differentiation of the map."
 We therefore perform the following procedure for eeneratiug suitable CAIB masks: We now present a method for estimating how mnmuch a CAMB data sot. deviates from Caussianitv. based ou standard statistical tests.," We therefore perform the following procedure for generating suitable CMB masks: We now present a method for estimating how much a CMB data set deviates from Gaussianity, based on standard statistical tests."
 We calculate the covariance matrix C from the fraction of hills. lakes aud saddles for 10000. Gaussian simulations.," We calculate the covariance matrix $\mathbf C$ from the fraction of hills, lakes and saddles for $40 000$ Gaussian simulations."
" We implement a standard -test on the form where d=Ny,CXjp, aud Ny, is a hill. lake or saddle density at threshold 7;."," We implement a standard $\chi^2$ -test on the form where $\mathbf d = X_{T_t} - \langle X \rangle_{T_t}$ and $X_{T_t}$ is a hill, lake or saddle density at threshold $T_t$."
 We then estimate the (P ou 10000 sinmlations. aud obtain a histocram of the distribution. as presented in Fieure [..," We then estimate the $\chi^2$ on $10 000$ simulations, and obtain a histogram of the distribution, as presented in Figure \ref{fig:chisq}."
 When performing the analvsis on WALAP data.we first calculate the \? and then compare it to the pre-calculated 47. distribution.," When performing the analysis on WMAP data,we first calculate the $\chi^2$ and then compare it to the pre-calculated $\chi^2$ distribution."
 We then count the percentage of the \?s that are above the experimental. ic. if 32% of the simulated. 4? values are larvecr than the 47 from experiueutal data this corresponds to la deviation from the Catssian expectation.," We then count the percentage of the $\chi^2$ s that are above the experimental, i.e. if $32 \%$ of the simulated $\chi^2$ values are larger than the $\chi^2$ from experimental data this corresponds to $1 \sigma$ deviation from the Gaussian expectation."
 A value of 5% would be cousisteut with a 28 deviation from the Caussiauity., A value of $5 \%$ would be consistent with a $2 \sigma$ deviation from the Gaussianity.
 As an alternative method we also use a diagonal covariance matrix for calculating the \? values. thus ignoring correlations between threshold levels.," As an alternative method we also use a diagonal covariance matrix for calculating the $\chi^2$ values, thus ignoring correlations between threshold levels."
 This is the method applied in Hauseuetal.(2001).., This is the method applied in \cite{hansen:2004b}.
 We consider the publicly available 5-vear WALAP data (Bennettetal.2003:Uhushaw2007). that can be obtained from the site.," We consider the publicly available 5-year WMAP data \citep{bennett:2003, hinshaw:2007} that can be obtained from the site."
 We primarily perform the analyses on a combined Voand W frequeucv band (91 aud 61 GIIZ) in order to keep noise and foregrounds to a münnmuimuu., We primarily perform the analyses on a combined V and W frequency band (94 and 61 GHz) in order to keep noise and foregrounds to a minimum.
" Ii addition to the S-vr WMAP data, we also simulate 50000 Gaussian ΟΔΟ maps based on the best-fit angular powerspectriun from the S-vear WAIAP release."," In addition to the 5-yr WMAP data, we also simulate $50 000$ Gaussian CMB maps based on the best-fit angular powerspectrum from the 5-year WMAP release."
 We convolve the simulated maps with the combined V aud WW dustruniental beam before adding Cassia noise., We convolve the simulated maps with the combined V and W instrumental beam before adding Gaussian noise.
 All CMD maps considered in this paper have Healpix resolution Nua=512., All CMB maps considered in this paper have Healpix resolution $N_{\textrm{side}}=512$.
 We focus on several different divisions of the sk., We focus on several different divisions of the sky.
 First. we cousider the full-sky WMAP data includiug the extended IKQs5 aud the IKQ?5 mask from Section 2.2..," First, we consider the full-sky WMAP data including the extended KQ85 and the KQ75 mask from Section \ref{sec:extended_mask}."
 We then investigate the northern aud southern galactic hemispheres individually. using only the NOQs5 mask.," We then investigate the northern and southern galactic hemispheres individually, using only the KQ85 mask."
 Due to the alieumieut of the WALAP satellite with the ecliptic plane. we also include an analysis of the northeru and southern ecliptic hemispheres.," Due to the alignment of the WMAP satellite with the ecliptic plane, we also include an analysis of the northern and southern ecliptic hemispheres."
 We then perform the full-esky analysis on a co-added Q-band data set. a co- V-baud set aud a co-added W-hancl set.," We then perform the full-sky analysis on a co-added Q-band data set, a co-added V-band set and a co-added W-band set."
 Finally. we re-analyze the co-added V|W data sets on hemisplieres centered around all the pixel ceuters on a map with Nuage= 2.," Finally, we re-analyze the co-added V+W data sets on hemispheres centered around all the pixel centers on a map with $N_{\textrm{side}}=2$ ."
 From these results. we obtain maps with preferred directions depicting the amount of deviation from Ciaussianity on the sk.," From these results, we obtain maps with preferred directions depicting the amount of deviation from Gaussianity on the sky."
We thank Luke Dones for a careful and helpful review.,We thank Luke Dones for a careful and helpful review.
 This research was supported by a grant from NASA’s Outer Planets Research program., This research was supported by a grant from NASA's Outer Planets Research program.
multiplied by the differential number counts down to the completeness limit of the survey.,multiplied by the differential number counts down to the completeness limit of the survey.
" Following the work of Madau&Pozzetti(2000),, we compute the EBL in the U band, J, measured in ergs1αιHz!Sr7}, using the following method: where we use the following correction from Vega to AB: Όλη=UVega+0.86."," Following the work of \cite{mp2000}, we compute the EBL in the U band, $I_{\nu}$ measured in $erg~s^{-1}~cm^{-2}~Hz^{-1}~Sr^{-1}$, using the following method: where we use the following correction from Vega to AB: $U_{AB}=U_{Vega}+0.86$."
 Fig.2 shows the EBL obtained by the LBC deep number counts and compare it with the results of Madau&Pozzetti(2000)., \ref{fig:ebl} shows the EBL obtained by the LBC deep number counts and compare it with the results of \cite{mp2000}.
". Our data agrees well with previous estimates of the EBL in the U band, and indeed it allows to derive precisely the peak of the differential EBL, at U=23.55+0.25, and put strong constraints to the contributions of faint galaxies down to U—27.0 to the integrated EBL in this band."," Our data agrees well with previous estimates of the EBL in the U band, and indeed it allows to derive precisely the peak of the differential EBL, at $U=23.55\pm0.25$, and put strong constraints to the contributions of faint galaxies down to $U=27.0$ to the integrated EBL in this band."
 At magnitudes U>24 our estimate of the EBL is significantly larger than the one of Madau&Pozzetti(2000)., At magnitudes $U\ge 24$ our estimate of the EBL is significantly larger than the one of \cite{mp2000}.
". Checking in detail Fig.1,, it is clear that the HDFS is slightly underdense with respect to the HDFN or other deep pencil beam surveys, probably due to cosmic variance effects, and this is reflected on the smaller EBL of Madau&Pozzetti(2000) compared to the LBC estimate."," Checking in detail \ref{fig:lognsu}, it is clear that the HDFS is slightly underdense with respect to the HDFN or other deep pencil beam surveys, probably due to cosmic variance effects, and this is reflected on the smaller EBL of \cite{mp2000} compared to the LBC estimate."
" The integrated galaxy contribution to the EBL in the UV band results vI(v)=3.428+0.068nW/m?/Sr at the effective wavelength of the LBC U-BESSEL filter,3590A,, taking into account the magnitude range U=17-"," The integrated galaxy contribution to the EBL in the UV band results $\nu I(\nu)=3.428\pm 0.068 nW/m^2/Sr$ at the effective wavelength of the LBC U-BESSEL filter, taking into account the magnitude range U=17-27."
 This estimate is higher than the one estimated by Madau&Pozzetti(2000) (2.87*025nW/m?/ Sr) although their value is still consistent with our measure due to their large uncertainties., This estimate is higher than the one estimated by \cite{mp2000} $2.87^{+0.58}_{-0.42}nW/m^2/Sr$ ) although their value is still consistent with our measure due to their large uncertainties.
" Moreover, our estimate reduces the uncertainties by an order of magnitudes, which is of fundamental importance to put strong constraints to galaxy evolution models."," Moreover, our estimate reduces the uncertainties by an order of magnitudes, which is of fundamental importance to put strong constraints to galaxy evolution models."
" Extrapolating the observed number counts in the U band down to flux=0 we derive an integrated EBL of vI(v)=3.727+0.084nW/m?/Sr, so in our survey to U=27 we are resolving 92% of EBL produced by galaxies."," Extrapolating the observed number counts in the U band down to flux=0 we derive an integrated EBL of $\nu I(\nu)=3.727\pm 0.084nW/m^2/Sr$, so in our survey to U=27 we are resolving $92\%$ of EBL produced by galaxies."
" The integrated flux from resolved galaxies, however, should be considered as a lower limit to the total EBL in the universe, since our estimate could be affected by different systematics:"," The integrated flux from resolved galaxies, however, should be considered as a lower limit to the total EBL in the universe, since our estimate could be affected by different systematics:"
where and the phase angle a: Finally. the flux of the reflected light from the planet at the orbital phase é follows from equations 1. and 2:: The planet orbiting its host star produces not only a flux variation (Equation 5)). but also a Doppler shift of the stellar spectrum reflected from the planet.,"where and the phase angle $\alpha$: Finally, the flux of the reflected light from the planet at the orbital phase $\phi$ follows from equations \ref{equ:2} and \ref{equ:3}: The planet orbiting its host star produces not only a flux variation (Equation \ref{equ:4}) ), but also a Doppler shift of the stellar spectrum reflected from the planet."
 The RV semi-amplitude Ky of that shift depends on the orbital inclination £. which is unknown for non-transiting planets.," The RV semi-amplitude $K_{\rm p}$ of that shift depends on the orbital inclination $i$ , which is unknown for non-transiting planets."
" A, can be expressed by where MyK, is the RV semi-amplitude of the star. and M, and M,sin are the stellar mass and the minimum mass of the planet. respectively."," $K_{\rm{p}}$ can be expressed by where $K_{\rm{s}}$ is the RV semi-amplitude of the star, and $M_{\rm{s}}$ and $M_{\rm{p}} \sin i$ are the stellar mass and the minimum mass of the planet, respectively."
 The largest possible amplitude Kj. occurs at the orbital inclination /=907.," The largest possible amplitude $K_{\rm{p,max}}$ occurs at the orbital inclination $i=90^{\circ}$."
 The instantaneous RV shift of the planetary signal with respect to the star depends on the orbital phase 6. High-precision RV measurements revealed the existence of a hot Jupiter orbiting the GO-type main-sequence star HD 75289A (Udry et al.," The instantaneous RV shift of the planetary signal with respect to the star depends on the orbital phase $\phi$, High-precision RV measurements revealed the existence of a hot Jupiter orbiting the G0-type main-sequence star HD 75289A (Udry et al."
 2000)., 2000).
 Table 1. summarises the parameters of the planet and its host star, Table \ref{tab:hd75289} summarises the parameters of the planet and its host star.
 We note in passing that the system also contains a faint low-mass stellar component. separated by =621AU from the primary (Mugrauer et al.," We note in passing that the system also contains a faint low-mass stellar component, separated by $\approx 621~\rm{AU}$ from the primary (Mugrauer et al."
 2004)., 2004).
 The planetary system of HD 75289A seemed to bear the opportunity to detect the reflected starlight for the following reasons: First. the brightness of HD 75289A in the visual (V=6.35 mag) enables a large amount of high-S/N spectra to be acquired within a short period of time. which helps to reduce photon noise. the dominant noise source.," The planetary system of HD 75289A seemed to bear the opportunity to detect the reflected starlight for the following reasons: First, the brightness of HD 75289A in the visual $V=6.35$ mag) enables a large amount of high-S/N spectra to be acquired within a short period of time, which helps to reduce photon noise, the dominant noise source."
 Second. the amount of starlight received from this planet is likely to be high due to the extremely small distance to its host star of only «a=0.048 AU (cf.," Second, the amount of starlight received from this planet is likely to be high due to the extremely small distance to its host star of only $a=0.048$ AU (cf."
 Equation 1))., Equation \ref{equ:2}) ).
 Third. a GO V star shows a rich absorption-line spectrum at optical wavelengths. a prerequisite for the data-synthesis approach outlined in Section 4.," Third, a G0 V star shows a rich absorption-line spectrum at optical wavelengths, a prerequisite for the data-synthesis approach outlined in Section 4."
" Using data of the stellar mass M,. the planetary minimum mass Mysini. and the RV semi-amplitude of the reflex motion of the star KY. the maximum possible RV semr-amplitude of the planet can be determined (from Equation 6)) to be Ky,=138.915.8km s."," Using data of the stellar mass $M_{\rm{s}}$, the planetary minimum mass $M_{\rm{p}}\sin i$, and the RV semi-amplitude of the reflex motion of the star $K_{\rm{s}}$, the maximum possible RV semi-amplitude of the planet can be determined (from Equation \ref{equ:doppler}) ) to be $K_{\rm p,max}=148.9\pm15.8~\rm{km~s^{-1}}$ ."
 An estimate of the planetary radius Ry can be obtained from a comparison of the radii determined for other hot Jupiters with the transit method., An estimate of the planetary radius $R_{\rm{p}}$ can be obtained from a comparison of the radii determined for other hot Jupiters with the transit method.
 To this end. we selected a subsample of known transiting hot Jupiters (Burrows et al.," To this end, we selected a subsample of known transiting hot Jupiters (Burrows et al."
 2007. and references therein). based on the following eriteria: (1) Planet mass between the minimum mass of HD 75289Ab and 2.3 times this value (for random orientation of the planetary orbit there is a 90 confidence that the true mass of HD 75289Ab lies in this interval). (," 2007, and references therein), based on the following criteria: (i) Planet mass between the minimum mass of HD 75289Ab and 2.3 times this value (for random orientation of the planetary orbit there is a 90 confidence that the true mass of HD 75289Ab lies in this interval). ("
1) Orbital radii similar to that of HD 75289Ab within 20%.. (,ii) Orbital radii similar to that of HD 75289Ab within 20. (
iit) Host star spectral types close to GO V (F8 V - GI V).,iii) Host star spectral types close to G0 V (F8 V - G1 V).
 The resulting group of objects contains OGLE 10b. OGLE ΠΙΟ. HAT-P-Ib. XO-Ib. HD 209458b.," The resulting group of objects contains OGLE 10b, OGLE 111b, HAT-P-1b, XO-1b, HD 209458b."
From this sample. we estimated the planetary radius of HD 75289Ab as the average of the radit of these five hot Jupiters finding,"From this sample, we estimated the planetary radius of HD 75289Ab as the average of the radii of these five hot Jupiters finding"
with /= 5 times the grid size. and where «ο>2.3.,"with $l = $ 5 times the grid size, and where $\Delta S / C_v > 2.3$."
 This most intense current laver is found to rise upward with the eruption of the flux rope., This most intense current layer is found to rise upward with the eruption of the flux rope.
 The associated post-reconnection lield lines are initially low hing and form a narrow sigmoid shaped bundle as can be seen in Figures I0((b) and 11((b)., The associated post-reconnection field lines are initially low lying and form a narrow sigmoid shaped bundle as can be seen in Figures \ref{fig10}( (b) and \ref{fig11}( (b).
 With time. the post-reconnection loops broaden ancl rise up. showing cusped apexes (Figures I0((e)(h) and Figures 11(GD(D)).," With time, the post-reconnection loops broaden and rise up, showing cusped apexes (Figures \ref{fig10}( (e)(h) and Figures \ref{fig11}( (d)(f))."
 The morphology. of the post-reconnection loops. which iransition from an initially narrow low-lving sigmoid bundle to a broad. sigmoid-shaped row of loops wil cusped apexes is in (qualitative agreement with the observed evolution of the post-[lare X-ray. brightening shown in Figures 9((a)(c)(e).," The morphology of the post-reconnection loops, which transition from an initially narrow low-lying sigmoid bundle to a broad, sigmoid-shaped row of loops with cusped apexes is in qualitative agreement with the observed evolution of the post-flare X-ray brightening shown in Figures \ref{fig9}( (a)(c)(e)."
 The foot points of the post-reconnection loops (panels (ο)(4) of Figures 10)) can be compared qualitatively with the evolution of the flare ribbons in the lower solar atmosphere as shown in the IBinode SOT observation (panels (b)(d)(E) of Figure 9))., The foot points of the post-reconnection loops (panels (c)(f)(i) of Figures \ref{fig10}) ) can be compared qualitatively with the evolution of the flare ribbons in the lower solar atmosphere as shown in the Hinode SOT observation (panels (b)(d)(f) of Figure \ref{fig9}) ).
 The ribbon corresponding to the positive polarity [oot points (orange ribbon in Figures 10((0)(D)()) of the loops is found to sweep ποπανανά across (he newly emerged positive polarity spot. similar to the apparent movement of (he positive polarity ribbon seen the observation (panels (b)(«l)(£) of Figure 9)) in relation to the observed positive emergimeg spot.," The ribbon corresponding to the positive polarity foot points (orange ribbon in Figures \ref{fig10}( (c)(f)(i)) of the post-reconnection loops is found to sweep southward across the newly emerged positive polarity spot, similar to the apparent movement of the positive polarity ribbon seen the observation (panels (b)(d)(f) of Figure \ref{fig9}) ) in relation to the observed positive emerging spot."
 For the ribbon corresponding to the negative polarity footvpoints (the vellow ribbon im Figures 10CCC)CD)Q)). ils eastern portion is found (to extend ancl sweep northward into the dominant pre-existing negative spot. while ils western hook-shaped portion is found (to sweep nortliwad across (he newly emerged negative spot.," For the ribbon corresponding to the negative polarity footvpoints (the yellow ribbon in Figures \ref{fig10}( (c)(f)(i)), its eastern portion is found to extend and sweep northward into the dominant pre-existing negative spot, while its western hook-shaped portion is found to sweep northward across the newly emerged negative spot."
 Similarly in (he SOT observation (panels (0)(dW) ol Fig 9)). for the negative polarity ribbon. the eastern portion sweeps northward into the dominant negative sunspot. while its western. upward curved hook-shaped portion is found to sweep northward across the minor. fragmented negative pores which have emerged (o the west of the main d-sunspot.," Similarly in the SOT observation (panels (b)(d)(f) of Fig \ref{fig9}) ), for the negative polarity ribbon, the eastern portion sweeps northward into the dominant negative sunspot, while its western, upward curved hook-shaped portion is found to sweep northward across the minor, fragmented negative pores which have emerged to the west of the main $\delta$ -sunspot."
 The modeled ribbons based on the footpoints certainly differ In many wavs in their shape ancl extent compared to the observed flare ribbons., The modeled ribbons based on the footpoints certainly differ in many ways in their shape and extent compared to the observed flare ribbons.
 But they capture some Κον qualitative features in the observed motions of the flare ribbons in relation to the photospheric magnetic [Iux concentrations., But they capture some key qualitative features in the observed motions of the flare ribbons in relation to the photospheric magnetic flux concentrations.
 We have presented an MIID model (hat qualitatively describes (he coronal magnetic field evolution of the eruptive flare in AR 10930 on December 13. 2006.," We have presented an MHD model that qualitatively describes the coronal magnetic field evolution of the eruptive flare in AR 10930 on December 13, 2006."
 The model assumes the eniergence of an east-west oriented magnetic [Iux rope into a pre-existing coronal magnetic field constructed. based. on the MDI full-disk magnetogram of the photospheric magnetic field at 20:51:01 UT on December 12., The model assumes the emergence of an east-west oriented magnetic flux rope into a pre-existing coronal magnetic field constructed based on the MDI full-disk magnetogram of the photospheric magnetic field at 20:51:01 UT on December 12.
 As described in Section 2.. a substantial smoothing of the observed photospheric magnetic flix density from the MDI magnetoeram is carried out such that the peak field strength on the lower boundary is reduced [rom 3000 G to," As described in Section \ref{sec:model}, a substantial smoothing of the observed photospheric magnetic flux density from the MDI magnetogram is carried out such that the peak field strength on the lower boundary is reduced from $\sim 3000$ G to"
If the Ix star is the sole coutributor to the near-infrared emission from © = 0-0 ool. then the orbital inclinaion in AQG2Z0-00 is 7 =387 — EY (i 38° — 50° for the extreme rauge of parameter values).,"If the K star is the sole contributor to the near-infrared emission from $\phi$ = 0 – 0.51, then the orbital inclination in A0620–00 is $i$ =$\arcdeg$ – $\arcdeg$ $i$ = $\arcdeg$ – $\arcdeg$ for the extreme range of parameter values)."
 This rauge is cousistent with the results of Shahbaz.Naylor&Charles(1991).. who found a best fit to their Ix light curve of 7 = 37° for g = 0.067 aud a confidence interval of 722 30° 15? for hat mass ratio (tjeir 22).," This range is consistent with the results of \citet{shahbaz1994}, who found a best fit to their K light curve of $i$ = $\arcdeg$ for $q$ = 0.067 and a confidence interval of $i \simeq$ $\arcdeg$ – $\arcdeg$ for that mass ratio (their 2)."
 The assumptio) tha the modulation of the light. curve at these orbital phases is surely ellipsoidal is not correct. i0wever.," The assumption that the modulation of the light curve at these orbital phases is purely ellipsoidal is not correct, however."
 The amplitude of the o = 0.25 peak relative to the light curve iniuima is smaller i the 1996 January light curve than in 1995 December or 1996 Decembe: light curves., The amplitude of the $\phi$ = 0.25 peak relative to the light curve minima is smaller in the 1996 January light curve than in 1995 December or 1996 December light curves.
 This leads toa lewer value for the orbital tuclination based on the 1996 Janulary data G= 387) and a rarve of j»ossible inclinations (38* — 117) that does not overlap those deteruinect rom the 1995/199¢) Deceuber light curves (7 = 1 — 50°)., This leads to a lower value for the orbital inclination based on the 1996 January data $i = 38\arcdeg$ ) and a range of possible inclinations $\arcdeg$ – $\arcdeg$ ) that does not overlap those determined from the 1995/1996 December light curves $i$ = $\arcdeg$ – $\arcdeg$ ).
 This change in the atiplitude of the nodulation even a ©0.25 indicates that there is some source of variable coitamniuation of he ellipsoidal modlation., This change in the amplitude of the modulation even at $\phi =0.25$ indicates that there is some source of variable contamination of the ellipsoidal modulation.
 Tt also gives a measure of the uncertainty in any deter:jinatiou of the inclination based oi just one observation epoch: the inclinations derived from our th‘ee liglit curves lave a rauge of 7°., It also gives a measure of the uncertainty in any determination of the inclination based on just one observation epoch: the inclinations derived from our three light curves have a range of $\arcdeg$ .
 In suiunary. models includiug only the effect of ellipsoidal variatiois of the lv star give a lower liiit to the inclination of 7>38 and show that there is a variabe distortion of he ellipsoidal variatious even at ©=0.25.," In summary, models including only the effect of ellipsoidal variations of the K star give a lower limit to the inclination of $i \geq 38\arcdeg$ and show that there is a variable distortion of the ellipsoidal variations even at $\phi =0.25$."
 We next modeled the light curves over the fll binary orbit. adii© all accretion disk and a bright spot to the IX star.," We next modeled the light curves over the full binary orbit, adding an accretion disk and a bright spot to the K star."
 We have complete orbital coverage ouly in he H bancl. so we cannot οΗΣΑΠ system parameters sich as the temperattre of the accretion disk aud bright spot wihb our models.," We have complete orbital coverage only in the H band, so we cannot constrain system parameters such as the temperature of the accretion disk and bright spot with our models."
 Rather. our goal was to find simple moces with reasonable pa‘alneter values that it the observed light curves iu order o determine the raiige of possible valles for the orbital inclination.," Rather, our goal was to find simple models with reasonable parameter values that fit the observed light curves in order to determine the range of possible values for the orbital inclination."
" Siuce the derived orbital itclinations do not ceyencl strongly ou thetlass rat]o. the tempe""allure ol the Ix star. nor the gravity «arkenine coellicier. we fixed their valles at yg = (067. Ts = |100 Ix aud ο = 0.08."," Since the derived orbital inclinations do not depend strongly on the mass ratio, the temperature of the K star, nor the gravity darkening coefficient, we fixed their values at $q$ = 0.067, $T_{2}$ = 4100 K and $\beta$ = 0.08."
" The inuer raclits of he accretion clisk was fixed at ϐ.001 Ry, atd the outer 'aclius at 0.5 Ry. which is the outer racLIs lu visible lieit found by lars1.Robinson&Wood(199L)."," The inner radius of the accretion disk was fixed at 0.001 $_{L_{1}}$ and the outer radius at 0.5 $_{L_{1}}$, which is the outer radius in visible light found by \citet{marsh1994}."
. To allow for a beamed bright spot on the disk rim. we set tle [laὁ hall-auge of the clisk to a 5all but non-zero value of 1°.," To allow for a beamed bright spot on the disk rim, we set the flare half-angle of the disk to a small but non-zero value of $^{\circ}$."
" Te bright spot also exteuds ono the top surface o. the disk rom 0.15 — 0.5 Ry, (Marsh.Robinson&Wood1991).", The bright spot also extends onto the top surface of the disk from 0.45 – 0.5 $_{L_{1}}$ \citep{marsh1994}.
. Since the b‘ight spot is not ecipsed and since the disk has only a small flare. the spot component on the t«ip of the clisk merely adds to e constant disk [flux aud is otherwise irrelevaut.," Since the bright spot is not eclipsed and since the disk has only a small flare, the spot component on the top of the disk merely adds to the constant disk flux and is otherwise irrelevant."
 The bright spo and accretion disk were each assured to emit as single-temperature blackbocdies. with the linear limib-darkeiing coelficients for e accretion disk obtained from Claret (1998)..," The bright spot and accretion disk were each assumed to emit as single-temperature blackbodies, with the linear limb-darkening coefficients for the accretion disk obtained from \citet{claret1998}. ."
 Previous observations of A0620—00 in quiescence iive found no evidence [or siguilicaut irradiatiou of tLe lx starby the disk. so we did 1ot include ος'adiation in our n1oclels.," Previous observations of A0620–00 in quiescence have found no evidence for significant irradiation of the K starby the disk, so we did not include irradiation in our models."
IPC field near the location of PSR J1015—5719 [#33] is likely an artifact resulting from the large effective area correction used.),IPC field near the location of PSR $-$ 5719 3] is likely an artifact resulting from the large effective area correction used.)
 However. the exposures may not have been deep enough to reveal such X-ray emission. usually expected at the level L.~10 (e.g.. Becker Trumper 1997).," However, the exposures may not have been deep enough to reveal such X-ray emission, usually expected at the level $L_x \sim
10^{-3} \dot E$ (e.g., Becker Trumper 1997)."
 Future observations of these Epulsar fields with andChandra could thus be instructive., Future observations of these pulsar fields with and could thus be instructive.
 In conclusion. we find that two recently discoverec1 Parkes pulsars (PSRs J1015—5719 and J1637—4642) coule1 plausibly generate at least. part of the >-ray flux observed from two unidentified EGRET sources (3EG 1014—5705 and 3EG J1639—4702. respectively). and either of PSRs 1412-0145 or JI413-6141 could be associated with 3EG J1410—6147 if they are located closer than their DM distances by a factor of ~4.," In conclusion, we find that two recently discovered Parkes pulsars (PSRs $-$ 5719 and $-$ 4642) could plausibly generate at least part of the $\gamma$ -ray flux observed from two unidentified EGRET sources (3EG $-$ 5705 and 3EG $-$ 4702, respectively), and either of PSRs $-$ 6145 or $-$ 6141 could be associated with 3EG $-$ 6147 if they are located closer than their DM distances by a factor of $\sim 4$."
 Cases #33-5 and 7 represent promising targets for the forthcoming andGLAST missions., Cases 3–5 and 7 represent promising targets for the forthcoming and missions.
 D.ET. was supported by CONICET and Fundaciénn Antorchas. and is on leave from IAR. Argentina.," D.F.T. was supported by CONICET and Fundaciónn Antorchas, and is on leave from IAR, Argentina."
 Y.M.B. acknowledges the support of the High Energy Astrophysics division at the CfA and the project., Y.M.B. acknowledges the support of the High Energy Astrophysics division at the CfA and the project.
 Partial support from the US DOE grant DE-FG02-91ER40671 (D.E.T.) and NASA grants NAS8-39073 (Y.M.B.) and NAG5-9095 (EC.) is acknowledged., Partial support from the US DOE grant DE-FG02-91ER40671 (D.F.T.) and NASA grants NAS8-39073 (Y.M.B.) and NAG5-9095 (F.C.) is acknowledged.
 We thank D. Nice and G. E. Romero for comments., We thank D. Nice and G. E. Romero for comments.
hereafter respectively R1. Rl. aud RG} has established hat the mean pulse profiles of radio pulsars often consist of compoucuts designatedtec as core aud coual having differing characteristics.,"hereafter respectively R1, R4, and R6) has established that the mean pulse profiles of radio pulsars often consist of components designated as core and conal having differing characteristics."
 On the basis of a large bov of daa concerning the pulse widths. it has |νο sueeestedOO hat the core enuiüssion cluanates frou the full polar cap at ana titude ràaq4w uch is the sauc ds hes tellar radius. R (usually taken as 10 Ian) (RL).," On the basis of a large body of data concerning the pulse widths, it has been suggested that the core emission emanates from the full polar cap at an altitude $r_{core}$ which is the same as the stellar radius, $R_*$ (usually taken as 10 km) (R4)."
 The coual cussion in] iferred O COue from a hollow cone at an altitule. Ücones of about 10-20 R (htiy," The conal emission is inferred to come from a hollow cone at an altitude, $r_{cone}$, of about 10-20 $R_*$ (R6)."
 This was the first lustance when differeut «Bnissdon altitudes were attritcad to differeit pulse conrponeuts;, This was the first instance when different emission altitudes were attributed to different pulse components.
 About the values of altitudes there is no 1inuinvo at preseut (kijak aud Cil 1998). but 1tds clear that core aud coual euussion altitudes could eenerallv be different.," About the values of altitudes there is no unanimity at present (Kijak and Gil 1998), but it is clear that core and conal emission altitudes could generally be different."
 Neeoues in view the previously quoted enmission altitudes of upto 2008... we take for hem a Oogeneral maxi fieureC» of the oler of 20% of the light evliuder radius rp=(οP/2a). although the altitudes could vary with the pulsar period P. the inclination angle between the magnetic aud rotation axes o. and the fraction of the emission cap that is active.," Keeping in view the previously quoted emission altitudes of upto $200 R_*$, we take for them a general maximum figure of the order of $20 \%$ of the light cylinder radius $r_L = (c\,P/2\,\pi)$, although the altitudes could vary with the pulsar period $P$, the inclination angle between the magnetic and rotation axes $\alpha$, and the fraction of the emission cap that is active."
 However. later we also conmuent on the possibility of having arbitrarily huge cussion altitudes.," However, later we also comment on the possibility of having arbitrarily large emission altitudes."
 It is obvious that components of pulsar raciation which emauate from differeut altitudes will show some differences in properties related to their emission altitudes., It is obvious that components of pulsar radiation which emanate from different altitudes will show some differences in properties related to their emission altitudes.
 Oue such is the often seen nou-coincidence of ceutres of the core components and imnid-poiuts of the conal pairs., One such is the often seen non-coincidence of centres of the core components and mid-points of the conal pairs.
 This phase offset between core aud conal conrponeuts through kincmatic effects provides a new method to learn about these chussion altitudes., This phase offset between core and conal components through kinematic effects provides a new method to learn about these emission altitudes.
 As we demoustrate later. it is uuencunmbered by the ignorance of detailslc of the emission niechanisuir aud allows interpretation of observations in1 more detail so as to coustrain the theory.," As we demonstrate later, it is unencumbered by the ignorance of details of the emission mechanism and allows interpretation of observations in more detail so as to constrain the theory."
 What follows is a framework for a systematic understanding of these phase offsets between core and conal components as being due to their differing altitudes., What follows is a framework for a systematic understanding of these phase offsets between core and conal components as being due to their differing altitudes.
 This approach has many advantages because one can learn about eniüssion altitudes by focussing ondifferent pulse components observed atone frequency rather than studyiug componeit atdifferent frequeucies., This approach has many advantages because one can learn about emission altitudes by focussing on pulse components observed at frequency rather than studying component at frequencies.
 Firstly this complements multifrequeucy considerations (Philips. 1992: Iijas and Cal. 1998).," Firstly this complements multifrequency considerations (Philips, 1992; Kijak and Gil, 1998)."
 More iniportautlyv it allows us to skirt the complicaIOUS arise Yonà our ignorantὉ of radis to frequency mapping aud related time delays etc. (, More importantly it allows us to skirt the complications arising from our ignorance of radius to frequency mapping and related time delays etc. (
Cordes 1992).,Cordes 1992).
" Firthermore. we expect that our stiv will clarify whether the pase offsets are as expected YOU ouly the coutribiuticMis We are considering or call for additional coutribuloli which provide us wih hints about the details o the emdssio1 process,"," Furthermore, we expect that our study will clarify whether the phase offsets are as expected from only the contributions we are considering or call for additional contributions which provide us with hints about the details of the emission process."
 Lastly let us nieution that core cussion comme frou the stellar surface las implica10119 for the pulsar matters equation of state leading to strong indications that pulsus are stranee stars rather than neutn stars (Ixapoor and Shukre 1999. 2001).," Lastly let us mention that core emission coming from the stellar surface has implications for the pulsar matter's equation of state leading to strong indications that pulsars are strange stars rather than neutron stars (Kapoor and Shukre 1999, 2001)."
 This mases it very maportaut to verity if core comporents truly emianuate on the stellar srface., This makes it very important to verify if core components truly emanate on the stellar surface.
 This interpretation of the remarkable pulse width relation of Rankin rests crucially on the assumption that core enission comes from thefull polar cap (RL)., This interpretation of the remarkable pulse width relation of Rankin rests crucially on the assumption that core emission comes from the polar cap (R4).
 We shall see what inipications this has for the core/cone offsets., We shall see what implications this has for the core/cone offsets.
 I1 the relation between the cussion altitudes of core/cone components aud the phase offsets due to these differius altitudes. one inevitable kincmatical contribution is due to differeutial aberration (Ikapoor aud Shure 1998 hereafter," In the relation between the emission altitudes of core/cone components and the phase offsets due to these differing altitudes, one inevitable kinematical contribution is due to differential aberration (Kapoor and Shukre 1998 hereafter"
It can be seen that the cluster masses are typically overestimated.,It can be seen that the cluster masses are typically overestimated.
 The strongest bias to high masses at time Τι (including the peak) shows that the unbound stars stay for a long time close to the cluster and contribute to the outer density profile in the fitting procedure., The strongest bias to high masses at time $T_1$ (including the peak) shows that the unbound stars stay for a long time close to the cluster and contribute to the outer density profile in the fitting procedure.
 This effect depends on the initial conditions., This effect depends on the initial conditions.
" At later times T3, T4 the cluster mass distribution becomes stable and the bias is independent of the age of the cluster."," At later times $T_3$, $T_4$ the cluster mass distribution becomes stable and the bias is independent of the age of the cluster."
" Figure 8 shows a histogram of the fraction on the sky per bin in Mi/Ma=(ri/r;)? for the parameter surface in Figure 7 corresponding to time 74, where the ""tidal"" masses M; and M, are calculated with Equation 15 from r; and rz, respectively."," Figure \ref{fig:histrt} shows a histogram of the fraction on the sky per bin in $M_t/M_{cl}=(r_t/r_J)^3$ for the parameter surface in Figure \ref{fig:aisurf1} corresponding to time $T_4$, where the “tidal” masses $M_t$ and $M_{\rm cl}$ are calculated with Equation \ref{eq:rjac} from $r_t$ and $r_J$, respectively."
 The solid line shows the distribution for projections in the range between —20?«b«+20°., The solid line shows the distribution for projections in the range between $-20^\circ < b < +20^\circ$ .
 The most frequent overestimation of the mass is M;/M_1=1.18., The most frequent overestimation of the mass is $M_t/M_{cl}=1.18$.
 The dashed line shows the distribution for projections in the range between —40°«b<440°., The dashed line shows the distribution for projections in the range between $-40^\circ < b < +40^\circ$.
 'The most frequent ratio is the same as for the solid line., The most frequent ratio is the same as for the solid line.
" However, the ratios extend up to M;/M,=1.7."," However, the ratios extend up to $M_t/M_{cl}=1.7$."
 The dotted line shows the distribution for all projections., The dotted line shows the distribution for all projections.
"The most frequent ratio is the same as for the other two lines, but the ratios extend up to M;/M,e2.1.","The most frequent ratio is the same as for the other two lines, but the ratios extend up to $M_t/M_{cl}\approx 2.1$."
" Figure 9 shows the mean mass (M;/M.1) as a function of the galactic latitude b, where M; is the mass of cluster stars within radius r; which has been obtained by the fitting procedure."," Figure \ref{fig:meanmass} shows the mean mass $\langle M_t/M_{\rm cl} \rangle$ as a function of the galactic latitude $b$, where $M_t$ is the mass of cluster stars within radius $r_t$ which has been obtained by the fitting procedure."
 Shown are the curves corresponding to times 7;—T4., Shown are the curves corresponding to times $T_1 - T_4$.
 All curves show a local minimum which is located roughly at b=O°., All curves show a local minimum which is located roughly at $b=0^\circ$.
 For time ΤΙ the mean overestimation of the mass reaches 2.5 in the direction of the Galactic poles., For time $T_1$ the mean overestimation of the mass reaches 2.5 in the direction of the Galactic poles.
 Usually only a small fraction of cluster stars are identified as members leading to an increased statistical uncertainty in the fitting procedure., Usually only a small fraction of cluster stars are identified as members leading to an increased statistical uncertainty in the fitting procedure.
" The uppermost plot in Figure 10 shows the parameter surface derived for the 400 most massive (i.e. the brightest) stars in the simulation at time T, (data courtesy of J. Beuria).", The uppermost plot in Figure \ref{fig:aisurf2} shows the parameter surface derived for the 400 most massive (i.e. the brightest) stars in the simulation at time $T_4$ (data courtesy of J. Beuria).
 The reason for the low values of r;/r; is mass segregation., The reason for the low values of $r_t/r_J$ is mass segregation.
 The stronger asymmetries in +b compared to the plots in Figure 7 are due to a slightly asymmetric distribution of the 400 mass-segregated stars in position space., The stronger asymmetries in $\pm b$ compared to the plots in Figure \ref{fig:aisurf1} are due to a slightly asymmetric distribution of the 400 mass-segregated stars in position space.
 In order to measure the statistical scatter in Τι we applied a bootstrap analysis., In order to measure the statistical scatter in $r_t$ we applied a bootstrap analysis.
 We divided the N-body snapshot file for time T4 with N=40404 particles into 100 small-number samples of Nsample=400 particles each (the remaining particles were dropped out of the analysis)., We divided the $N$ -body snapshot file for time $T_4$ with $N=40404$ particles into 100 small-number samples of $N_{\rm sample}=400$ particles each (the remaining particles were dropped out of the analysis).
 For each of these small-number samples we, For each of these small-number samples we
We note that using reduction methods exactly like those detailed for H 1743-322. we have detected Fe XXV and Fe XXVI lines in a 60 ksee Chandra/HETGS spectrum of Cygnus X-I obtained on 3 March 2004. when the source was in à canonical low/hard state (Miller et 22004c. in prep.).,"We note that using reduction methods exactly like those detailed for H $-$ 322, we have detected Fe XXV and Fe XXVI lines in a 60 ksec /HETGS spectrum of Cygnus X-1 obtained on 3 March 2004, when the source was in a canonical low/hard state (Miller et 2004c, in prep.)."
 Cygnus X-1 is known to produce a compact. steady radio jet in this state (Stirling et 22001).," Cygnus X-1 is known to produce a compact, steady radio jet in this state (Stirling et 2001)."
 Highly tonized Fe absorption lines have also been detected in the neutron star system GX 1341 (Sidoli et 22002)., Highly ionized Fe absorption lines have also been detected in the neutron star system GX $+$ 1 (Sidoli et 2002).
 Interestingly. GX 1341 1s also a radio emitter in which a delay between its X-ray spectral hardness and radio brightness has been observed which 1s very similar to that seen in GRS 19154105 (Homan et 220043). strongly suggesting à compact jet is also at work in this system.," Interestingly, GX $+$ 1 is also a radio emitter in which a delay between its X-ray spectral hardness and radio brightness has been observed which is very similar to that seen in GRS $+$ 105 (Homan et 2004a), strongly suggesting a compact jet is also at work in this system."
 However. simultaneous radio observations of H 1743-322 with the VLA/VLBA reveal that the radio source was quite strong during the first and second observations nearly 10 mJy at GHz frequencies) — but only 2 mJy during the third observation. and less than | mJy during the fourth observation (Rupen et 22004. in prep.).," However, simultaneous radio observations of H $-$ 322 with the VLA/VLBA reveal that the radio source was quite strong during the first and second observations --- nearly 10 mJy at GHz frequencies) — but only 2 mJy during the third observation, and less than 1 mJy during the fourth observation (Rupen et 2004, in prep.)."
 Thus. strong radio activity (indicative of a jet) is seen during the second observation (where absorption lines are not observed). and absorption lines are clearly observed in the third and fourthChandra observations when the radio flux was low.," Thus, strong radio activity (indicative of a jet) is seen during the second observation (where absorption lines are not observed), and absorption lines are clearly observed in the third and fourth observations when the radio flux was low."
 At minimum. this suggests that if jets and tonized absorption are causally related. they are not related in a simple way.," At minimum, this suggests that if jets and ionized absorption are causally related, they are not related in a simple way."
 Both may merely be the effect of a strongly active corona. for instance.," Both may merely be the effect of a strongly active corona, for instance."
 Aoreover. the detection of highly tonized Fe absorption lines in high inclination. neutron star bimaries may signify that it is unlikely that ionized absorption lines are tied to jet production through a mechanism like ΑΕΙ.," Moreover, the detection of highly ionized Fe absorption lines in high inclination neutron star binaries may signify that it is unlikely that ionized absorption lines are tied to jet production through a mechanism like AEI."
 Similar lines have also been seen in the persistent emission of the (edge-on) dipping neutron star binaries X 1254-690 (Boirin Parmar 2003). X 1624-490 (Parmar et 22002). MXB 1658-29 (Sidoli et 22001) and 4U 1916-053 (Boirin et 22004).," Similar lines have also been seen in the persistent emission of the (edge-on) dipping neutron star binaries X $-$ 690 (Boirin Parmar 2003), X $-$ 490 (Parmar et 2002), MXB $-$ 29 (Sidoli et 2001) and 4U $-$ 053 (Boirin et 2004)."
 In contrast to the absorption lines in the black hole systems H 1743-322. GX 339-4. and XTE J1650—500. the absorption lines in these edge-on neutron star binaries are not significantly blue-shifted.," In contrast to the absorption lines in the black hole systems H $-$ 322, GX $-$ 4, and XTE $-$ 500, the absorption lines in these edge-on neutron star binaries are not significantly blue-shifted."
 Although the appearance of lines in both black hole and neutron star systems suggests a related physical origin. the lack of blue-shifts in the lines seen in neutron stars may hint at a different physical picture.," Although the appearance of lines in both black hole and neutron star systems suggests a related physical origin, the lack of blue-shifts in the lines seen in neutron stars may hint at a different physical picture."
 Without significant blue-shifts. it is difficult to clearly tie the lines to an outflow.," Without significant blue-shifts, it is difficult to clearly tie the lines to an outflow."
 The lines may not be due to absorption in an outflow. but might be due to absorption in ambient gas within the system.," The lines may not be due to absorption in an outflow, but might be due to absorption in ambient gas within the system."
" The absorption lines in MXB 1658—29., for instance. are not observed to vary with orbital phase or even during dipping events (Sidoli et 22001). which suggests that the absorbing gas is not as close to the compact object as the absorbing gas in black hole We wish to thank CXC Director Harvey Tananbaum and the CXC staff. and Jean Swank and theRXTE staff. for making these TOO observations."," The absorption lines in MXB $-$ 29, for instance, are not observed to vary with orbital phase or even during dipping events (Sidoli et 2001), which suggests that the absorbing gas is not as close to the compact object as the absorbing gas in black hole We wish to thank CXC Director Harvey Tananbaum and the CXC staff, and Jean Swank and the staff, for making these TOO observations."
 We thank John Houck for assistance with ISIS., We thank John Houck for assistance with ISIS.
 JMM acknowledges helpful discussions with Mike Nowak. David Huenemoerder. Aneta Siemiginowska. Rob Fender. Herman Marshall. and Patrick Wojdowski.," JMM acknowledges helpful discussions with Mike Nowak, David Huenemoerder, Aneta Siemiginowska, Rob Fender, Herman Marshall, and Patrick Wojdowski."
 JMM gratefully acknowledges support from the NSF through its Astronomy and Astrophysics Postdoctoral Fellowship program., JMM gratefully acknowledges support from the NSF through its Astronomy and Astrophysics Postdoctoral Fellowship program.
 WHGL and JH gratefully acknowledge support from NASA., WHGL and JH gratefully acknowledge support from NASA.
 DS acknowledges support from the SAO Clay Fellowship., DS acknowledges support from the SAO Clay Fellowship.
 This research has made use of the data and resources obtained through the HEASARC on-line service. provided by NASA-," This research has made use of the data and resources obtained through the HEASARC on-line service, provided by NASA-GSFC."
the Two Micron All Sky Survey. probing from ~2.9 kpe.,"the Two Micron All Sky Survey, probing from $\sim2-9$ kpc."
 They detected no significant correlations at latitudes |b]—50°. but did detect correlations on smallscales (1°. ~LOO pe) at lowerlatitudes. which they attributed to structure in the thick disc.," They detected no significant correlations at latitudes $|b|\ga50\degr$, but did detect correlations on smallscales $1\degr$, $\sim100$ pc) at lower latitudes, which they attributed to structure in the thick disc."
 Lemon (2004) performed a similar analysis for nearby F-type stars in the Tillennium Galaxy Catalogue and found no significant clustering., Lemon (2004) performed a similar analysis for nearby F-type stars in the Millennium Galaxy Catalogue and found no significant clustering.
 Starkenburg (2009) used a correlation function in dimensions to identify substructures in the Spaghetti beam survey of the distant halo (Morrison 2000: Dohm-Palmer 2000)., Starkenburg (2009) used a correlation function in dimensions to identify substructures in the Spaghetti pencil-beam survey of the distant halo (Morrison 2000; Dohm-Palmer 2000).
 With this method they obtained a significant detection of clustering and set a lower limit on the number of qalo stars in all substructures., With this method they obtained a significant detection of clustering and set a lower limit on the number of halo stars in all substructures.
 Similarly. Schlaufman (2009) constrained the mass fraction of the halo in detectable substructure by estimating the completeness of their overdensity detection algorithm.," Similarly, Schlaufman (2009) constrained the mass fraction of the halo in detectable substructure by estimating the completeness of their overdensity detection algorithm."
 Starkenburg and Schlaufman concluded that :10% (by number of stars) and ~304 (by volume) of the Milky Way halo belongs to groups meeting their respective definitions of phase space substructure., Starkenburg and Schlaufman concluded that $>10\%$ (by number of stars) and $\sim30\%$ (by volume) of the Milky Way halo belongs to groups meeting their respective definitions of phase space substructure.
 These methods were tested on “mock catalogues! of tracer stars based on simplified models of the stellar halo., These methods were tested on `mock catalogues' of tracer stars based on simplified models of the stellar halo.
 The work of Xue (2009. 2011) is particularly relevant.," The work of Xue (2009, 2011) is particularly relevant."
 Xue (2009) considered the pairwise radial velocity separation of a sample of 2558 halo BHB stars as a function of their separation in space. but found no evidence of clustering.," Xue (2009) considered the pairwise radial velocity separation of a sample of 2558 halo BHB stars as a function of their separation in space, but found no evidence of clustering."
 From comparisons to the simulations of Bullock Johnston (2005). Xue concluded that a pairwise velocity statistic was not capable of detecting structure against a more smoothly distributed background in phase space (made up from stars in phase-mixed streams).," From comparisons to the simulations of Bullock Johnston (2005), Xue concluded that a pairwise velocity statistic was not capable of detecting structure against a more smoothly distributed background in phase space (made up from stars in phase-mixed streams)."
 However. their observed signal was not compared to an expected signal from random realizations.," However, their observed signal was not compared to an expected signal from random realizations."
 More recently. Xue (2011) studied an enlarged catalogue of BHBs comprising more than 4000 stars.," More recently, Xue (2011) studied an enlarged catalogue of BHBs comprising more than 4000 stars."
 They quantified clustering in this larger sample using a four-dimensional metric (similar to that of Starkenburg 2009). finding a significant excess of clustering on small scales by comparison to smooth models.," They quantified clustering in this larger sample using a four-dimensional metric (similar to that of Starkenburg 2009), finding a significant excess of clustering on small scales by comparison to smooth models."
 The conclusions of this more recent study by Xue agree with our own. as we discuss further in7.," The conclusions of this more recent study by Xue agree with our own, as we discuss further in."
. The statistic. we develop in this paper also builds on the approach of Starkenburg (2009)., The statistic we develop in this paper also builds on the approach of Starkenburg (2009).
 We detine a two-point correlation function based on a metric combining pairwise separations in four readily obtained phase space observables for halo stars cangular position. radial distance and radial velocity).," We define a two-point correlation function based on a metric combining pairwise separations in four readily obtained phase space observables for halo stars (angular position, radial distance and radial velocity)."
 We apply this statistic to the catalogue of BHB stars published by Xue and demonstrate that a significant clustering signal can be recovered from these data., We apply this statistic to the catalogue of BHB stars published by Xue and demonstrate that a significant clustering signal can be recovered from these data.
 A clustering metric of the kind we propose can be tuned to probe a specitic scale in phase space by adjusting the weight given to each of its components., A clustering metric of the kind we propose can be tuned to probe a specific scale in phase space by adjusting the weight given to each of its components.
 In the stellar halo. however. many ‘components’ may be superimposed with a complex assortment of scales and morphologies in phase space (Johnston 2008: Cooper 2010).," In the stellar halo, however, many `components' may be superimposed with a complex assortment of scales and morphologies in phase space (Johnston 2008; Cooper 2010)."
 For this reason. it is not clear. a priori. what sort of signal to expect. or which scales are most relevant.," For this reason, it is not clear, a priori, what sort of signal to expect, or which scales are most relevant."
" We find that we cannot identify an ""optimal metric.", We find that we cannot identify an `optimal' metric.
 Instead. we make a fiducial choice based on the self-consistent accreted halo models of Cooper(2010).," Instead, we make a fiducial choice based on the self-consistent accreted halo models of Cooper(2010)."
 These incorporate an ab initio galaxy formation model using high-resolution cosmological N-body simulations from the Aquarius project (Springel 2008)., These incorporate an ab initio galaxy formation model using high-resolution cosmological N-body simulations from the Aquarius project (Springel 2008).
 We apply our fiducial metric to the data and to these models., We apply our fiducial metric to the data and to these models.
 We find that even though both the metric and our choice of sealing are simple. this approach has the power to diseriminate quantitatively between qualitatively different stellar haloes.," We find that even though both the metric and our choice of scaling are simple, this approach has the power to discriminate quantitatively between qualitatively different stellar haloes."
 We deseribe the basis of our method in and the SDSS DR6 BHB catalogue of Xue (2008) in3., We describe the basis of our method in and the SDSS DR6 BHB catalogue of Xue (2008) in.
. In we describe our simulations 4.13) and our procedure for constructing mock catalogues 4.23)., In we describe our simulations ) and our procedure for constructing mock catalogues ).
 describes our choice of a fiducial metric., describes our choice of a fiducial metric.
 In we apply our method to quantify clustering in the SDSS data and compare this with our mock catalogues., In we apply our method to quantify clustering in the SDSS data and compare this with our mock catalogues.
 Our conclusions are given in7., Our conclusions are given in.
". The most readily obtained phase-space observables for halo stars are their Galactic angular coordinates. / and b. heliocentric radial distance. myer. and heliocentric line-of-sight velocity. ej,4."," The most readily obtained phase-space observables for halo stars are their Galactic angular coordinates, $l$ and $b$, heliocentric radial distance, $r_{\mathrm{hel}}$, and heliocentric line-of-sight velocity, $v_{\mathrm{hel}}$."
 Using its angular position and distance estimate. each star can be assigned a three-dimensional position vector in galactocentric Cartesian coordinates. 7(Y.Y.Z). and a radial velocity corrected for the Solar and Local Standard of Rest motions. v.," Using its angular position and distance estimate, each star can be assigned a three-dimensional position vector in galactocentric Cartesian coordinates, $\bmath{r}\,(X,Y,Z)$, and a radial velocity corrected for the Solar and Local Standard of Rest motions, $v$."
 We begin by defining a scaling relation (metric). A. which combines these observables into a simple ‘phase-space separation’ between two stars: Here. |r;rj| is the separation of a pair of stars in coordinate space (in Kiloparsecsy. and ο)0; is the difference in their radial velocities tin kilometres per second).," We begin by defining a scaling relation (metric), $\Delta$, which combines these observables into a simple `phase-space separation' between two stars: Here, $|\bmath{r_i}-\bmath{r_j}|$ is the separation of a pair of stars in coordinate space (in kiloparsecs), and $v_{i}-v_{j}$ is the difference in their radial velocities (in kilometres per second)."
" The scaling factor a). has units of kpckim1s, such that A has units of kpc."," The scaling factor $w_{v}$ has units of $\mathrm{kpc\,km^{{-}1}\,s}$, such that $\Delta$ has units of kpc."
" The choice of the is arbitrary unless a particular ""phase space scale’ of interest can be identified.", The choice of $w_{v}$ is arbitrary unless a particular `phase space scale' of interest can be identified.
 This is not straightforward: we discuss some possible choices below., This is not straightforward; we discuss some possible choices below.
 The aim of this paper is to explore £(4X). the cumulative two- correlation function of halo stars in the metric defined by Equation |..," The aim of this paper is to explore $\xi({\Delta})$, the cumulative two-point correlation function of halo stars in the metric defined by Equation \ref{eqn:delta_metric}."
 Throughout. we use the estimator Here D2D(«&A) counts the number of pairs in the sample separated by less than A. and GRA(<A) is the equivalent count for pairs of random points within the survey volume. averaged over several realizations.," Throughout, we use the estimator Here $DD(<\Delta)$ counts the number of pairs in the sample separated by less than $\Delta$ , and $\left \langle RR(<\Delta)
\right \rangle$ is the equivalent count for pairs of random points within the survey volume, averaged over several realizations."
 To generate these realizations we ‘shuffle’ the data by randomly reassigning Val:Chel) pairs to different (/.b) coordinates drawn from the oriinalcatalogue.," To generate these realizations we `shuffle' the data by randomly reassigning $(r_{\mathrm{hel}},v_{\mathrm{hel}})$ pairs to different $(l,b)$ coordinates drawn from the original."
.. Similar methods for quantifying the clustering of stars in a four-dimensional space are described by Starkenburg (2009. applied to a sample of giant stars from the Spaghetti survey) and Xue (2011. appliedto a sample of BHB stars from SDSS).," Similar methods for quantifying the clustering of stars in a four-dimensional space are described by Starkenburg (2009, applied to a sample of giant stars from the Spaghetti survey) and Xue (2011, appliedto a sample of BHB stars from SDSS)."
" Our A metric is very similar to the 9,4 metric ofStarkenburg et al.", Our $\Delta$ metric is very similar to the $\delta_{4\mathrm{d}}$ metric ofStarkenburg et al.
 in the limit of small angular separations?., in the limit of small angular .
. We have veritied, We have verified
"An interesting feature of this (vpe of breaking of 53 is that the neutrino mixing matrix is completely determined from the retrocireulant perturbation matrix ancl is. (hus. independent of the parameters appearing in the zeroth order neutrino mass matrix Al,.","An interesting feature of this type of breaking of $S_3$ is that the neutrino mixing matrix is completely determined from the retrocirculant perturbation matrix and is, thus, independent of the parameters appearing in the zeroth order neutrino mass matrix $M_\nu$."
 As a consequence. the deviations from TDM are entirely determined in terms of single unknown perturbation parameter À.," As a consequence, the deviations from TBM are entirely determined in terms of single unknown perturbation parameter $\lambda$."
 The neutrino mixing angles upto the third order perturbation are given by 5 solar and reactor neutrino mixing angles are related by . As pointed out earlier in (his section. the second and third order perturbation matrices can be interchanged.," The neutrino mixing angles upto the third order perturbation are given by and The solar and reactor neutrino mixing angles are related by As pointed out earlier in this section, the second and third order perturbation matrices can be interchanged."
" As a consequence of this interchange. the predictions for all the parameters remain intact except for the 2-3 mixing anele Go, which is given by lt can be seen that the 2-3 mixing angle im this case is shifted to below maximal by ihe same amount as it was above maximal in the earlier case."," As a consequence of this interchange, the predictions for all the parameters remain intact except for the 2-3 mixing angle $\theta _{23}$ which is given by It can be seen that the 2-3 mixing angle in this case is shifted to below maximal by the same amount as it was above maximal in the earlier case."
The mass distribution of the extrasolar planets was recognized to be a kev feature of the growing new population since the first [ew detections (e.g.. Dasri Marcy 1997: Mayor. Queloz Udry 1998; Mazeh. Goldberg Latham 1998: HLeacox 1999: Mazel 1999: Stepinski Black 2000).,"The mass distribution of the extrasolar planets was recognized to be a key feature of the growing new population since the first few detections (e.g., Basri Marcy 1997; Mayor, Queloz Udry 1998; Mazeh, Goldberg Latham 1998; Heacox 1999; Mazeh 1999; Stepinski Black 2000)."
 Recent studies showed that (he mass distribution is probably flat or slightly decreasing in log AI (Jorissen. Mayor Udry. 2001: Zucker Mazeh. 2001: Tabachnik Tremaine 2002: Lineweaver Grether 2002). and has a distinet eutoff al about LO Jupiter," Recent studies showed that the mass distribution is probably flat or slightly decreasing in log M (Jorissen, Mayor Udry 2001; Zucker Mazeh 2001; Tabachnik Tremaine 2002; Lineweaver Grether 2002), and has a distinct cutoff at about 10 Jupiter"
temperature and degeneraey. parameter c: at the time of the mereer were significant.,temperature and degeneracy parameter $\psi$ at the time of the merger were significant.
 TFhis is not easy to study faithfully as our calculations often fail to converge when the point of helium ignition has become even mildly degenerate (or 2)., This is not easy to study faithfully as our calculations often fail to converge when the point of helium ignition has become even mildly degenerate $\psi \gtrsim2$ ).
 Llowever. we adopted an artifical way to test. this.," However, we adopted an artifical way to test this."
 For a merger product which was not degenerate enough to cause numerical problems. we switched. olf the nuclear evolution whilst allowing the helium-burning energv generation to continue.," For a merger product which was not degenerate enough to cause numerical problems, we switched off the nuclear evolution whilst allowing the helium-burning energy generation to continue."
 We then forciblv cooled the core as much as possible., We then forcibly cooled the core as much as possible.
 For lower helium. \WD masses (e.g. 0.2 ancl 0.25 AL). we were able to reach core temperatures as low as 10 Ix. and in all cases below 10' Ix. The initial core temperature appeared to have no discernible οσοι on the long-term evolution of the merger product once the artificial cooling was removed: the core temperature increased. again due to heat flowing inwards from the burning shell.," For lower helium WD masses (e.g. 0.2 and 0.25 $\rm M_{\odot}$ ), we were able to reach core temperatures as low as $10^{6}$ K, and in all cases below $10^{7}$ K. The initial core temperature appeared to have no discernible effect on the long-term evolution of the merger product once the artificial cooling was removed; the core temperature increased again due to heat flowing inwards from the burning shell."
 We also note that compressional heating of the accreting star during the merger should reduce the degeneracy olf the merger products compared to our hydrostatic calculations. potentially further reducing any small clleet of the core temperature at the time of the merger.," We also note that compressional heating of the accreting star during the merger should reduce the degeneracy of the merger products compared to our hydrostatic calculations, potentially further reducing any small effect of the core temperature at the time of the merger."
 However. we cannot euarantec that this effect can be totally ignored: similar calculations have been performed. by Lben (1990). who argued that the prior thermal history of the core was not negligible in the evolution of the post-mereer star.," However, we cannot guarantee that this effect can be totally ignored; similar calculations have been performed by Iben (1990), who argued that the prior thermal history of the core was not negligible in the evolution of the post-merger star."
 The results of our calculations are shown in Figs., The results of our calculations are shown in Figs.
 2 and 3.., \ref{fig:logg_tracks} and \ref{fig:HR_and_times}.
 The majority of the helium-burning phase of the merger products is spent close to the majority of known Le-sdO stars in the cllective temperature — surface eravity diagram., The majority of the helium-burning phase of the merger products is spent close to the majority of known He-sdO stars in the effective temperature – surface gravity diagram.
 This is particularly true for mergers of the AL. post-sdB stars., This is particularly true for mergers of the $\rm~M_{\odot}$ post-sdB stars.
 Furthermore. the trend. from the OA to M. models suggests that the lowest mass sdB stars in the population (0.3M.: see Fig. 1))," Furthermore, the trend from the 0.4 to $\rm~M_{\odot}$ models suggests that the lowest mass sdB stars in the population $\approx 0.3\rm~M_{\odot}$; see Fig. \ref{fig:initdist}) )"
 might help to fill in the Ilow-temperature end of the distribution., might help to fill in the low-temperature end of the distribution.
 Pronounced, Pronounced
 The reaction of vibrationally excited Πο aud atomic oxveen las little effect on the abuudauce of ΟΠ for most of the conditions explored.,by The reaction of vibrationally excited $_2$ and atomic oxygen has little effect on the abundance of OH for most of the conditions explored.
" It does coutribute to OTL formation just for fractional abundances of Tote= 1) in excess of > and for kinetic temperatures around 300 I&. For πιο lower temperatures the activation barriers of the reactions Ποιο> 0) | O prevent to reach a high enough OF formation rate whereas for mach larecr temperatures the rate coustaut chuhancements AX are low and ΟΠ is imostly formed by the reaction of Ποιο= 0) | ©. As concerns the reactions of vibrationally excited IL, with Πο. ΟΠ. aud CN. they have alinost no muüpact on the chemistry as the rate constant enhancenients are just moderate and require of a verv large fraction of vibrationally excited IT». to compete with the correspouding reactions of Ποιο= 0)."," It does contribute to OH formation just for fractional abundances of $_2$ $v=1$ ) in excess of $^{-4}$ and for kinetic temperatures around 300 K. For much lower temperatures the activation barriers of the reactions $_2$ $v>0$ ) + O prevent to reach a high enough OH formation rate whereas for much larger temperatures the rate constant enhancements $\Delta k$ are low and OH is mostly formed by the reaction of $_2$ $v=0$ ) + O. As concerns the reactions of vibrationally excited $_2$ with $^+$, OH, and CN, they have almost no impact on the chemistry as the rate constant enhancements are just moderate and require of a very large fraction of vibrationally excited $_2$ to compete with the corresponding reactions of $_2$ $v=0$ )."
 lu sununuaryv. from the five reactions of vibrationally excited IIs cousidercd here. ouly that with C! las a laree Hupact on the abuudances of some species; especially CID). in those regions Which are not too won (with kineticη PMtemperatures below a few hundreds of degrees Ixeblvi and which have a moderately high fraction of vibrationally excited Πο (in excess of 10 ©).," In summary, from the five reactions of vibrationally excited $_2$ considered here, only that with $^+$ has a large impact on the abundances of some species, especially $^+$, in those astronomical regions which are not too warm (with kinetic temperatures below a few hundreds of degrees Kelvin) and which have a moderately high fraction of vibrationally excited $_2$ (in excess of $10^{-7}-10^{-6}$ )."
 Reactions of Ποιο> 0) with Ie!. OIT. and CN are less likely to affect the chemistry of interstellar clouds unless large amounts of vibrationally excited IT» are present.," Reactions of $_2$ $v>0$ ) with $^+$, O, OH, and CN are less likely to affect the chemistry of interstellar clouds unless large amounts of vibrationally excited $_2$ are present."
 Vibrationally excited molecular bydrogen is coumonly observed in PDRs. where the Πο ο>0 levels are excited by FUWV fluorescence. and in shocked gas where the excitation is mainiy collisional.," Vibrationally excited molecular hydrogen is commonly observed in PDRs, where the $_2$ $v>0$ levels are excited by FUV fluorescence, and in shocked gas where the excitation is mainly collisional."
 Amone these two types of regious the chemistry of vibrationally excited [> is expected to be less important iu shocked regions since kinetic temperatures are usually very high aud the IL level populations are not as far from thermalization as they are iu PDRs (Burton1992)., Among these two types of regions the chemistry of vibrationally excited $_2$ is expected to be less important in shocked regions since kinetic temperatures are usually very high and the $_2$ level populations are not as far from thermalization as they are in PDRs \citep{bur92}.
. Hereafter we thus focus ou the case of PDRs to investigate the effects of the reactions of vibrationally excited IT». onto the eas chemical composition., Hereafter we thus focus on the case of PDRs to investigate the effects of the reactions of vibrationally excited $_2$ onto the gas chemical composition.
 For that purpose we have utilized an updated version of the Meudou PDRcode'.. a photochemical nodel ofa a plane parallel aud. stationarv PDR (LePetitettal.20062008).. to compute the chewical and plivsica Structure as a function of the cloud depth for various PDR-like regions.," For that purpose we have utilized an updated version of the Meudon PDR, a photochemical model of a plane parallel and stationary PDR \citep{lep06,goi07,gon08}, to compute the chemical and physical structure as a function of the cloud depth for various PDR-like regions."
 Penetration oΕΙΝ radiation strongv depends on dust grain properties through dust absorption aud scattering of FUV photons which. iu addition. determine the efficiency of the donunaut heating mechanisuis of the easons photoelectric heating).," Penetration of FUV radiation strongly depends on dust grain properties through dust absorption and scattering of FUV photons which, in addition, determine the efficiency of the dominant heating mechanisms of the gas (photoelectric heating)."
 In this work we use standard dust properties. ie. erain sizes follow a power-law distribution (Alathisetal.1977:Coicoechea&LeBour-lot2007) and explicitly determine the eas teniperature eracdient by solving the cloud thermal balance (see LePetitetal.2006)).," In this work we use standard dust properties, i.e. grain sizes follow a power-law distribution \citep{mat77,goi07} and explicitly determine the gas temperature gradient by solving the cloud thermal balance (see \citealt{lep06}) )."
 We have used the chemical network described iu the previous section and the reactious of excited IT; shown in Table ??.., We have used the chemical network described in the previous section and the reactions of excited $_2$ shown in Table \ref{table-h2-reac}.
 For the reaction between Πω aud C! we adopt the rotational state-specific rate constant expressions (No 2 in Table 773) for II» (jy<7) and assume the Langevin value (No 3 in Table ??)3 for IT Cj>7 or ο>0). states for which the reaction becomes exothermic.," For the reaction between $_2$ and $^+$ we adopt the rotational state-specific rate constant expressions (No 2 in Table \ref{table-h2-reac}) ) for $_2$ $j\leq$ 7) and assume the Langevin value (No 3 in Table \ref{table-h2-reac}) ) for $_2$ $j>7$ or $v>0$ ), states for which the reaction becomes exothermic."
 Under thermal conditions this approach results ina value equivalent to the thermal rate coustaut (No 1 in Table ??)). at worst a factor of ~2 larger in the high temperature range (L000 In). where the rate constant starts to be dominated by Πο ¢>0 states.," Under thermal conditions this approach results in a value equivalent to the thermal rate constant (No 1 in Table \ref{table-h2-reac}) ), at worst a factor of $\sim$ 2 larger in the high temperature range $>$ 1000 K), where the rate constant starts to be dominated by $_2$ $v>0$ states."
 We cousider (de-jexcitation of IT» levels by collisious with IT». IT. and We (LeDourlotetal.1999).," We consider (de-)excitation of $_2$ levels by collisions with $_2$ , H, and He \citep{leb99}."
. Collisions with II! are also included based ou the recent work of IIuestis(2008).. although they are found to be not iuportaut.," Collisions with $^+$ are also included based on the recent work of \citet{hue08}, although they are found to be not important."
 Concerning collisions with clectrous. tle rates of vibrational de-cxcitation of Il; ο=1>0 are typically a few 19 cu s+ (from the cross sections recommended by Yoouetal. 2008)). which nuplies a critical deusitv of clectrons of ~ 10? 7.," Concerning collisions with electrons, the rates of vibrational de-excitation of $_2$ $v=1
\rightarrow 0$ are typically a few $^{-10}$ $^3$ $^{-1}$ (from the cross sections recommended by \citealt{yoo08}) ), which implies a critical density of electrons of $\sim$ $^3$ $^{-3}$."
" Iu low density regions. such as diffuse clouds. the electron density is much lower aud celectron-upact vibrational excitation is negligible. although im deuser regious witli a high ionization degree. such as planetary nebulae. clectrous could play a role iu producing nou-theriual II, vibrational populatious."," In low density regions, such as diffuse clouds, the electron density is much lower and electron-impact vibrational excitation is negligible, although in denser regions with a high ionization degree, such as planetary nebulae, electrons could play a role in producing non-thermal $_2$ vibrational populations."
 The implementation of (cle-Jexcitation of IT» bv collisious with electrons is however not straightforward as onlv a few state-to-state cross sections are available (Yoonetal.2008)., The implementation of (de-)excitation of $_2$ by collisions with electrons is however not straightforward as only a few state-to-state cross sections are available \citep{yoo08}.
". Wo are inainly interested ino investigating how Huportantis the effect of the chemistry of vibrationally excited Tl, iu the various regions.", We are mainly interested in investigating how importantis the effect of the chemistry of vibrationally excited $_2$ in the various regions.
 We thus have systematically run for the different sources a model iu which we adopt the thermal rate constaut expressions for all reactions of Πο iodel)) aud. another oue inchiding the state-specific rate constants for the reactions of excited IT» detailed in Table ?? model)).," We thus have systematically run for the different sources a model in which we adopt the thermal rate constant expressions for all reactions of $_2$ ) and another one including the state-specific rate constants for the reactions of excited $_2$ detailed in Table \ref{table-h2-reac}
 )."
 We first focus on a well known diffuse cloud such as ¢Oph. whose paracters (see Table ??)) have been taken from vanDishoeck&Black (1986)..," We first focus on a well known diffuse cloud such as $\zeta$Oph, whose parameters (see Table \ref{table-column-densities}) ) have been taken from \citet{van86}. ."
 Iu Fig., In Fig.
 3. we show as a function of Aq the abundance of some species while the, \ref{fig-zoph} we show as a function of $A_V$ the abundance of some species while the
enrichment of the system (White & Rees. 1015).,"enrichment of the system (White $\&$ Rees, 1978)."
 In this picture massive ellipticals form at. relatively recent epochs through major mergers between spiral ealaxies (e.g. I|xaulfimann & White. 1993) occurring over a large reclshilt interval.," In this picture massive ellipticals form at relatively recent epochs through major mergers between spiral galaxies (e.g. Kauffmann $\&$ White, 1993) occurring over a large redshift interval."
 Evidences favouring the hierarchical scenario are represented by the observed. interacting galaxies (on-going mergers). ellipticals with disturbed morphologies (i.e. counter-rotating cores. dust lanes. ripples. c.g. Ixormencdvy & Djorgovski. 1989). the morphology-density relation in clusters (Dressler et al.," Evidences favouring the hierarchical scenario are represented by the observed interacting galaxies (on-going mergers), ellipticals with disturbed morphologies (i.e. counter-rotating cores, dust lanes, ripples, e.g. Kormendy $\&$ Djorgovski, 1989), the morphology-density relation in clusters (Dressler et al.,"
 LOOT). the relatively large values of the £1; index in a non negligible fraction of nearby cllipticals (Gonzalez 1993. ‘Trager et al 1998).," 1997), the relatively large values of the $H_{\beta}$ index in a non negligible fraction of nearby ellipticals (Gonzalez 1993, Trager et al 1998)."
 We refer the reader to Paper lL. Peebles (2002). Calura et al. (," We refer the reader to Paper I, Peebles (2002), Calura et al. ("
2004) for a more detailed discussion on the topic and references.,2004) for a more detailed discussion on the topic and references.
 Llere we recall that in Paper E we showed that the majority of the optical properties of elliptical galaxies can be simultaneously. reproduced: under the assumption that the formation process. occurring at high redshift (as in the monolithic collapse scenario). is stronger and faster in more massive objects with respect to the less massive ones.," Here we recall that in Paper I we showed that the majority of the optical properties of elliptical galaxies can be simultaneously reproduced under the assumption that the formation process, occurring at high redshift (as in the monolithic collapse scenario), is stronger and faster in more massive objects with respect to the less massive ones."
 In fact. the κΑςΡΟcv] ratio in the cores of ellipticals increases with galactic mass (Worthey ct al.," In fact, the $[<Mg/Fe>_V]$ ratio in the cores of ellipticals increases with galactic mass (Worthey et al."
 902: Weiss et al., 1992; Weiss et al.
 1995: Ixuntschner 2000. Ixuntschner et. al..," 1995; Kuntschner 2000, Kuntschner et al.,"
 2001). being always larger than zero in the core of bright galaxies (c.g. Faber et al.," 2001), being always larger than zero in the core of bright galaxies (e.g. Faber et al.,"
 1992: Carollo et al..," 1992; Carollo et al.,"
 1993: Davies et al.," 1993; Davies et al.,"
 1993: Worthey et al.," 1993; Worthey et al.,"
 1992)., 1992).
 This strongly suggests that the star formation (SE) lasted for a period shorter than the time at which the Fe pollution from tvpe Ia SNe becomes dominant (see Weiss et al..," This strongly suggests that the star formation (SF) lasted for a period shorter than the time at which the Fe pollution from type Ia SNe becomes dominant (see Weiss et al.,"
 1995. Matteucci. 1994).," 1995, Matteucci, 1994)."
 οσον. a similar suggestion was proposed by Cranato et al. (," Recently, a similar suggestion was proposed by Granato et al. ("
2004) in the scheme of the joint formation of ellipticals and QSOs.,2004) in the scheme of the joint formation of ellipticals and QSOs.
 In order to complete the studs. we aim at. verifing whether Paper LE conclusions can be achieved with a more complicated formation history. namely t0 study in a simple wav the elfect of merger induced starbursts on the properties ofa dominant stellar population assembled at high redshift.," In order to complete the study, we aim at verifing whether Paper I conclusions can be achieved with a more complicated formation history, namely to study in a simple way the effect of merger induced starbursts on the properties of a dominant stellar population assembled at high redshift."
" A complementary. view on the analysis made in this paper is presented by Thomas (1999) and ""Ehomas et al. (", A complementary view on the analysis made in this paper is presented by Thomas (1999) and Thomas et al. (
1999).,1999).
 In »wticular. they compared the outcome of a fast collapse to he results obtained by the merger of two spirals ancl found difficult to reproduce the «AMg/Fezv] overabundance in he latter case. unless very [at LIMES are invoked during he merger.," In particular, they compared the outcome of a fast collapse to the results obtained by the merger of two spirals and found difficult to reproduce the $[<Mg/Fe>_V]$ overabundance in the latter case, unless very flat IMFs are invoked during the merger."
ὃν Furthermore. since the ςAlg/ke pe relation seems to be independent from the environment and correlates with galactic age (the older is the galaxy the righer is its «Mg/Pe: v]. Thomas et al (2002) concluded hat SE histories (SELIs) based on the hierarchical clustering scenario (Ixaulfmann Charlot. 1998) cannot. reproduce he observed chemical. properties of elliptical galaxies.," Furthermore, since the $[<Mg/Fe>_V]$ $\sigma$ relation seems to be independent from the environment and correlates with galactic age (the older is the galaxy the higher is its $[<Mg/Fe>_V]$ ), Thomas et al (2002) concluded that SF histories (SFHs) based on the hierarchical clustering scenario (Kauffmann Charlot, 1998) cannot reproduce the observed chemical properties of elliptical galaxies."
 On he other hand. by means of broad-band photometry alone. Shiova Bekki (1998) constrained the epoch of the major merger.," On the other hand, by means of broad-band photometry alone, Shioya Bekki (1998) constrained the epoch of the major merger."
 They found that mergers between spirals are allowed even at 250.93. but their models cannot account for the Ale enhancement.," They found that mergers between spirals are allowed even at z=0.3, but their models cannot account for the Mg enhancement."
 Therefore. it seems necessary to couple the information coming from the chemistry to those [rom the broad-band colours in order to better constrain the SELL of spheroids.," Therefore, it seems necessary to couple the information coming from the chemistry to those from the broad-band colours in order to better constrain the SFH of spheroids."
 Due to the fact that we are not. modelling he dynamical evolution. our treatment of the galaxy merging is quite simple and focused only to study how the mean chemical and photometric properties of the stellar »opulations change as a consequence of late episodes of star ormation triggered. by gascous mergers.," Due to the fact that we are not modelling the dynamical evolution, our treatment of the galaxy merging is quite simple and focused only to study how the mean chemical and photometric properties of the stellar populations change as a consequence of late episodes of star formation triggered by gaseous mergers."
" In. particular. we analvsed: dillerent. scenarios involving high mass spheroids (i.c. initial gas mass My,=10422A/.). as predicted: by ""aper Es best model. that we consider ourcasc.. and we took it as a basis for studying simulated aceretion of »imordial eas trigeering a starburst."," In particular, we analysed different scenarios involving high mass spheroids (i.e. initial gas mass $M_{lum}=10^{11-12}M_{\odot}$ ), as predicted by Paper I's best model, that we consider our, and we took it as a basis for studying simulated accretion of primordial gas triggering a starburst."
 We explored. several cases by varving the duration of the starburst and. its intensity., We explored several cases by varying the duration of the starburst and its intensity.
 This paper is organized. as follows: in section 2 we present the model., This paper is organized as follows: in section 2 we present the model.
 In section 3 we discuss the results and in section 4 we draw some conclusions., In section 3 we discuss the results and in section 4 we draw some conclusions.
" In order to be consistent with the results of Paper L. we assume a cosmological model with O,,=0.3. O4=0.7 and Hf,=TOkms‘Alpei"," In order to be consistent with the results of Paper I, we assume a cosmological model with $\Omega_m=0.3$, $\Omega_{\Lambda}=0.7$ and $H_o=70 \rm \,km\, s^{-1}\, Mpc^{-1}$."
 For all models we assume that the star formation starts at redshift 2/25., For all models we assume that the star formation starts at redshift $z_f$ =5.
" Variations in 2, clo not significantly allect the conclusions.", Variations in $z_f$ do not significantly affect the conclusions.
 The adopted: chemical. evolution. model is based on that presented in Paper L In this particular case. however. we consider our model galaxies as a single zone extending out to one cllective radius. with instantaneous mixing of gas.," The adopted chemical evolution model is based on that presented in Paper I. In this particular case, however, we consider our model galaxies as a single zone extending out to one effective radius, with instantaneous mixing of gas."
 Moreover we take explicitly into account a possible mass [ow due to the galactic wind and a possible secondary. episode of eas accretion., Moreover we take explicitly into account a possible mass flow due to the galactic wind and a possible secondary episode of gas accretion.
 Therefore. the equations of chemical evolution take the following form: where (0) is the normalized. mass density of the clement at the time in the ISM.," Therefore, the equations of chemical evolution take the following form: where $G_i (t)$ is the normalized mass density of the element at the time in the ISM."
 Αν) is defined. as the abundance by mass of the clement7., $X_i (t)$ is defined as the abundance by mass of the element.
" We address the reader to the original papers of Matteucci & Cireggio (1986). Alatteucci & ""Tornambe! (LOST). Gibson (1997) and Paper L for à comprehensive discussion. of these equations."," We address the reader to the original papers of Matteucci $\&$ Greggio (1986), Matteucci $\&$ Tornambe' (1987), Gibson (1997) and Paper I for a comprehensive discussion of these equations."
 By means of this equation we can calculate the evolution of 2] elemental species., By means of this equation we can calculate the evolution of 21 elemental species.
 In particular. the integrals at the right-hand. side of the equations give the rate at which the clement is restored into the interstellar medium both as newlv-svnthesized and already existing element by low- and," In particular, the integrals at the right-hand side of the equations give the rate at which the element is restored into the interstellar medium both as newly-synthesized and already existing element by low- and"
The MZ relation is best fitted by with the Lo errors in cach of the parameters in equation NS even bv ay=(007.050.0L and σε=0.08.,"The MZ relation is best fitted by with the $1\sigma$ errors in each of the parameters in equation \ref{eq:mzfit} given by $\sigma_{A} = 0.007, \sigma_{B} = 0.01$ and $\sigma_{C} = 0.03$."
 Figure 6. shows the DEEP2 MZ relation at 5~OLS conrpared to the local MZ relation from SDSS (see appendix)., Figure \ref{fig:mz} shows the DEEP2 MZ relation at $z\sim0.8$ compared to the local MZ relation from SDSS (see appendix).
 The black filled circles are the data sorted into 15 mass bins., The black filled circles are the data sorted into 15 mass bins.
 Each bin contains 90 data points aud the errors on the data are determined by raucdomly selecting 90 data poiuts in each bin aud taking the median of the subsample., Each bin contains $\sim 90$ data points and the errors on the data are determined by randomly selecting $\sim90$ data points in each bin and taking the median of the subsample.
 We take the standard deviation of this distribution as the error in cach bin., We take the standard deviation of this distribution as the error in each bin.
 We note that this scatter characterizes the distribution of metallicities within each mass bin and therefore accounts for both lueasurement uncertainties and intrinsic scatter in the mnctallicity determination., We note that this scatter characterizes the distribution of metallicities within each mass bin and therefore accounts for both measurement uncertainties and intrinsic scatter in the metallicity determination.
 The erev lines are the 16% aud SI percentile contours of the data., The grey lines are the $16\%$ and $84\%$ percentile contours of the data.
 The large dispersion at the high mass cud is due to the large iuterval im mass covered bv the highest mass bius., The large dispersion at the high mass end is due to the large interval in mass covered by the highest mass bins.
" Overall. we fud that the high-ass (AL>ο”. DEEP? salaxies have similar metallicities (within the errors of £0.05 dex) to local galaxies at the same stellar nass,"," Overall, we find that the high-mass $M > 10^{10.5} M_\odot$ ) DEEP2 galaxies have similar metallicities (within the errors of $\pm 0.05$ dex) to local galaxies at the same stellar mass."
 At lower masses (AL<LO! oar). the DEEP? ealaxies have lower moetallicities than the SDSS sample for a given stellar mass. up to a Aog(O/IT)~0.15 dex.," At lower masses $M < 10^{10.5} M_\odot$ ), the DEEP2 galaxies have lower metallicities than the SDSS sample for a given stellar mass, up to a $\Delta log(O/H) \sim 0.15$ dex."
 Fiewe 7T shows the LZ velation derived frou the DEEP? sample aud the local sample from SDSS., Figure \ref{fig:lz} shows the LZ relation derived from the DEEP2 sample and the local sample from SDSS.
" Similar to the MZ relation. the LZ relation is determined for ealaxics with AL,>10°?AL..."," Similar to the MZ relation, the LZ relation is determined for galaxies with $M_{\ast} > 10^{9.2}M_{\odot}$."
 The filled circles are the aedian inetalliifies mo 15 Lhunuinositv bius., The filled circles are the median metallicities in 15 luminosity bins.
 The solidcurve is a linear fit to the biuned data and is parameterized by the equation where Vp=Mp|21., The solidcurve is a linear fit to the binned data and is parameterized by the equation where $X_B = M_B + 21$.
 The errors in the ft aud binued data of the LZ relation are derived in the same nauner as or the MZ relation., The errors in the fit and binned data of the LZ relation are derived in the same manner as for the MZ relation.
 The ervey lines are the 1654 aud 8X rcentile contours of the data., The grey lines are the $16\%$ and $84\%$ percentile contours of the data.
 The dashed curve is the ocal LZ relation froin SDSS taken from the appendix of this paper., The dashed curve is the local LZ relation from SDSS taken from the appendix of this paper.
 The ft made is over the full range of uaenitudes 21Mpx 15) of the local LZ relation., The fit made is over the full range of magnitudes $-21<M_{B}<-18$ ) of the local LZ relation.
 However. the local LZ relation appears to turn over at Ap<—20.," However, the local LZ relation appears to turn over at $M_{B}<-20$."
 The dot-dashed curve is a fit to the linear xut of the local LZ relation with Afp>20., The dot-dashed curve is a fit to the linear part of the local LZ relation with $M_{B}> -20$.
 The slope of this fit aud the DEEP2 LZ relation are consistent o within the errors., The slope of this fit and the DEEP2 LZ relation are consistent to within the errors.
 KNOL find that the slope of the LZ relation for ;=(0.6O8 and 2=US.LOU are O.1172EO.017 and =O.A3140.032 respectively., KK04 find that the slope of the LZ relation for $z = 0.6-0.8$ and $z = 0.8-1.0$ are $-0.117\pm0.017$ and $-0.134\pm0.032$ respectively.
 To within he errors. the slopes derived by KI&01 agree with DEEP? and local LZ relation slopes as well.," To within the errors, the slopes derived by KK04 agree with DEEP2 and local LZ relation slopes as well."
 Iu this section we will discuss our key assuniptions (55.1). the evolution of the B-hand Iuninositv as a function of stellar mass in the population of blue oOealaxies (55.2) aud compare our determination of the MZ relation with other authors (65.3).," In this section we will discuss our key assumptions $\S \, 5.1$ ), the evolution of the B-band luminosity as a function of stellar mass in the population of blue galaxies $\S \, 5.2$ ) and compare our determination of the MZ relation with other authors $\S \, 5.3$ )."
 We list and address some key assumptions mado in tliis paper: 1., We list and address some key assumptions made in this paper: 1.
" All ealaxies with AL,>10°74S. are on the upper brauch of the Ro; diagnostic.", All galaxies with $M_\ast \geq 10^{9.2} M_\odot$ are on the upper branch of the $R_{23}$ diagnostic.
 At the lower mass cud of the DEEP? MZ relation the dispersion is simall aud consistent with the dispersion found in the local MZ relation (TOL)., At the lower mass end of the DEEP2 MZ relation the dispersion is small and consistent with the dispersion found in the local MZ relation (T04).
" This simall dispersion sugeests that at Do(US. the vast majority of galaxies with AL,>10°F are on the upper mietallicity branch of the Ro; diagnostic."," This small dispersion suggests that at $z\sim0.8$, the vast majority of galaxies with $M_\ast > 10^{9.2} M_\odot$ are on the upper metallicity branch of the $R_{23}$ diagnostic."
" Additionally. TOL have observed a coutinuous slope down o AL,~105?3/.."," Additionally, T04 have observed a continuous slope down to $M_\ast \sim 10^{8.5} M_\odot$."
 A breakdown of this assmuptiou would result in the Που measure of metallicity in low uetallicity galaxies to be significantly overestimated (ldex) resulting in a flattenine of the MZ relation at ower duasses., A breakdown of this assumption would result in the $R_{23}$ measure of metallicity in low metallicity galaxies to be significantly overestimated $\sim 1$ dex) resulting in a flattening of the MZ relation at lower masses.
 This flattening is observed to occur at AL.=1092A/. aud eiven the small dispersion. we wpothesize that musplacement on the upper brauch is nünor above this mass.," This flattening is observed to occur at $M_{\ast} = 10^{9.2}M_\odot$ and given the small dispersion, we hypothesize that misplacement on the upper branch is minor above this mass."
 Near infrared spectroscopy o obtain additional line ratios for a represcutative subsample of DEEP2 ealaxies is needed to test this ivpothliesis., Near infrared spectroscopy to obtain additional line ratios for a representative subsample of DEEP2 galaxies is needed to test this hypothesis.
 2., 2.
 Metals ire iustautaneouslv recveled., Metals are instantaneously recycled.
 “Thuis approximation was first introduced— bx 7., This approximation was first introduced by \citet{Schmidt1963}.
 The approximation is valid when metals produced in lnassive stars. such as oxvgen. are cousidered and if the SER is not subject to extreme variations on short timescale (?)..," The approximation is valid when metals produced in massive stars, such as oxygen, are considered and if the SFR is not subject to extreme variations on short timescale \citep{Pagel1997}."
" Frou, their study of DEEP2 galaxies suce c=LI1.7 o determine that the dominant foriu of evolution is a gradual decline in star formation."," From their study of DEEP2 galaxies since $z=1.1$, \citet{Noeske2007a} determine that the dominant form of evolution is a gradual decline in star formation."
 3., 3.
 We have attempted to remove ACN using enpirical luits on the Που. ratio aud the color binodality., We have attempted to remove AGN using empirical limits on the $R_{23}$ ratio and the color bimodality.
 This removes most of the galaxies whose cussion is dominated Jy ACN. but composite galaxies iu our saluple may be blue aud not exceed the eupirical Ro» limit.," This removes most of the galaxies whose emission is dominated by AGN, but composite galaxies in our sample may be blue and not exceed the empirical $R_{23}$ limit."
 AGN contamination in composite galaxy spectra increases the Roy ratio due to the high HH> ratio., AGN contamination in composite galaxy spectra increases the $R_{23}$ ratio due to the high $[OIII]/H\beta$ ratio.
 Ou the upper metallicity branch of the|OZTTI]|/ Ro; diagnostic. the aetallicity is a monotonically decreasing function of Ros.," On the upper metallicity branch of the $R_{23}$ diagnostic, the metallicity is a monotonically decreasing function of $R_{23}$."
 Therefore. ACN contamination will lower the metallicity estimate.," Therefore, AGN contamination will lower the metallicity estimate."
 However. due to their relatively siadl numbers in the DEEP? blue galaxy sample (7?) and the simall dispersion in the MZ relation. ACN contanunation is likely to play a minor role iu the MZ velation (7)..," However, due to their relatively small numbers in the DEEP2 blue galaxy sample \citep{Weiner2007} and the small dispersion in the MZ relation, AGN contamination is likely to play a minor role in the MZ relation \citep{Lamareille2009}."
 Near infrared spectroscopy to obtain additional AGN-sensitive cuiissiou-line ratios is required o verify ΣΕΠΠ 1., Near infrared spectroscopy to obtain additional AGN-sensitive emission-line ratios is required to verify this assumption.
 , 4.
Since its introduction by ?.. the iuifial stellar nass ποοι (IME) - describing the mass distribution of stars at birth - has larecly been taken to be uuiversal and invariant.," Since its introduction by \citet{Salpeter1955}, , the initial stellar mass function (IMF) - describing the mass distribution of stars at birth - has largely been taken to be universal and invariant."
 Dadirect evidences for variations in the ocal IME have been provided by the comparison of fa cluission to the far-ultraviolet flux (?) and the Ifa Hux to broadband color iudex (?).., Indirect evidences for variations in the local IMF have been provided by the comparison of $H\alpha$ emission to the far-ultraviolet flux \citep{Meurer2009} and the $H\alpha$ flux to broadband color index \citep{Hoversten2008}.
 Comparisons of the rate of buuinositv evolution to that of color evolution xovides indirect observational evidence for a possible redshift evolution iu the IME. (?).., Comparisons of the rate of luminosity evolution to that of color evolution provides indirect observational evidence for a possible redshift evolution in the IMF \citep{vanDokkum2008}.
 Furthermore. IME variatious have been cited as the possible cause of the MZ relation (?7)..," Furthermore, IMF variations have been cited as the possible cause of the MZ relation \citep{Koppen2007}. ."
 We work uuder the assumption that the IMPE is universal., We work under the assumption that the IMF is universal.
 5., 5.
 MZ aud LZ relation studies assmue that the absence of low surface brightness ealaxies (LSBCs) do uot sienificautly bias the deteriuünuation of the MZ aud LZ relations., MZ and LZ relation studies assume that the absence of low surface brightness galaxies (LSBGs) do not significantly bias the determination of the MZ and LZ relations.
 ?point out that LSBCs will be absent, \citet{Impey1997} point out that LSBGs will be absent
10312 (DPI: S. Hawley).,10312 (PI: S. Hawley).
 The Space Telescope Luagine Spectrograph (STIS. Iximball οἱ al.," The Space Telescope Imaging Spectrograph (STIS, Kimball et al."
 1993). was used with the medium resolution NUV echelle grating (I22320M) ancl the 0.2x0.2 arcsec aperture.," 1998), was used with the medium resolution NUV echelle grating (E230M) and the 0.2x0.2 arcsec aperture."
 The E230M spectra have a nearly constant (£1.5'%.)) dispersion of 4.88 km/s per pixel., The E230M spectra have a nearly constant $\pm 1.5$ ) dispersion of 4.88 km/s per pixel.
 The line spread function (LSF) for the 0.2x0.2 arcsec aperture has a EWILIM of 1.3-1.9 pixels over the wavelength range2400-2800A.. corresponding to a resolving power of R=33.000.," The line spread function (LSF) for the 0.2x0.2 arcsec aperture has a FWHM of 1.8-1.9 pixels over the wavelength range, corresponding to a resolving power of R=33,000."
 These are the only STIS spectra of an M cdiwai flare star al NUV wavelengths that exist in the LST archive., These are the only STIS spectra of an M dwarf flare star at NUV wavelengths that exist in the HST archive.
 Data were obtained during six consecutive orbits. wilh exposure limes ranging from 32 to 40 minutes. over a 6.5 hour total timespan.," Data were obtained during six consecutive orbits, with exposure times ranging from 32 to 40 minutes, over a 6.5 hour total timespan."
 The MAMA detector timetag mode was used for the observations. resulting in an event list for each exposure that contains the detection time aud location of every detected photon.," The MAMA detector timetag mode was used for the observations, resulting in an event list for each exposure that contains the detection time and location of every detected photon."
 To construct lieht. curves of the integrated. light obtained. during each exposure. we used an extraction window that was ll pixels wide. centered on each echelle order. to assign all timetag events to either source or background regions on the ALANA detector.," To construct light curves of the integrated light obtained during each exposure, we used an extraction window that was 11 pixels wide, centered on each echelle order, to assign all timetag events to either source or background regions on the MAMA detector."
 We then divided each exposure into 100 uniform time bins (19-24 s in duration dependiusg on the individual exposure lengths which range from 32-40 min. see Table 1). obtaining separate lieht curves for the source and backeround regions.," We then divided each exposure into 100 uniform time bins (19-24 s in duration depending on the individual exposure lengths which range from 32-40 min, see Table 1), obtaining separate light curves for the source and background regions."
 For each exposure. we fitted the background lisht curve with a sixth order polynomial. scaled the result bv the ratio οἱ source Lo background pixels. and (then subtracted this predicted background from (he source light curve to obtain a total observed stellar count rate for each time bin.," For each exposure, we fitted the background light curve with a sixth order polynomial, scaled the result by the ratio of source to background pixels, and then subtracted this predicted background from the source light curve to obtain a total observed stellar count rate for each time bin."
 Errors for each time bin include Poisson noise from the subtracted background., Errors for each time bin include Poisson noise from the subtracted background.
 Figure | illustrates the resulting light curves for (he six exposures., Figure 1 illustrates the resulting light curves for the six exposures.
 The fLIaring; and quiet intervals F1-10 and QI-8 were assigned by qualitative inspection of the light curves. and are labeled on the figure.," The flaring and quiet intervals F1-10 and Q1-8 were assigned by qualitative inspection of the light curves, and are labeled on the figure."
 Table 1 provides relevant information about each spectral interval., Table 1 provides relevant information about each spectral interval.
 Column 1 lists the number of the exposure containing the time interval., Column 1 lists the number of the exposure containing the time interval.
 Exposure number N maps to LST dataset 059&010NO., Exposure number N maps to HST dataset o59k010N0.
 The second and third columnis give the start aud end lime of each interval. expressed in minutes since the beginning of each exposure.," The second and third columns give the start and end time of each interval, expressed in minutes since the beginning of each exposure."
 Column 4 gives the label assigned to each lime interval. with Q labels indicating «quiet. intervals and F labels indicating flare intervals.," Column 4 gives the label assigned to each time interval, with Q labels indicating quiet intervals and F labels indicating flare intervals."
 The numbers were assigned sequentially in (me [ου individual intervals., The numbers were assigned sequentially in time for individual intervals.
 Columns 5 and 6 give the total count rate and standard: deviation lor each interval., Columns 5 and 6 give the total count rate and standard deviation for each interval.
 Although the Marine ancl quiet intervals were assigned qualitatively. it is clear from these data that the flaring intervals have both higher count rates ancl larger scatter.," Although the flaring and quiet intervals were assigned qualitatively, it is clear from these data that the flaring intervals have both higher count rates and larger scatter."
 The exceptions are the late decay phases of the (wo laree flares. labeled as F2¢ and F9c.," The exceptions are the late decay phases of the two large flares, labeled as F2c and F9c."
 Since we wanted to investigate the evolution of the emission lines in (hese late phases. we retain {hose as flare intervals.," Since we wanted to investigate the evolution of the emission lines in these late phases, we retain those as flare intervals."
 Also note that quiet interval Q7 has a relatively high count rate and shows significant low level variability in the light curve. but no obvious flares.," Also note that quiet interval Q7 has a relatively high count rate and shows significant low level variability in the light curve, but no obvious flares."
 We include it as a (quiet interval after insuring that the QT spectrum shows no more than variation compared to the summed quiet spectrum over all other intervals., We include it as a quiet interval after insuring that the Q7 spectrum shows no more than variation compared to the summed quiet spectrum over all other intervals.
 Since QT represents almost, Since Q7 represents almost
Figure 4 presents the frequency analysis of the regular component (Fig.,Figure 4 presents the frequency analysis of the regular component (Fig.
 4a) and of the chaotic component (Fig., 4a) and of the chaotic component (Fig.
 4b) of the system., 4b) of the system.
 We are only interested in revealing the “disc” resonances on the rotation plane and therefore the frequency analysis is made in the 2D projection of the orbits on the y—z plane., We are only interested in revealing the “disc” resonances on the rotation plane and therefore the frequency analysis is made in the 2D projection of the orbits on the $y-z$ plane.
" However, there exist also vertical resonances that influence mainly the edge-on profiles of our systems (seePatsisetal.2002)."," However, there exist also vertical resonances that influence mainly the edge-on profiles of our systems \citep[see][]{b22}."
. The “frequency ratio” q is a parameter given by the relation where Q is the angular velocity of a particle in the inertial frame of reference and κ. is the epicyclic (radial) frequency (Binney&Tremaine1987)., The “frequency ratio” $q$ is a parameter given by the relation where $\Omega$ is the angular velocity of a particle in the inertial frame of reference and $\kappa$ is the epicyclic (radial) frequency \citep{b31}.
". Positive q values correspond to resonances inside corotation while negative q values correspond to resonances outside corotation, which are related only to chaotic orbits."," Positive $q$ values correspond to resonances inside corotation while negative $q$ values correspond to resonances outside corotation, which are related only to chaotic orbits."
 In Fig., In Fig.
" 4a we see that the main resonance of the regular orbits (inside corotation) is the 2:1 resonance (or inner Lindblad resonance), while there is also a small number of particles around the 3:1 resonance and an even smaller number around corotation (q—0)."," 4a we see that the main resonance of the regular orbits (inside corotation) is the 2:1 resonance (or inner Lindblad resonance), while there is also a small number of particles around the 3:1 resonance and an even smaller number around corotation $(q=0)$."
 The vast majority of these orbits support the bar of the system., The vast majority of these orbits support the bar of the system.
" Figure 4b shows that a rather small percentage of chaotic orbits inside corotation lie near important resonances (e.g. 2:1, 3:1 and 4:1)."," Figure 4b shows that a rather small percentage of chaotic orbits inside corotation lie near important resonances (e.g. 2:1, 3:1 and 4:1)."
 These represent sticky chaotic orbits mainly supporting the outer layers of the bar (seeHar-soula&Kalapotharakos2009) and the innermost parts of the spiral structure., These represent sticky chaotic orbits mainly supporting the outer layers of the bar \citep[see][]{b8} and the innermost parts of the spiral structure.
" The rest of the chaotic orbits move around corotation and outside it and some of them stick around specific resonances (q—0, -0.5, -1, -1.5, -2, -2.5, -"," The rest of the chaotic orbits move around corotation and outside it and some of them stick around specific resonances $q$ =0, -0.5, -1, -1.5, -2, -2.5, -3)."
" The gq=0 population corresponds to chaotic orbits located around corotation and close to the PLi, PL2, PL and PL; families (nomenclatureafterVoglisetal.2006a)."," The $q=0$ population corresponds to chaotic orbits located around corotation and close to the $PL_1$, $PL_2$, $PL_4$ and $PL_5$ families \citep[nomenclature after][]{b6}."
. Below we will show that the chaotic orbits sticking along the asymptotic manifolds originating from these unstable periodic orbits are responsible for the spiral structure of the system., Below we will show that the chaotic orbits sticking along the asymptotic manifolds originating from these unstable periodic orbits are responsible for the spiral structure of the system.
 In this section we discuss the role of apocentres or pericentres in revealing the morphological features of the system and we relate their geometrical loci with the locus of the velocity minima on the rotation plane., In this section we discuss the role of apocentres or pericentres in revealing the morphological features of the system and we relate their geometrical loci with the locus of the velocity minima on the rotation plane.
" The distribution of the radial velocities 7 of the N-body orbits has a preferable concentration around the value zero, indicating that the apocentres and pericentres should play an important role for the system."," The distribution of the radial velocities $\dot{r}$ of the $N$ -body orbits has a preferable concentration around the value zero, indicating that the apocentres and pericentres should play an important role for the system."
" The geometrical loci of the minima of the velocities vyz on the rotation plane should correspond to the maxima of the density distribution on the configuration space, because particles spend most of their time there."," The geometrical loci of the minima of the velocities $v_{yz}$ on the rotation plane should correspond to the maxima of the density distribution on the configuration space, because particles spend most of their time there."
 It is of interest to investigate how these loci are related to the apocentres or the pericentres of the orbits., It is of interest to investigate how these loci are related to the apocentres or the pericentres of the orbits.
 In Fig., In Fig.
 5 we plot the effective potential of a 2D, 5 we plot the effective potential of a 2D
" The evolution of millisecond pulsars can be consistently described in terms of the equilibrium period distribution (D) of MSXPs at the end of the LMXB evolution the mass accretion rates (M) of the progenitor population during the recycling process Galactic birth rates (R), and the dominant energy loss mechanism after the onset of radio emission.","							 The evolution of millisecond pulsars can be consistently described in terms of the equilibrium period distribution $D$ ) of MSXPs at the end of the LMXB evolution the mass accretion rates $\dot{M}$ ) of the progenitor population during the recycling process Galactic birth rates $R$ ), and the dominant energy loss mechanism after the onset of radio emission."
" We devise a semi-analytical evolution function € to parametrize the evolution of millisecond pulsars after the accretion phase, which can be described in closed form as: where PDF is the probability distribution function."," We devise a semi-analytical evolution function $\mathcal{E}$ to parametrize the evolution of millisecond pulsars after the accretion phase, which can be described in closed form as: where $\mathcal{PDF}$ is the probability distribution function."
" The shape parameters a and £ define the distributions (i.e., D,M,R for k—1,2,3) for the Beta functiond!]citepEHP00 inferred from observations at each Monte-Carlo realization 664499T."," The shape parameters $\alpha$ and $\beta$ define the distributions (i.e., $D,\dot{M}, R$ for k=1,2,3) for the Beta \\citep{EHP00} inferred from observations at each Monte-Carlo realization “r”."
" The evolution function £ is built by randomly choosing initiation seeds from a period distribution D, which is then convolved via the standard model to consequently sample the P— parameter space."," The evolution function $\mathcal{E}$ is built by randomly choosing initiation seeds from a period distribution $D$, which is then convolved via the standard model to consequently sample the $P-\dot{P}$ parameter space."
" For the observed MSXPs, the period distribution which seeds will be randomly chosen from is the observed Pysxp distribution (table 1)."," For the observed MSXPs, the period distribution which seeds will be randomly chosen from is the observed $P_{MSXP}$ distribution (table 1)."
" We uniquely construct a “relaxed multidimensional Kolmogorov-Smirnov (K-S) filter"" (fig. 3) ", We uniquely construct a “relaxed multidimensional Kolmogorov-Smirnov (K-S) filter” (fig. \ref{fig:probdist}) )
to check population consistencies by calculating the 2D K-S probabilities (Pops) between observed MSRPs and the synthetic population that is formed by these properly evolved progenitor seeds., to check population consistencies by calculating the 2D K-S \citep{FF87} probabilities $P_{2DK-S}$ ) between observed MSRPs and the synthetic population that is formed by these properly evolved progenitor seeds.
 T'he filtering is reiterated for each realization to obtain synthetic populations with consistent distributions as: which is then used to construct the PDF in Equation refstat.eq.., The filtering is reiterated for each realization to obtain synthetic populations with consistent distributions as: which is then used to construct the $\mathcal{PDF}$ in Equation \\ref{stat.eq}. .
 Nominally any PepxK—s>0.2 value would imply consistent populations in a 2D K-S test., Nominally any $P_{2DK-S} > 0.2$ value would imply consistent populations in a 2D K-S test.
 By allowing 0.005<P5pg.s0.2 with lower fractions (see fig. By.," By allowing $0.005 < P_{2DK-S} < 0.2$ with lower fractions (see fig. \ref{fig:probdist})),"
 we oversample outliers to, we oversample outliers to
for the 1.5m. only displaced by a certain amount. (,"for the 1.5m, only displaced by a certain amount. ("
This displacement could in principle be determined from the data. but not verv reliably given the small numbers involved. and in anv case (he inaccuracy of the original measurements argues against (rving to be too precise.,"This displacement could in principle be determined from the data, but not very reliably given the small numbers involved, and in any case the inaccuracy of the original measurements argues against trying to be too precise."
 We will take it to be one magnitude.), We will take it to be one magnitude.)
 We mav now write an ileralive expression [or the telescope selection functions: We probably do not know (he ratio of total numbers well at this faint end of the ΠΙΟΓΙΟ. so we will leave (hem as a normalization constant to be determined.," We may now write an iterative expression for the telescope selection functions: We probably do not know the ratio of total numbers well at this faint end of the function, so we will leave them as a normalization constant to be determined."
 Taking only extragalactic objects. their histograms are shown in Figure (11)).," Taking only extragalactic objects, their histograms are shown in Figure \ref{nsraw}) )."
 We are clearlv going to have trouble with the small number of objects per bin., We are clearly going to have trouble with the small number of objects per bin.
 We will Lake two measures to deal with that: first. smooth the data with a tliree-bin ruining top hat ‘either before or after carrving out the division): second. demand that (he selection function be flat al brighter magnitudes (as long as an object is bright enough. it should be seen) and nonotonicallv declining thereafter.," We will take two measures to deal with that: first, smooth the data with a three-bin running top hat (either before or after carrying out the division); second, demand that the selection function be flat at brighter magnitudes (as long as an object is bright enough, it should be seen) and monotonically declining thereafter."
 The ratio of southern to northern objects fIuctuates in the bins from 23.0 to about 23.6 nagnitudes per square arc second. depending on the smoothing: partly Chis is due to statistics. but possibly also due to a somewhat ragged bright limit on our candidates.," The ratio of southern to northern objects fluctuates in the bins from 23.0 to about 23.6 magnitudes per square arc second, depending on the smoothing; partly this is due to small-number statistics, but possibly also due to a somewhat ragged bright limit on our candidates."
 At 23.8 to 24.0 it steadies. and declines thereafter.," At 23.8 to 24.0 it steadies, and declines thereafter."
 We (take this to indicate that there is no brightness-dependent selection effect due to the 1.511 up to 24.0. ancl so none for the 2.1 up to 25.0.," We take this to indicate that there is no brightness-dependent selection effect due to the 1.5m up to 24.0, and so none for the 2.1m up to 25.0."
 The next step is to apply our derived. fiction (o a set of observed data: and reconstruct. as [ar as we can. the histogram due to the original brightness distribution and the eve-and-plate selection function.," The next step is to apply our derived function to a set of observed data and reconstruct, as far as we can, the histogram due to the original brightness distribution and the eye-and-plate selection function."
 As noted by the referee. our measured surface brightnesses include Galactic extinction.," As noted by the referee, our measured surface brightnesses include Galactic extinction."
" This would not be important Lor our purposes if it affected all objects impartially,", This would not be important for our purposes if it affected all objects impartially.
 Unfortunately. we find. (using extinction measures [rom Schlegeletal.(1998) via NED) that the more heavily extincted objects tend to be brighter.," Unfortunately, we find (using extinction measures from \citet{SF98} via NED) that the more heavily extincted objects tend to be brighter."
 This makes sense: a faint object will be harder to find in a region of high extinction because of confusion with bright Galactic nebulosity., This makes sense: a faint object will be harder to find in a region of high extinction because of confusion with bright Galactic nebulosity.
 We could in principle correct for Chis effect statistically. but the small number of objects we have woull make such a correction very uncertain.," We could in principle correct for this effect statistically, but the small number of objects we have would make such a correction very uncertain."
 Instead. we simply omit those objects whose extinction is larger than the uneertainty in their measured surface brightness.," Instead, we simply omit those objects whose extinction is larger than the uncertainty in their measured surface brightness."
 This is equivalent to confining our analvsis to areas of (he skv well clear of the Milky Wavy., This is equivalent to confining our analysis to areas of the sky well clear of the Milky Way.
 We now take the extragalactic northern objects. minus those with high extinction. aud apply a three-bin running top hat smoothing.," We now take the extragalactic northern objects, minus those with high extinction, and apply a three-bin running top hat smoothing."
 At the 25.0 magnitudes per square arc second bin and fainter we apply a correction for the 2.1nm selection function., At the 25.0 magnitudes per square arc second bin and fainter we apply a correction for the 2.1m selection function.
 The result is shown in, The result is shown in
"We have demonstrated our method on data from the UKIDSS LAS, combining morphology statistics measured in the Y, J, H and K bands (or whatever subset of these a source was detected in).","We have demonstrated our method on data from the UKIDSS LAS, combining morphology statistics measured in the $Y$ , $J$ $H$ and $K$ bands (or whatever subset of these a source was detected in)."
" The morphology statistic used, veyswithsimilarcitepIrwinetal:2010, , representsapowerf ulmeansof datacompr eseigiegaomitbedilkipiaged candegnéinitsalchoatelakeituac (pein formation forth cclassi ficationtechniques)."," The morphology statistic used, \\citep{Irwin_etal:2010}, represents a powerful means of data compression from the full image, and contains almost all the useful information for the faint sources (which are the main motivation for the development of sophisticated classification techniques)."
" However, theexistingUKIDS Sdatapreductsinehudadulyhadr probabilities, vidicchoabiondprior knowledge, atantthebaeperigorous icclassifications, andtheapplicationofthe Bayesianmethoddesilap aural cankhoss ihbeld qairfactabiive wsdibilédtopS Sinformation consistentamannerasispossiblewithoutusingcolourin formationaU"," However, the existing UKIDSS data products include only heuristic combinations of the band-specific classifications, and the application of the Bayesian method developed here makes it possible to extract all the useful UKIDSS information on a source's morphology in as self-consistent a manner as is possible without using colour information as well."
"NIdéhecSSdette ao ÉpiGagindpanmuadiiandueotat beaver identclassi ficationof faintsources, forwhichtheavailablemegigeeremendsehiysw wmthokonmpprokeastal qgiccitnlorumetiemo"," In particular, the use of prior information avoids the overly-confident classification of faint sources, for which the available measurements contain little morphological information."
"f Our test sample of UKIDSS LAS sources was chosen to lie in the multiply-scanned SDSS Stripe 82 region, giving us independent and almost totally reliable classifications of all our sources. ("," Our test sample of UKIDSS LAS sources was chosen to lie in the multiply-scanned SDSS Stripe 82 region, giving us independent and almost totally reliable classifications of all our sources. ("
"This is a rare situation outside simulations, and an opportunity that could be used for a number of similar testing schemes.)","This is a rare situation outside simulations, and an opportunity that could be used for a number of similar testing schemes.)"
" Converting the posterior probabilities intoclass labels,the Bayesian classifier achieves an error rate of 0.068 "," Converting the posterior probabilities intoclass labels,the Bayesian classifier achieves an error rate of $0.068$ "
(Hotanetal.2004) and corrected for parallactic angle and the orientation of the feed.,"\citep{Hot04}
 and corrected for parallactic angle and the orientation of the feed."
 The position angles are also corrected for rotation through the interstellar medium using the nominal Faradayrotation measure., The position angles are also corrected for Faraday rotation through the interstellar medium using the nominal rotation measure.
" The area of sky containing PSRJJI622-4950 was covered previously by the Parkes Multibeam Pulsar Survey (Manchesteretal.2001) but, somewhat surprisingly given its large apparent flux density and DM, the pulsar was not detected in these data."," The area of sky containing J1622–4950 was covered previously by the Parkes Multibeam Pulsar Survey \citep{Man01}
 but, somewhat surprisingly given its large apparent flux density and DM, the pulsar was not detected in these data."
" However, that survey did find two other radio pulsars in close proximity to PSRJJ1622—4950 and monitored them for a number of years, namely PSRs J1623-4949 and J1622—4944 at an angular separation of 11’ and 7' respectively."," However, that survey did find two other radio pulsars in close proximity to J1622–4950 and monitored them for a number of years, namely PSRs J1623--4949 and J1622–4944 at an angular separation of $'$ and $'$ respectively."
 These archival data were reprocessed using the period and DM of JJ1622—4950., These archival data were reprocessed using the period and DM of J1622–4950.
" The first detection was made in observations recorded in June 1999 (MJD = 51334), and it was then detected in 11 observations between October 2000 (MJD = 51844) and July 2002 (MJD = 52458) (see the of 1))."," The first detection was made in observations recorded in June 1999 (MJD = 51334), and it was then detected in 11 observations between October 2000 (MJD = 51844) and July 2002 (MJD = 52458) (see the top panel of Fig. \ref{Fig:Lightcurve}) )."
" However, in the 14 subsequent observationstop panelpreceding Fig.August 2006 (MJD = 53975), the pulsar was detected only once."," However, in the 14 subsequent observations preceding August 2006 (MJD = 53975), the pulsar was detected only once."
 No further archival data of the two nearby pulsars were available after that date., No further archival data of the two nearby pulsars were available after that date.
" The Methanol Multibeam pulsar survey at 6.3GGHz (Batesetal.inprep.) also covered the part of the sky containing JJ1622—4950, and by reprocessing the relevant data, the pulsar was detected in an observation taken in April 2007 (MJD = 54211)."," The Methanol Multibeam pulsar survey at GHz \citep{Bat10}
 also covered the part of the sky containing J1622–4950, and by reprocessing the relevant data, the pulsar was detected in an observation taken in April 2007 (MJD = 54211)."
" Observations of the pulsar were made on 8 Dec 2009 and 27 Feb 2010 with the Australia Telescope Compact Array (ATCA), an east-west synthesis telescope located near Narrabri, NSW, which consists of six 22-m antennas on a 6-km track."," Observations of the pulsar were made on 8 Dec 2009 and 27 Feb 2010 with the Australia Telescope Compact Array (ATCA), an east-west synthesis telescope located near Narrabri, NSW, which consists of six 22-m antennas on a 6-km track."
" The observations were carried out at 5.5 and 9.0 GHz with a bandwidth of 2simultaneously GHz at each frequency subdivided into 2048 spectral channels, and full Stokes parameters."," The observations were carried out simultaneously at 5.5 and 9.0 GHz with a bandwidth of 2 GHz at each frequency subdivided into 2048 spectral channels, and full Stokes parameters."
 The source was tracked for 12-hours in each of the EW352 and 750B array , The source was tracked for 12-hours in each of the EW352 and 750B array configurations.
Initial data reduction and analysis configurations.were carried out with the MIRIAD using standard techniques., Initial data reduction and analysis were carried out with the MIRIAD using standard techniques.
" After flagging bad data, the primary calibrator (PKS 1934-638) was used for flux density and bandpass calibration and the secondary calibrator (PKS J1613-586) was used to solve for antenna gains, phases and polarization leakage terms."," After flagging bad data, the primary calibrator (PKS 1934-638) was used for flux density and bandpass calibration and the secondary calibrator (PKS J1613-586) was used to solve for antenna gains, phases and polarization leakage terms."
 We also obtained a 20-ks observation performed with ACIS-I on 10 July 2009 to search for an X-ray counterpart., We also obtained a 20-ks observation performed with ACIS-I on 10 July 2009 to search for an X-ray counterpart.
 Figure 2 shows the resultant radio and X-ray images., Figure \ref{Fig:ATCA-CXO} shows the resultant radio and X-ray images.
" Although there are a number of radio and X-ray sources present in the field of view, a highly polarized radio point source is coincident with the brightest X-ray source, CXOU J162244.8—495054, which is also the counterpart of a source seen in archival ASCA data, AX J162246-4946."," Although there are a number of radio and X-ray sources present in the field of view, a highly polarized radio point source is coincident with the brightest X-ray source, CXOU J162244.8–495054, which is also the counterpart of a source seen in archival ASCA data, AX J162246–4946."
 We note that CXOU J162244.8-495054 was one of only two X-ray sources in the field without an optical/infrared counterpart., We note that CXOU J162244.8–495054 was one of only two X-ray sources in the field without an optical/infrared counterpart.
 We are therefore confident that the polarized source seen in the radio is indeed the pulsar and that it has an X-ray counterpart., We are therefore confident that the polarized source seen in the radio is indeed the pulsar and that it has an X-ray counterpart.
also utilize the Galaxy Zoo 1 survey results tottetal. 2011).,"also utilize the Galaxy Zoo 1 survey results \citep[GZ1,][]{Lintott2011}."
". The Galaxy Zoo (Lintottetal.2008) contains a total of 900,000 SDSS galaxies with morphological classifications (Lintottetal.2011)."," The Galaxy Zoo \citep{Lintott2008}
 contains a total of 900,000 SDSS galaxies with morphological classifications \citep{Lintott2011}."
". While this study does not focus exclusively on the LRG sample, it should be noted that if it is possible to improve the Photo-Z estimates for these objects as shown herein it could also improve the estimation of cosmological parameters (e.g.Blake&Bridle2005;Pad-2010;Zunckel,Gott&Lunnan using the SDSS as well as upcoming surveys such 2011)as citepCuesta2011,Eisenstein2011,, BigBOSS (Schlegeletal. 2009),, and possibly Euclid (Sorba&Sawicki2011),, not to mention citepIvezic08.."," While this study does not focus exclusively on the LRG sample, it should be noted that if it is possible to improve the Photo-Z estimates for these objects as shown herein it could also improve the estimation of cosmological parameters \citep[e.g.][]{BB2005,Padmanabhan2007,Percival2010,Reid2010,Zunckel2011} using the SDSS as well as upcoming surveys such as \\citep{Cuesta2011,Eisenstein2011}, BigBOSS \citep{Schlegel2009}, and possibly Euclid \citep{Sorba2011}, not to mention \\citep{Ivezic08}."
" It could also contribute to more reliable Photo-Z errors, as required for weak-lensing surveys (Bernstein&Huterer2010;Kitching,HeavensMiller and Baryonic Acoustic Oscillation measurements, which are also dependent upon accurate Photo-Z estimation of LRGs (Roigetal.2008)."," It could also contribute to more reliable Photo-Z errors, as required for weak-lensing surveys \citep{Bernstein2010,Kitching2011} and Baryonic Acoustic Oscillation measurements, which are also dependent upon accurate Photo-Z estimation of LRGs \citep{Roig2008}."
. All of the data used herein have been obtained via the SDSS casjobsserver®., All of the data used herein have been obtained via the SDSS casjobs.
. In order to obtain results consistent with Paper II for both the MGS and LRG samples we use the same photometric quality flags (I, In order to obtain results consistent with Paper II for both the MGS and LRG samples we use the same photometric quality flags (!
BRIGHT and !,BRIGHT and !
BLENDED and !,BLENDED and !
SATURATED) and redshift quality and but using the SDSS DR7 instead of (zConf>0.95earlier SDSS zWarning=0)releases.,SATURATED) and redshift quality $>$ 0.95 and zWarning=0) but using the SDSS DR7 instead of earlier SDSS releases.
" These data are cross-matched in casjobs with columns 14--16 in Table 2 of Lintottetal.(2011) extracting the galaxies flagged as ‘spiral’, ‘elliptical’ or ‘uncertain’."," These data are cross-matched in casjobs with columns 14–16 in Table 2 of \cite{Lintott2011} extracting the galaxies flagged as `spiral', `elliptical' or `uncertain'."
" The galaxies “flagged as ‘elliptical’ or ‘spiral’ require 80 per cent of the vote in that category after the debiasing procedure has been applied; all other galaxies are flagged 'uncertain"""" (Lintottetal."," The galaxies “flagged as `elliptical' or `spiral' require 80 per cent of the vote in that category after the debiasing procedure has been applied; all other galaxies are flagged `uncertain'"" \citep{Lintott2011}."
 Debiasing is the processes of correcting for small 2011)..biases in spin direction and color., Debiasing is the processes of correcting for small biases in spin direction and color.
 See Section 3.1 in Lintottetal.(2011) for more details on debiasing., See Section 3.1 in \cite{Lintott2011} for more details on debiasing.
" Note that the GZ1 sample is based upon the MGS, but the MGS contains LRGs as well."," Note that the GZ1 sample is based upon the MGS, but the MGS contains LRGs as well."
 This is why we can analyze both of these samples., This is why we can analyze both of these samples.
" However, the actual LRG survey goes fainter than the MGS and so we do not find LRG galaxies fainter than the MGS limit of S17.77."," However, the actual LRG survey goes fainter than the MGS and so we do not find LRG galaxies fainter than the MGS limit of $_{petrosian}\lesssim$ 17.77."
 See Straussetal.(2002) and EisensteinYpetrosianetal.(2001) for details on the MGS and LRG samples., See \cite{Strauss02} and \cite{Eisenstein01} for details on the MGS and LRG samples.
 A number of points from both the LRG and MGS were eliminated because of either bad values (e.g. -9999) or because they were considered outliers from the main distribution of points., A number of points from both the LRG and MGS were eliminated because of either bad values (e.g. -9999) or because they were considered outliers from the main distribution of points.
" The former offenders included: petroR90.ii (13 points in the MGS sample, 1 point in the LRG), mE1-ii (43 points, 5 points), petroR90Err.ii points, 1262 mRrCcErr.ii (22 points, 12 points)."," The former offenders included: i (13 points in the MGS sample, 1 point in the LRG), i (43 points, 5 points), i (7177 points, 1262 points), i (22 points, 12 points)."
"(7177 The reason for points),eliminating bad mE1_ii points is that we use it for calculating aE.ii from Table 2 of Banerjietal.(2010).", The reason for eliminating bad i points is that we use it for calculating i from Table 2 of \cite{Banerji2010}.
". A small number of outliers were also removed from the MGS sample, but totalled only 27 points."," A small number of outliers were also removed from the MGS sample, but totalled only 27 points."
 No such outlier points were removed in the LRG sample., No such outlier points were removed in the LRG sample.
" This leaves us with a total of 437,273 MGS and 68,996 LRG objects."," This leaves us with a total of 437,273 MGS and 68,996 LRG objects."
" Using the GZ1 classifications in the MGS there are 45,249 ellipticals, 119,369 spirals and 272,655 uncertain (~ 62%))."," Using the GZ1 classifications in the MGS there are 45,249 ellipticals, 119,369 spirals and 272,655 uncertain $\sim$ )."
" For the LRG sample there are 27,227 ellipticals and 13,495 spirals leaving 28,274 uncertain (~41%))."," For the LRG sample there are 27,227 ellipticals and 13,495 spirals leaving 28,274 uncertain $\sim$ )."
 Using the morphological classifications from the Galaxy Zoo project first data release etal.2011) we attempt to calculate Photo-Zs (Lintottfor 4 different samples and four combinations of primary and secondary isophotal shape estimates from the SDSS as seen in Table 1.., Using the morphological classifications from the Galaxy Zoo project first data release \citep{Lintott2011} we attempt to calculate Photo-Zs for 4 different samples and four combinations of primary and secondary isophotal shape estimates from the SDSS as seen in Table \ref{tbl-1}.
 A larger variety of input combinations were tried including those in Table 1 of Banerjietal.(2010)., A larger variety of input combinations were tried including those in Table 1 of \cite{Banerji2010}.
". However, we only report those found with the lowest root mean square error (rmse) in Table 1 of this paper."," However, we only report those found with the lowest root mean square error (rmse) in Table \ref{tbl-1} of this paper."
 The results using the Banerjietal.(2010) suggested isophotal shape estimates as well as others tested in, The results using the \cite{Banerji2010} suggested isophotal shape estimates as well as others tested in
respectively.,respectively.
" This indicates that the median (absolute value of Ay at which we find half of the total mass of the histogram) of the Stokes V amplitudes has increased by just a factor of 1.27 when increasing the spatial resolution from 1.92"" to 0.32"".", This indicates that the median (absolute value of $A_V$ at which we find half of the total mass of the histogram) of the Stokes $V$ amplitudes has increased by just a factor of 1.27 when increasing the spatial resolution from $''$ to $''$.
" Regarding the linear polarization histograms, the right panel of Fig."," Regarding the linear polarization histograms, the right panel of Fig."
" 2. shows that the mean values are 5.2x1077. and 4.1x1071. for the 0.32"" and the low resolution data. respectively."," \ref{hist_amplitudes_varia_res} shows that the mean values are $\times 10^{-4} I_\mathrm{c}$ and $\times 10^{-4} I_\mathrm{c}$ for the $''$ and the low resolution data, respectively."
" Accordingly, the median linear polarization signal increases by a factor of 1.27 when increasing the spatial resolution from 1.92"" to 0.32""."," Accordingly, the median linear polarization signal increases by a factor of 1.27 when increasing the spatial resolution from $''$ to $''$."
 The previous values for the median of the circular and linear polarization amplitudes indicate that the quiet Sun magnetism does not change much when increasing the spatial resolution from a few aresee to sub-aresec scales., The previous values for the median of the circular and linear polarization amplitudes indicate that the quiet Sun magnetism does not change much when increasing the spatial resolution from a few arcsec to sub-arcsec scales.
 We now focus on the whole shape of the histograms., We now focus on the whole shape of the histograms.
" First, the circular polarization amplitude histograms do not present a Gaussian shape throughout the full range of amplitudes, manifesting signatures of intermittency (c.g.,Stenflo&Holzreuter2003:AsensioRamos 2009b)."," First, the circular polarization amplitude histograms do not present a Gaussian shape throughout the full range of amplitudes, manifesting signatures of intermittency \citep[e.g.,][]{stenflo03,andres09}."
". Additionally, MartinezGonzálezetal.(2008a) also found that these extended tails are modified when observing at different heliocentric angles, an indication that these tails correspond to non-isotropic structures."," Additionally, \cite{marian_andres_08} also found that these extended tails are modified when observing at different heliocentric angles, an indication that these tails correspond to non-isotropic structures."
" However, the core of the histograms can be well reproduced by a Gaussian function."," However, the core of the histograms can be well reproduced by a Gaussian function."
" The standard deviation found for the different spatial resolutions considered are: 63»=1.0x107. 0,54=8.6x1071. and Cis)=C1928.4x1071."," The standard deviation found for the different spatial resolutions considered are: $\sigma_{0.32}=1.0 \times 10^{-3} I_\mathrm{c}$, $\sigma_{1.28}=8.6\times 10^{-4} I_\mathrm{c}$, and $\sigma_{1.60}=\sigma_{1.92}=8.4 \times 10^{-4} I_\mathrm{c}$."
" This demonstrates that the core of the histograms at different spatial resolutions, i. e.. for amplitudes between +2.0x1077. are Gaussian and almost indistinguishable."," This demonstrates that the core of the histograms at different spatial resolutions, i. e., for amplitudes between $\pm 2.0 \times 10^{-3} I_\mathrm{c}$ are Gaussian and almost indistinguishable."
" The percentage of the observed area in the Sun covered by pixels with amplitudes in the core isπα.82%,, and for the 0.32"", 1.28-1.60"" and 1.92"" data sets, respectively."," The percentage of the observed area in the Sun covered by pixels with amplitudes in the core is, and for the $''$, $''$ and $''$ data sets, respectively."
 A deviation from the Gaussian shape occurs for pixels with larger circular polarization amplitudes and they do display a trend with the spatial resolution., A deviation from the Gaussian shape occurs for pixels with larger circular polarization amplitudes and they do display a trend with the spatial resolution.
" Therefore, we can state that ~20% of the pixels show an increase in the magnetic flux density when increasing the spatial resolution, a consequence of the small number of signal cancellations at 0.32”"," Therefore, we can state that $\sim$ of the pixels show an increase in the magnetic flux density when increasing the spatial resolution, a consequence of the small number of signal cancellations at $''$ ."
We now look in more detail at the DNO. and QPO behaviours in cach of the four light curves.,We now look in more detail at the DNO and QPO behaviours in each of the four light curves.
 None of the E'Ts shows any coherent periocdicities at periods shorter than 13 S. We divided the light curve into 10. equal. subsections (of length  600 s). the first six of which contain no evidence for DNOs. after which a clear signal at 22.05 s is seen in the seventh section.," None of the FTs shows any coherent periodicities at periods shorter than 13 s. We divided the light curve into 10 equal subsections (of length $\sim$ 600 s), the first six of which contain no evidence for DNOs, after which a clear signal at 22.05 s is seen in the seventh section."
 The next obvious DNO in the ET occurs in the combined ninth and tenth sections at à period of 24.78 s and also à very strong QPO at 400 s. More details are given in Table 2.., The next obvious DNO in the FT occurs in the combined ninth and tenth sections at a period of 24.78 s and also a very strong QPO at 400 s. More details are given in Table \ref{dno7tab2}.
 We now introduce an vl© diagram. as defined in Paper IV. which gives the Observed —- Calculated amplitude and phase variations relative to a fixed. period.," We now introduce an $A - \phi$ diagram, as defined in Paper IV, which gives the Observed – Calculated amplitude and phase variations relative to a fixed period."
 The 2bo diagram. of the final part (UTC 0.90. — 0.93) of the light curve. relative to the period 24.78 s. is shown in Fig. 4..," The $A - \phi$ diagram of the final part (UTC 0.90 – 0.93) of the light curve, relative to the period 24.78 s, is shown in Fig. \ref{SALT016omc1}."
 Up until 0.918 the DNO has a constant period. after which the »eriod. lengthens. as part of the general increase in. period. normally seen in the decline phase of the light curve.," Up until 0.918 the DNO has a constant period, after which the period lengthens, as part of the general increase in period normally seen in the decline phase of the light curve."
" During his latter part there is strong amplitude modulation (mice ΠΟ}, seen to be approximately in antiphase with the ΟΡΟ in the light curve (upper panel)."," During this latter part there is strong amplitude modulation (middle panel), seen to be approximately in antiphase with the QPO in the light curve (upper panel)."
 This is an indication that we are seeing the svnodic and not the sidereal DNO — the DNO rotating beamed radiation is reprocessed olf the 400 s QPO travelling wave., This is an indication that we are seeing the synodic and not the sidereal DNO – the DNO rotating beamed radiation is reprocessed off the 400 s QPO travelling wave.
 This implies ono = 23.33 s. We divided the run into five subsections., This implies $P_{\rm DNO}$ = 23.33 s. We divided the run into five subsections.
 Phe E'Ts naturally show very strong DNO peaks. their. periods. increasing hrough the run. with a mean period of 25.64 s and mean amplitude 26.6 mimag in the short first section ancl periods al 27.65 s (23.3 mmae) in the second section. increasing o 28.82 s (19.5 mmag) towards the end.," The FTs naturally show very strong DNO peaks, their periods increasing through the run, with a mean period of 25.64 s and mean amplitude 26.6 mmag in the short first section and periods at 27.65 s (23.3 mmag) in the second section, increasing to 28.82 s (19.5 mmag) towards the end."
 There are short ived appearances of IpDNOs. e.g. at 86.6 s and 92.8 s. This general behaviour only serves to add: more weight to the same correlations previously found. (figure 3 of Paper IL and igure 1 of Paper IV).," There are short lived appearances of lpDNOs, e.g. at 86.6 s and 94.8 s. This general behaviour only serves to add more weight to the same correlations previously found (figure 3 of Paper I and figure 1 of Paper IV)."
 Fig., Fig.
 5r shows the 46 diagram for the second. part of he light curve (subsections 2 to 5). here split into only two sections. which shows a sudden change of behaviour roughly midway through. where the mean DNO period (adopted for calculating ©) changes from 27.65 s to 28.50 s. The first section shows evclical phase variations (Fig. 5)))).," \ref{SALT004omc1} shows the $A - \phi$ diagram for the second part of the light curve (subsections 2 to 5), here split into only two sections, which shows a sudden change of behaviour roughly midway through, where the mean DNO period (adopted for calculating $\phi$ ) changes from 27.65 s to 28.50 s. The first section shows cyclical phase variations (Fig. \ref{SALT004omc1}) ),"
 an FV of which gives a period of 475 c 13 s. but an ET of the same piece of light curve itself shows no significant QPO period demonstrating (as occasionally seen. before) that although the rotating DNO beam intersects the OQPO travelling wave. the latter is not necessarily visible in the direction of the observer.," an FT of which gives a period of 475 $\pm$ 13 s, but an FT of the same piece of light curve itself shows no significant QPO period demonstrating (as occasionally seen before) that although the rotating DNO beam intersects the QPO travelling wave, the latter is not necessarily visible in the direction of the observer."
 The FP (Pig. 6)), The FT (Fig. \ref{SALT004ft1}) )
 of the short first part of the run shows a peak at 13.39 s. which has not been previously seen in any of the extensive studies of VW. Livi.," of the short first part of the run shows a peak at 13.39 s, which has not been previously seen in any of the extensive studies of VW Hyi."
 We see from figure 10 of Paper IV that the expected QPO fundamental at 7 = 1.19 0.20 is  400) 550s with a first harmonic probably appearing curing this at  300J) s. Ehe evelic phase variation seen in the short first part of the run corresponds to the QPO 'undamental., We see from figure 10 of Paper IV that the expected QPO fundamental at $T$ = 0.10 – 0.20 d is $\sim$ 400 – 550 s with a first harmonic probably appearing during this at $\sim$ 300 s. The cyclic phase variation seen in the short first part of the run corresponds to the QPO fundamental.
 Phe 13.39 s modulation is not the harmonic of he observed DNO (the sidereal period at 25.64 8). but iit is he first harmonic of the svnodie period (ic. yy /2) then he beat between sideral and. svnodic periods is at 602s. which is in the neighbourhood of the OPO period expected rom the general correlation with 7.," The 13.39 s modulation is not the harmonic of the observed DNO (the sidereal period at 25.64 s), but if it is the first harmonic of the synodic period (i.e., $P_{\rm SYN}$ /2) then the beat between sideral and synodic periods is at 602s, which is in the neighbourhood of the QPO period expected from the general correlation with $T$ ."
 In the first section of, In the first section of
"nucleus and the star in the GMOS acquisition image, which has been taken before the spectra, we measure à corresponding separation of 2.354-0.08"".","nucleus and the star in the GMOS acquisition image, which has been taken before the spectra, we measure a corresponding separation of $2.35\pm0.08$."
". Using the light distribution of the spectrum along the slit, and fitting two Gaussians to the light peaks, we measure a separation of 2.45+0.1."," Using the light distribution of the spectrum along the slit, and fitting two Gaussians to the light peaks, we measure a separation of $2.45\pm0.1$."
" ? did not report any displacement of the kinematical centre, and therefore, we will not consider the small discrepancy between our photometric and kinematical determinations (one pixel, or one fifth of the FWHM seeing) as significant; it may result from the dust extinction."," \citet{nowak2007} did not report any displacement of the kinematical centre, and therefore, we will not consider the small discrepancy between our photometric and kinematical determinations (one pixel, or one fifth of the FWHM seeing) as significant; it may result from the dust extinction."
 We adopt the kinematical centre as a reference., We adopt the kinematical centre as a reference.
 The analysis was performed with the full-spectrum fitting package (?).., The analysis was performed with the full-spectrum fitting package \citep{ulyss}.
 The present difficulty of contamination by a foreground or background object was already met on other occasions., The present difficulty of contamination by a foreground or background object was already met on other occasions.
" In ? and ?,, the interloper was separately modelled and then subtracted before studying the target."," In \cite{bouchard2010} and \citet{makarova2010}, the interloper was separately modelled and then subtracted before studying the target."
" In the present case, the contamination reaches and extends over a large range of the object."," In the present case, the contamination reaches and extends over a large range of the object."
 We therefore tried to optimise the decomposition., We therefore tried to optimise the decomposition.
 The ULySS package allows one to fit a spectrum against a constrained linear combination of non-linear, The ULySS package allows one to fit a spectrum against a constrained linear combination of non-linear
(Greisen1966:Zalsepin&Ixuzmin1966).,"\citep{g,zk}."
. This delines a GZlx horizon located roughly al a distance LOO—200 Mpe (ILararietal.2006)., This defines a GZK horizon located roughly at a distance $\sim 100-200$ Mpc \citep{harari06}.
. Astrophysical sources located outside the GZlx horizon are not expected to contribute appreciably to the observed Πας of energetic cosmic ravs., Astrophysical sources located outside the GZK horizon are not expected to contribute appreciably to the observed flux of energetic cosmic rays.
 Motivated bv this expectation. we investigate the correlation of PAO events with Ηλ AGNs located within the GZIX horizon.," Motivated by this expectation, we investigate the correlation of PAO events with 1LAC AGNs located within the GZK horizon."
" We define the GZlx horizon as being located at a comoving distance of 200 Alpe. which corresponds (0 z,,,,=0.048. assuming the ACDM cosmology with Ly)=71kans!Mpe|. Q,,=0.270. O4=0.730 and Οι=0 (Ixomatsuetal. 2009).."," We define the GZK horizon as being located at a comoving distance of 200 Mpc, which corresponds to $z_{\rm max}=0.048$, assuming the $\Lambda$ CDM cosmology with $H_0 = 71 \ \textrm{km} \ \textrm{s}^{-1} \ \textrm{Mpc}^{-1}$, $\Omega_m = 0.270$, $\Omega_{\rm \Lambda} = 0.730$ and $\Omega_k=0$ \citep{wmap}. ."
 We use the redshift information provided in the LILAC (Abdoetal.2010b) to discard the AGNs within the PAO field of view for which z>zg4;., We use the redshift information provided in the 1LAC \citep{1lac} to discard the AGNs within the PAO field of view for which $z>z_{\rm max}$.
 Only 9 AGNs remain. including the radio galaxies Centaurs A (the nearest radio galaxv) ancl MST. the Sevlert 2 NGC 4945. the Seviert 1 ESO ai? and (he starburst galaxy NGC 253 (as previously noted bv Dermer&Razzaque 2010)).," Only 9 AGNs remain, including the radio galaxies Centaurus A (the nearest radio galaxy) and M87, the Seyfert 2 NGC 4945, the Seyfert 1 ESO 323-G77 and the starburst galaxy NGC 253 (as previously noted by \citealt{dermer10}) )."
 Onlv three AGNs within the GZIx horizon are BL Lacertae objects., Only three AGNs within the GZK horizon are BL Lacertae objects.
 There is no FSRQ AGN within theGZIx horizon. with the nearest FSRQ blazar in the [LAC being located at 2=0.15.," There is no FSRQ AGN within theGZK horizon, with the nearest FSRQ blazar in the 1LAC being located at $z=0.15$."
 LLAC sources inside the GZIx horizon are shown as squares in Figure 6.., 1LAC sources inside the GZK horizon are shown as squares in Figure \ref{fig:1lacmap}.
 Figure 9 shows the correlation signal between [LAC AGNs with :<0.048 and CHECKS. with ο reaching up to 20°.," Figure \ref{fig:gzk} shows the correlation signal between 1LAC AGNs with $z \leq 0.048$ and UHECRs, with $\psi$ reaching up to $20^\circ$."
 The noteworthy feature here is that (here are three characteristic angular scales where (he probability of (he observed configuration being explained by an isotropic [hix of cosmic rays is minimized: e12 2.37. vy>se0.57 and iwz16.97.," The noteworthy feature here is that there are three characteristic angular scales where the probability of the observed configuration being explained by an isotropic flux of cosmic rays is minimized: $\psi_1 \approx 2.3^\circ$ , $\psi_2 \approx 6.5^\circ$ and $\psi_3 \approx 16.9^\circ$."
" The first local minimum of the probability signal located al c82.3"" corresponds to only 4 associations. with Pc)=6x10? and ο)=ασ."," The first local minimum of the probability signal located at $\psi_1 \approx 2.3^\circ$ corresponds to only 4 associations, with $P(\psi_1)=6 \times 10^{-5}$ and $S(\psi_1)=4\sigma$."
" The second local minimum of the probability signal located al ος226.5"" corresponds to 7 measured correlations whereas 0.9 chance correlations are expected. wilh P(t.)=5-x10 and Ses)=5e."," The second local minimum of the probability signal located at $\psi_2 \approx 6.5^\circ$ corresponds to 7 measured correlations whereas 0.9 chance correlations are expected, with $P(\psi_2)=5 \times 10^{-7}$ and $S(\psi_2)=5\sigma$."
" The absolute minimum of P is characterized bv vyz16.97. N,=1T. AN.=5.0. Ples)=5x105 and S(t4)=5.1o."," The absolute minimum of $P$ is characterized by $\psi_3 \approx 16.9^\circ$, $N_o=17$, $N_c=5.7$, $P(\psi_3)=5 \times 10^{-8}$ and $S(\psi_3)=5.4\sigma$."
" Withinthe distance ey. 4 PAO events correlate with3 AGNs: NGC 4945. Centaurus A and 4C--04.77: ad cs. 7 PAO events correlate with d AGNs: NGC 4945. Centatrus A. ESO 323-CG77 and 4C+04,77: at s. 17 events are associated with 7 AGNs: NGC 4945. Centaurus A. ESO 323-G77. 4C-04.77. NGC 253. NGC 1218 and RXJOO08.02-1450."," Withinthe distance $\psi_1$, 4 PAO events correlate with 3 AGNs: NGC 4945, Centaurus A and 4C+04.77; at $\psi_2$, 7 PAO events correlate with 4 AGNs: NGC 4945, Centaurus A, ESO 323-G77 and 4C+04.77; at $\psi_3$, 17 events are associated with 7 AGNs: NGC 4945, Centaurus A, ESO 323-G77, 4C+04.77, NGC 253, NGC 1218 and RXJ0008.0+1450."
 Table 1. lists the 9 ILAC AGNs located inside the GZlx horizon. including their IFGL names. Galactic coordinates (longitude / and latitude 56). identifications and redshifts according to (2010a.b)..," Table \ref{tab:agn} lists the 9 1LAC AGNs located inside the GZK horizon, including their 1FGL names, Galactic coordinates (longitude $l$ and latitude $b$ ), identifications and redshifts according to \citet{1fgl,1lac}. ."
 We also list the number of PAO events potentiallycorrelated with each AGN al (he angular scales cé» and vy., We also list the number of PAO events potentiallycorrelated with each AGN at the angular scales $\psi_2$ and $\psi_3$ .
 The ellipses in Figure 6 correspond to circles of radius vs centered onthe GZlx. AGNs., The ellipses in Figure \ref{fig:1lacmap} correspond to circles of radius $\psi_3$ centered onthe GZK AGNs.
clear that hourelass-like aud spiral-like. feaures found in residual images of the kind shown in Fie.,clear that hourglass-like and spiral-like features found in residual images of the kind shown in Fig.
" | nist not. af face value. be taken as indication for the presence ofa true bar or true spiral structure νο, disk structure)."," 1 must not, at face value, be taken as indication for the presence of a true bar or true spiral structure (i.e. disk structure)."
 However. as mentioned iu the iutroductiou and as demonstrated in the following section. the residual features in our four dwarfs clearly persist when scrutinized with au ΠΣΑΤ) masking incthod.," However, as mentioned in the introduction and as demonstrated in the following section, the residual features in our four dwarfs clearly persist when scrutinized with an unsharp masking method."
 Their strength makes it difficult to, Their strength makes it difficult to
 — l.l dex).,$\sim$ $-$ 1.4 dex).
 In particular. there is an anticorrelation among aud for the most metal-rich stars.," In particular, there is an anticorrelation among and for the most metal-rich stars."
 This could most readily be explained bv the euriclinent. of CNO-processed material (with iuereased N anc decreased C)., This could most readily be explained by the enrichment of CNO-processed material (with increased N and decreased C).
 However. that oulv 1e niost nctalazich stars seen to show this auticorrelation is puzzling.," However, that only the most metal-rich stars seem to show this anticorrelation is puzzling."
 A careful examination of the location of these stars in the colouranagnitude diagram shows that the nost mctalrich stars are not the reddest ones;, A careful examination of the location of these stars in the colour-magnitude diagram shows that the most metal-rich stars are not the reddest ones.
 Moreover. jese stars have rather intermediate luuinosities (as compared to the overall sample of stars) at these colours.," Moreover, these stars have rather intermediate luminosities (as compared to the overall sample of stars) at these colours."
 This indicates that the observed anticorrelatiou is not due to (not corrected for) temperature and/or Inuuinosity effects or tha these stars occupy a special region iui the colom-iuaguitudeo diagram., This indicates that the observed anticorrelation is not due to (not corrected for) temperature and/or luminosity effects or that these stars occupy a special region in the colour-magnitude diagram.
 Tusteack. it could be that the eurichmeut history for these stars was indeed very different from hat of the earlier populations.," Instead, it could be that the enrichment history for these stars was indeed very different from that of the earlier populations."
 The scatter as ineasured by the rus of the best-fittine liue is much larger for aud tthan for the a clemeuts Me aud Ca (as measured by aad (C22)., The scatter as measured by the rms of the best-fitting line is much larger for and than for the $\alpha$ elements Mg and Ca (as measured by and ).
 Under the asstunptiou that nunostlv depends on N while lis a nmieasure of C abundance. we conclude that there is indeed a wide rauge of C aud N abuudances. similar to what was found for ROB stars by ?..," Under the assumption that mostly depends on N while is a measure of C abundance, we conclude that there is indeed a wide range of C and N abundances, similar to what was found for RGB stars by \citet{N/DC:95}."
 It should be also noted that the average index (653839|GCIT1300)/2 (= iu Fie. 16)), It should be also noted that the average index $(\delta S3839+\delta CH4300)/2$ $=$ in Fig. \ref{cn3}) )
 shows a strong correlation with increasing aand a siguificautly simaller scatter as compared to the individual indices and1300., shows a strong correlation with increasing and a significantly smaller scatter as compared to the individual indices and.
 This was also observed by ?.. albeit for RGB stars.," This was also observed by \citet{H/R:00}, albeit for RGB stars."
 The fact that the CN/CIT anticorrelation is not only visible ou the RGB but is also detected for almost nnevolved stars on the SGB strongly sugeests that mixiug effects cannot be the dominant cause of the increase of C and N for higher iietallicities., The fact that the CN/CH anticorrelation is not only visible on the RGB but is also detected for almost unevolved stars on the SGB strongly suggests that mixing effects cannot be the dominant cause of the increase of C and N for higher metallicities.
 Instead. this shows that the," Instead, this shows that the"
shooting method. adopting a rigid inner boundary condition at the stellar surface and seeking prograde confined. modes that decay asaox. (,"shooting method, adopting a rigid inner boundary condition at the stellar surface and seeking prograde confined modes that decay as $x\to\infty$. ("
In practice this is done by imposing a rigid outer boundary condition and verifving that the eigenvalue is completely insensitive to the value of the outer racdus provided it is sullicientIy large.),In practice this is done by imposing a rigid outer boundary condition and verifying that the eigenvalue is completely insensitive to the value of the outer radius provided it is sufficiently large.)
 The results are shown in the left panel of Fig. 3.., The results are shown in the left panel of Fig. \ref{f:ot}.
 Here we confirm that. confined prograde modes exist in the 2D theories only for sullicientLy large values of Q. as described in Section 3-4..," Here we confirm that confined prograde modes exist in the 2D theories only for sufficiently large values of $\tilde Q$, as described in Section \ref{s:qcrit}."
 In contrast. the 3D theory allows such moces to be obtained for any value of Q.," In contrast, the 3D theory allows such modes to be obtained for any value of $\tilde
Q$."
 Furthermore. the precession rates are much larger in the 3D theory.," Furthermore, the precession rates are much larger in the 3D theory."
 Since the realistic values of Q are probably in the vicinity of 10 or smaller. it is clear that the three-dimensional elfects are of essential importance in the case of a rigid inner boundary condition.," Since the realistic values of $\tilde Q$ are probably in the vicinity of $10$ or smaller, it is clear that the three-dimensional effects are of essential importance in the case of a rigid inner boundary condition."
 In the right panel of Fig., In the right panel of Fig.
" 3 we show comparable results for the free inner boundary condition considered. bv '""apaloizou&Savonije(2006).. meaning that the Lagrangian oressure perturbation Is zero ator=A."," \ref{f:ot} we show comparable results for the free inner boundary condition considered by \citet{2006A&A...456.1097P}, meaning that the Lagrangian pressure perturbation is zero at $r=R$."
" In the secular theory his is equivalent to 0,47=0. or à|2r0,i0."," In the secular theory this is equivalent to $\p_rE=0$, or $u+2r\p_r u=0$."
 As described » Papaloizou&Savonije(2006).. à free inner boundary condition allows prograde modes to be found. for smaller values of Q in the 2DS theory.," As described by \citet{2006A&A...456.1097P}, a free inner boundary condition allows prograde modes to be found for smaller values of $\tilde Q$ in the 2DS theory."
 Lt is still (ruc. however. that he threc-dimensional elfects have an important. cllect on he results.," It is still true, however, that the three-dimensional effects have an important effect on the results."
 Since the eigenfunctions obtained with a free inner boundary condition are generally. peaked at the inner radius. there may also be a conflict between the nonlinear development of such an eccentric mode and the existence of a stellar surface. unless there is a wide gap between the star and the disc.," Since the eigenfunctions obtained with a free inner boundary condition are generally peaked at the inner radius, there may also be a conflict between the nonlinear development of such an eccentric mode and the existence of a stellar surface, unless there is a wide gap between the star and the disc."
 To convert these eigenvalues into physical units. we again make use of the stellar models in Pitt&Larmanec (2006).," To convert these eigenvalues into physical units, we again make use of the stellar models in \citet{2006A&A...447..277F}."
. We then find that the precession period is where Cp ranges from approximately 1.6 at the lower end to approximately 2.3 at the upper end (or 2.4 for ~ Cas and Ls for 59 Cye)., We then find that the precession period is where $C_P$ ranges from approximately $1.6$ at the lower end to approximately $2.3$ at the upper end (or $2.4$ for $\gamma$ Cas and $1.8$ for $59$ Cyg).
 This conversion factor is independent. of assumptions regarding the stellar rotation rate. except inasmuch as the rotation alfects the equatorial radius. but Cp is inversely proportional to the assumed disc Lemperatiure.," This conversion factor is independent of assumptions regarding the stellar rotation rate, except inasmuch as the rotation affects the equatorial radius, but $C_P$ is inversely proportional to the assumed disc temperature."
 The shapes of the confined modes in the 3D theory are illastratecd in Fig. 4.., The shapes of the confined modes in the 3D theory are illustrated in Fig. \ref{f:ef}.
 For larger values of Q. the mocde is increasingly confined in the inner part of the disc.," For larger values of $\tilde Q$, the mode is increasingly confined in the inner part of the disc."
 Although higher-order modes may exist. one of which is referred. to in Fig. 3..," Although higher-order modes may exist, one of which is referred to in Fig. \ref{f:ot},"
 we focus here on the fundamental confined mode with the simplest racial structure., we focus here on the fundamental confined mode with the simplest radial structure.
 Since wis close to 1 in the 3D theory for reasonable values of Q and with a rigid inner boundary condition. precession rates in the vicinity of 1 νι are. obtained.," Since $\tilde\omega$ is close to $1$ in the 3D theory for reasonable values of $\tilde Q$ and with a rigid inner boundary condition, precession rates in the vicinity of $1$ $3\,\mathrm{yr}$ are obtained."
 Observed. evele times. which can vary significantlv even for the same star. are more usually in the range 5 10vr (Okazaki1997).," Observed cycle times, which can vary significantly even for the same star, are more usually in the range $5$ $10\,\mathrm{yr}$ \citep{1997A&A...318..548O}."
 This discrepancy might occur. because the dises have cilferent temperature or density profiles [rom those assumed here. either. of which would alfect the precession period.," This discrepancy might occur because the discs have different temperature or density profiles from those assumed here, either of which would affect the precession period."
 Non-lincarity of the eccentric mode niw also increase the precession period by altering. the distribution of eccentricity so that it is less peaked in the inner part of the disc., Non-linearity of the eccentric mode may also increase the precession period by altering the distribution of eccentricity so that it is less peaked in the inner part of the disc.
 Furthermore. non-isothermal elfects or radiative forces might be important.," Furthermore, non-isothermal effects or radiative forces might be important."
 In some cases. including ~ Cas and 59 Cve. the presence of a relative close binary companion may also allect the dynamics of the eccentric mode. although it would. seem. most likely to contribute to the progracde precession and. therefore to decrease. the precession period.," In some cases, including $\gamma$ Cas and $59$ Cyg, the presence of a relative close binary companion may also affect the dynamics of the eccentric mode, although it would seem most likely to contribute to the prograde precession and therefore to decrease the precession period."
a significance of >3.20 in the standard Gaussian random. statistically isotropic cosinoloey.,"a significance of $\geq 3.2 \sigma$ in the standard Gaussian random, statistically isotropic cosmology."
 This result. first observed in CODE data. has been strengthened with WALAP first-vear data (Ilinshawetal.1996)... as well as laterWALAP observations (Copietal.etal. 2011).. though questions have been raised regarding its sienilicance 2010:Efstathiouetal.Ma2011) In this brief report. we explore the directional contributions to the angular (wo-point correlation function.," This result, first observed in COBE data, has been strengthened with WMAP first-year data \citep{DMR4}, as well as laterWMAP observations \citep{wmap123,Copi_nolarge,Sarkar}, though questions have been raised regarding its significance \citep{PeirisPontzen,Efstathiou2010,Ma2011}
 In this brief report, we explore the directional contributions to the angular two-point correlation function."
 Directional information is lost in the original definition of C'(8). as Eq. (1))," Directional information is lost in the original definition of $C(\theta)$, as Eq. \ref{eq:Ctheta}) )"
 assumes statistical isotropy ancl averages over all directions on the skv., assumes statistical isotropy and averages over all directions on the sky.
 Our goal is to provide a generalization of Eq. (1)).," Our goal is to provide a generalization of Eq. \ref{eq:Ctheta}) ),"
 and ascertain if there are specific directions which result in the unusually low correlation., and ascertain if there are specific directions which result in the unusually low correlation.
 The angular two-point correlation Iunction is defined by Eq. (1)).," The angular two-point correlation function is defined by Eq. \ref{eq:Ctheta}) ),"
 where the average is taken over all pairs of points separated by angle 0 0n the sky., where the average is taken over all pairs of points separated by angle $\theta$ on the sky.
 We now propose a new statistical quantity. a fixed-vertex correlation Iunction where © is a direction in the sky. and the averageis taken over a ring of pixels separated bv 0 from. the direcon. $9.," We now propose a new statistical quantity, a fixed-vertex correlation function where $\hat\Omega$ is a direction in the sky, and the averageis taken over a ring of pixels separated by $\theta$ from the direction $\hat\Omega$."
 In. particular.. notice. that. meC(0)py=ολ.Aup.(Q;.0)!.," In particular, notice that $C(\theta)=\frac{1}{N}\sum_{i=1}^N C(\hat\Omega_i,
\theta)$."
. Maps οἱ. C(O.90) [or 6 = 60° and 8 = 120° are shown in Figure 1.," Maps of $C(\hat{\Omega},\theta)$ for $\theta$ = $60^\circ$ and $\theta$ = $120^\circ$ are shown in Figure 1."
 In order to quantily the low power in (he angular two-point Iuiction between approximately 60 and 1307. the statistic (Spergeletal.2003) has been widely used.," In order to quantify the low power in the angular two-point function between approximately $60\degr$ and $180\degr$ , the statistic \citep{Spergel2003}
 has been widely used."
 In order (to examine possible in C(0). we attempt to," In order to examine possible in $C(\theta)$ , we attempt to"
emphasizing that the 2.10 keV fIux obtained by ascribing all the far-IR to the SB component and applving the [ar-I to X-ray. correlation of Ranalli. Comastr Seli (2003) lies. in. the range ~1 6.x10|! erg t >7. which. is. on average ~300 . lower than what envisaged for the AGN (see Table 2)).,"emphasizing that the 2–10 keV flux obtained by ascribing all the far-IR to the SB component and applying the far-IR to X-ray correlation of Ranalli, Comastri Setti (2003) lies in the range $\sim$ $6 \times 10^{-14}$ erg $^{-1}$ $^{-2}$, which is on average $\sim$ 300 lower than what envisaged for the AGN (see Table \ref{t2}) )."
 The median 210 keV luminosity precdieted for the AGN component is Ly~2.5xLO! erg |., The median 2–10 keV luminosity predicted for the AGN component is $L_\rmn{X} \sim 2.5 \times 10^{44}$ erg $^{-1}$.
 This is actually a conservative evaluation: if we adopt the direct 6 jan to 2LO keV correction found by Lutz et al. (, This is actually a conservative evaluation: if we adopt the direct 6 $\mu$ m to 2–10 keV correction found by Lutz et al. (
2004a). it should be revised upwards by a [actor of 72.,"2004a), it should be revised upwards by a factor of $\sim$ 2."
 This places our sources well above (he X-ray intrinsic uminositv of the local population of Compton-thick AGN recently identified by Goulding el al. (, This places our sources well above the X-ray intrinsic luminosity of the local population of Compton-thick AGN recently identified by Goulding et al. (
2011). allowing us to explore the parameter space of highly obscured. AGN in the ow-recshift/hieh-Iuminosity region.,"2011), allowing us to explore the parameter space of highly obscured AGN in the low-redshift/high-luminosity region."
 Concerning the unpublished. X-ray data. IRAS 01003—2238 was observed by on 2009 December 08 for 35 ks.," Concerning the unpublished X-ray data, IRAS $-$ 2238 was observed by on 2009 December 08 for 35 ks."
 The EPIC pn and MOS cameras operated in [ull frame mode., The EPIC pn and MOS cameras operated in full frame mode.
 Alter the filtering of high-background periods. the useful pn exposure is ~28 ks.," After the filtering of high-background periods, the useful pn exposure is $\sim$ 28 ks."
 Incicentally. a previous snapshot (+10 ks) taken with in 2003 vielded only 20 counts al 0.58.0 keV. preventing a detailed spectral modelling (Teng et al.," Incidentally, a previous snapshot $\sim$ 10 ks) taken with in 2003 yielded only 20 counts at 0.5–8.0 keV, preventing a detailed spectral modelling (Teng et al."
 2005)., 2005).
 The other five sources. instead. were observed with the ACIS-S detector.," The other five sources, instead, were observed with the ACIS-S detector."
 The basic details about all the observations are summarized in Table 1.., The basic details about all the observations are summarized in Table \ref{t1}.
 The data reduction was performed following the standard procedures. using the latest versions of the and packages lor aud event files. respectively.," The data reduction was performed following the standard procedures, using the latest versions of the and packages for and event files, respectively."
 The source and background spectra were extracted from circular regions with radius of Chandra)) aud XAMAI-Neicton))., The source and background spectra were extracted from circular regions with radius of ) and ).
 In each case the X-ray. emission from the source is entirely collected within the chosen aperture. corresponding to ~310 kpe in the images.," In each case the X-ray emission from the source is entirely collected within the chosen aperture, corresponding to $\sim$ 3–10 kpc in the images."
 No clear morphological structure can be appreciated., No clear morphological structure can be appreciated.
 The analvsis has been performed on unbinned spectra using the v12.5 filling package. and C-statistic has been adopted due to the low number of counts.," The analysis has been performed on unbinned spectra using the v12.5 fitting package, and $C$ -statistic has been adopted due to the low number of counts."
 Quoted uncertainties aid upper limits are given al the 90 per cent confidence level., Quoted uncertainties and upper limits are given at the 90 per cent confidence level.
The fundamental questions of physics appear on (wo levels. (he microscopic and macroscopic.,"The fundamental questions of physics appear on two levels, the microscopic and macroscopic."
 We begin by asking: Given the basic building blocks and their interactions. we want to know: llow lar have we advanced today in our uuderstaiding of these different aspects?," We begin by asking: Given the basic building blocks and their interactions, we want to know: How far have we advanced today in our understanding of these different aspects?"
 According to our present state of knowledge. the ultimate constituents are quarks. leptons. eluons. photons. intermediate vector bosous (Z/W7) and Higgs bosons - in a conservative count (no antiparticles etc.).," According to our present state of knowledge, the ultimate constituents are quarks, leptons, gluons, photons, intermediate vector bosons $Z/W^{\pm}$ ) and Higgs bosons - in a conservative count (no antiparticles etc.),"
 sixteen in all. wilh gravitation not vel in the eame.," sixteen in all, with gravitation not yet in the game."
 Their interactions were orignally classified as strong. electromagnetic. weak and gravitation. leaving a more general scheme as a challenge.," Their interactions were orignally classified as strong, electromagnetic, weak and gravitation, leaving a more general scheme as a challenge."
 The first unification brought electroweak theory. the second combined this with strong interactions to the standard model.," The first unification brought electroweak theory, the second combined this with strong interactions to the standard model."
 The origin of all the different. basic constituents. as well as the role of gravitation. are still open. wailing for the theory of evervthing (TOE).," The origin of all the different basic constituents, as well as the role of gravitation, are still open, waiting for the theory of everything (TOE)."
 In ancient times. the basic states of matter were earth. water. air and fire: today. we have solids. liquids. gases ancl plasmas.," In ancient times, the basic states of matter were earth, water, air and fire; today we have solids, liquids, gases and plasmas."
 In addition. (here now is a multitude of others: insulators. conductors and superconductors. fluids and superfluids. ferromagnets. spin elasses. gelatines and many more.," In addition, there now is a multitude of others: insulators, conductors and superconductors, fluids and superfluids, ferromagnets, spin glasses, gelatines and many more."
 And the question of the possible states of matter brings us (o a new kind of physics; the knowledge of the elementary constituents and (heir inleractions in general does not predict the structure of the possible complex states of many constituents., And the question of the possible states of matter brings us to a new kind of physics; the knowledge of the elementary constituents and their interactions in general does not predict the structure of the possible complex states of many constituents.
 The study of complex systenis becomes even more general. less dependent on (he microstructure. when we ask for the transitions between (he different states.," The study of complex systems becomes even more general, less dependent on the microstructure, when we ask for the transitions between the different states."
 We have phase transitions. depending on the singular behaviour of the partition Function determined bv (he respective cdvnamies. as well as clustering ancl percolation transitions. determined by (he connectivity aspects of the svstem.," We have phase transitions, depending on the singular behaviour of the partition function determined by the respective dynamics, as well as clustering and percolation transitions, determined by the connectivity aspects of the system."
 Bul we then find that scaling ancl renormalization concepts lead to a universal description of critical phenomena. and critical exponents define universality Classes which contain quite different interaction forms.," But we then find that scaling and renormalization concepts lead to a universal description of critical phenomena, and critical exponents define universality classes which contain quite different interaction forms."
 When we study stronglyex. interactinge matter. we are therefore led to aspects which are relevant not only to QCD. but to the understanding of complex svstems in general.," When we study strongly interacting matter, we are therefore led to aspects which are relevant not only to QCD, but to the understanding of complex systems in general."
 What happens to stronglv interacting matter in the limit of high temperatures and densities?, What happens to strongly interacting matter in the limit of high temperatures and densities?
 This question has fascinated ΡΕΠΕ ever since (he discovery of the strong force and the multiple hadron production it leads to., This question has fascinated physicists ever since the discovery of the strong force and the multiple hadron production it leads to.
 Let us look at some of the features that have emerged over tlie vears., Let us look at some of the features that have emerged over the years.
no emission on scales greater than 1 kpe.,no emission on scales greater than 1 kpc.
 An exception to this rule is the CSO 01084-3388 which has some faint extended emission 1990)., An exception to this rule is the CSO 0108+388 which has some faint extended emission .
. Another exceptional source is (he CSO Jlits+5924 (NGC 3894) which also has faint kpe-scale emission and has been seen (o be slowly variable on timescales of vears 1998)., Another exceptional source is the CSO J1148+5924 (NGC 3894) which also has faint kpc-scale emission and has been seen to be slowly variable on timescales of years .
. J114384-5924 is relatively nearby at a redshift of 0.01085 1939).. and has a luminosity more than two orders of magnitude less than all other members of the CSO Class.," J1148+5924 is relatively nearby at a redshift of 0.01085 , and has a luminosity more than two orders of magnitude less than all other members of the CSO class."
 In 822 we present observations of seven CSOs used as calibrators in an intensive VLA monitoring campaign., In 2 we present observations of seven CSOs used as calibrators in an intensive VLA monitoring campaign.
 The calibration procedure is described in 833., The calibration procedure is described in 3.
 In 844 and 355 we consider primary and secondary systematic errors which allows us (to provide. in 366. a prescription lor obtaining hieh-precision monitoring.," In 4 and 5 we consider primary and secondary systematic errors which allows us to provide, in 6, a prescription for obtaining high-precision monitoring."
 Lastly. in 877 we consider how variability studies can be used to constrain physical models lor CSOs.," Lastly, in 7 we consider how variability studies can be used to constrain physical models for CSOs."
 The sample observed for (his paper consisted of seven CSOs chosen from (he samples of ancl (1996b)., The sample observed for this paper consisted of seven CSOs chosen from the samples of and .
. These particular sources were observed as part of a gravitational lens monitoring campaign and were chosen because (hey were bright and were located in the RA range defined bv the lens svstems., These particular sources were observed as part of a gravitational lens monitoring campaign and were chosen because they were bright and were located in the RA range defined by the lens systems.
 The sources and (heir coordinates are listed in Table 1.., The sources and their coordinates are listed in Table \ref{tab_obs}.
 The observations were made with the at 8.5 Gllz during its A. Dn. and D configurations. vielding tvpical angulaw resolutions of 0722 in the A conliguration to 0777 in the D configuration.," The observations were made with the at 8.5 GHz during its A, BnA, and B configurations, yielding typical angular resolutions of 2 in the A configuration to 7 in the B configuration."
 The monitoring program consisted of 88 observations of the sample between 1999 June 21 and 2000 February 1d. giving an average spacing of 2.7 d between epochs.," The monitoring program consisted of 88 observations of the sample between 1999 June 21 and 2000 February 14, giving an average spacing of 2.7 d between epochs."
 The tvpical observing block was one or two hours in length. and the (vpical integration time on each of the CSOs was 45 or GO sec.," The typical observing block was one or two hours in length, and the typical integration time on each of the CSOs was 45 or 60 sec."
 For each epoch. some fraction of (he CSO sample was observed al a large hour angle. since the CSOs are widely spaced on the sky and the observing blocks were short.," For each epoch, some fraction of the CSO sample was observed at a large hour angle, since the CSOs are widely spaced on the sky and the observing blocks were short."
 La fact. it was not always possible to observe the entire sample of CSOs in an observing block.," In fact, it was not always possible to observe the entire sample of CSOs in an observing block."
 However. (he incomplete sampling should not signilicantlv affect the conclusions of this paper.," However, the incomplete sampling should not significantly affect the conclusions of this paper."
 All of the CSOs were observed al least 68 times. giving quite adequate determinations of their light curves.," All of the CSOs were observed at least 68 times, giving quite adequate determinations of their light curves."
 In addition to the CSOs and the lens svstems. (wo pliase calibrators were observed as part of the monitoring program (Table 2)).," In addition to the CSOs and the lens systems, two phase calibrators were observed as part of the monitoring program (Table \ref{tab_pcal}) )."
 The (vpical integration times on the phase calibrators were between 2 and 3 min., The typical integration times on the phase calibrators were between 2 and 3 min.
 These two sources will also be discussed in this paper, These two sources will also be discussed in this paper
kim s1 relative. to the quasar (uncertaintiesMN due to slit. elfects have been taken into account).,km $^{-1}$ relative to the quasar (uncertainties due to slit effects have been taken into account).
 Phe faintest kinematic component is broader 260-4 LINES70) and is more highly ionized /LL3—7.0250.4)., The faintest kinematic component is broader $\pm$ 70) and is more highly ionized $\beta$ $\pm$ 0.4).
. His. possible that this. is the extension of the intermediate ionized. region. detected between the quasar and the knot. which. as proposed above. could be an ELLER associated with the quasar.," It is possible that this is the extension of the intermediate ionized region detected between the quasar and the knot, which, as proposed above, could be an EELR associated with the quasar."
 This is supported by the spatial continuity observed for the OLLI] line emission between this region and the fainter kinematic component of the knot along the slit., This is supported by the spatial continuity observed for the [OIII] line emission between this region and the fainter kinematic component of the knot along the slit.
 Therefore. the morphology of SDSS 0511-00 is reminiscent of a galaxy witha continuum tidal tail which is a signature of an merger event.," Therefore, the morphology of SDSS J0217-00 is reminiscent of a galaxy with a continuum tidal tail which is a signature of an merger event."
 The quasar is associated with a companion star forming knot where the gas is photoionized bv the voung stars., The quasar is associated with a companion star forming knot where the gas is photoionized by the young stars.
" An EELR. is detected across 4"" or 19 kpe between both objects and it might connect them. physically.", An EELR is detected across $\sim$ $\arcsec$ or 19 kpc between both objects and it might connect them physically.
 No extended: structures are detected in the broad. or the narrow band images (Fig., No extended structures are detected in the broad or the narrow band images (Fig.
 S) of this quasar., 8) of this quasar.
 The faint object. niarked as Cl in the figure falls within the slit.," The faint object, marked as G1 in the figure falls within the slit."
 The spectrum shows that it is a faint continuum source of unknown z., The spectrum shows that it is a faint continuum source of unknown $z$.
 The spatial profile of the OLI]. line (Fig., The spatial profile of the [OIII] line (Fig.
" 9) is dominated by a spatially unresolved component of FWILL O.57£0.05"" (vs. 0.6740.04"" ENLIN seeing measured on the images).", 9) is dominated by a spatially unresolved component of FWHM $\pm$ $\arcsec$ (vs. $\pm$ $\arcsec$ FWHM seeing measured on the images).
 There is no evidence for extended line emission alone PA 171. associated. with this quasar αἲ surface brightness. levels 230=3.5. Ls7 erg 1 72 , There is no evidence for extended line emission along PA 171 associated with this quasar at surface brightness levels $\ga$ $\sigma$ $\times$ $^{-18}$ erg $^{-1}$ $^{-2}$ $^{-2}$.
The objects discussed below were observed. during the 2009 run., The objects discussed below were observed during the 2009 run.
 As mentioned in 92. no narrow or internicciate emission line|continuum images are available.," As mentioned in $\delta$ 2, no narrow or intermediate emission line+continuum images are available."
 Therefore. except in two cases where the continuum dominated images revealed: or suggested. a galaxy interaction or a peculiar quasar morphology. the slit was placed. blindly or through objects near the quasar. with the goal of checking their z.," Therefore, except in two cases where the continuum dominated images revealed or suggested a galaxy interaction or a peculiar quasar morphology, the slit was placed blindly or through objects near the quasar, with the goal of checking their $z$."
 The narrow band continuum image is shown in Fig., The narrow band continuum image is shown in Fig.
 10., 10.
 The emission line galaxies €1 and G2 fall within the slits., The emission line galaxies G1 and G2 fall within the slits.
 C2 lies at 220.221., G2 lies at $z$ =0.221.
 GL emits a single emission line whose most likely identification is OLIJA3727 at 220.950 (rather than Lye at > =4.98). given the strong continuum emission blucwards of the line.," G1 emits a single emission line whose most likely identification is $\lambda$ 3727 at $z$ =0.950 (rather than $\alpha$ at $z=$ 4.98), given the strong continuum emission bluewards of the line."
 The spatial distribution of OIL] (Fig., The spatial distribution of [OIII] (Fig.
" 11) along PA 65 is dominated by a compact component which is successfully represented by a Gaussian function of PWLIAI=1.05+0.05"" (vs. seeing FWLAI=0.88+0.08).", 11) along PA 65 is dominated by a compact component which is successfully represented by a Gaussian function of $\pm$ $\arcsec$ (vs. seeing $\pm$ 0.08).
 Taking errors into account ancl possible seeing variations during the spectroscopic observations it is not. possible to discern whether the OLLI] profile along PAGS is spatially resolved., Taking errors into account and possible seeing variations during the spectroscopic observations it is not possible to discern whether the [OIII] profile along PA65 is spatially resolved.
" Phe emission lines are unresolved along PA 153 (PWIIMA0.8240.04"").", The emission lines are unresolved along PA 153 $\pm$ $\arcsec$ ).
 The broad band image is shown in Fig., The broad band image is shown in Fig.
 12., 12.
 Several galaxies, Several galaxies
The new versionH of our code preseuted here solves the primitiveHma equationsH ofH usingH the dynamical core originally developed by Hoskins&Sitmmuuous(1975).. converted for hot Jupiter studies in Menou&Rauscher(2009) aud used in subsequent papers (Rauscher&Menou2010:Pernaetal.2010a.b:Miller-RicciKempton&aus‘ther 2011).,"The new version of our code presented here solves the primitive equations of using the dynamical core originally developed by \citet{Hoskins1975}, converted for hot Jupiter studies in \citet{MR09} and used in subsequent papers \citep{RM10,Perna2010a,Perna2010b,Kempton2011}."
. The previous version of this code (IGCNIL) lieated the atmosphere through a Newtouiau relaxation scheme., The previous version of this code (IGCM1) heated the atmosphere through a Newtonian relaxation scheme.
 Here we use au updated version of the code (GCS.deF.Forsteretal.οί00) with several improvements. includiug a simple radiative trausfer scheme. whicl we have adapted to moclel gaseous exoplanets.," Here we use an updated version of the code \citep[IGCM3,][]{Forster2000} with several improvements, including a simple radiative transfer scheme, which we have adapted to model gaseous exoplanets."
 This code also coutalus a standard dry couvectiou scheme., This code also contains a standard dry convection scheme.
 The static stability throughout each vertical columu is calculated aud auy layers fouud to be couvectively uustable are adjusted to coustaut potential temperature (eutropy). with the vertical integral of enthalpy couservec.," The static stability throughout each vertical column is calculated and any layers found to be convectively unstable are adjusted to constant potential temperature (entropy), with the vertical integral of enthalpy conserved."
 The primary driver of attnospheric circulation on close-in extrasolar planets is heating from the stellar irradiation. whieh is absorbed aid re-emitted throughout the atinosphere.," The primary driver of atmospheric circulation on close-in extrasolar planets is heating from the stellar irradiation, which is absorbed and re-emitted throughout the atmosphere."
 We employ the common assumption that the incicent radiation is predominantly at optical wavelengths. while the re-emitted radiation is in the iufrared.," We employ the common assumption that the incident radiation is predominantly at optical wavelengths, while the re-emitted radiation is in the infrared."
 We cai then separate our treatinent of radiative trausfer iuto two wavelength ranges. with downward opical flux absorbed as it peuetrates deeper into the atmosphere aud iulrared flix absorbed and emitted isotropically throughout the atiuosphere.," We can then separate our treatment of radiative transfer into two wavelength ranges, with downward optical flux absorbed as it penetrates deeper into the atmosphere and infrared flux absorbed and emitted isotropically throughout the atmosphere."
 We calculate the upward and downward uet [Iuxes ii eac1 vertical column of the atmosphere aud then solve for the localized specific heaine asi dP/dt=(gfcy)dP /dP. where dT/dl is the temperature tendency. g is the gravitatioral ace'eleration (assumec to be constant through the model). cj is the specific heat. aud dF/dP is the verIca derivative of te net vertical [lux (pressure. 2. is our vertical coordinate).," We calculate the upward and downward net fluxes in each vertical column of the atmosphere and then solve for the localized specific heating as: $dT/dt = (g/c_p)dF/dP$ , where $dT/dt$ is the temperature tendency, $g$ is the gravitational acceleration (assumed to be constant through the model), $c_p$ is the specific heat, and $dF/dP$ is the vertical derivative of the net vertical flux (pressure, $P$, is our vertical coordinate)."
 Upward flux. ir the cirection of decreasiug pressure. is defiued to be positive.," Upward flux, in the direction of decreasing pressure, is defined to be positive."
 For au atmosphere in local radiative equiibriuu (i.e. withot dynamics) dF7/dP=0. either because there is a constant [Iux through tle all105here (e.g. an outward heat flux from the planet interior) or because the optical heatjs) Is exacly oiauced by iurared cooling.," For an atmosphere in local radiative equilibrium (i.e. without dynamics) $dF/dP=0$, either because there is a constant flux through the atmosphere (e.g. an outward heat flux from the planet interior) or because the optical heating is exactly balanced by infrared cooling."
 The heating rate feeds into the euergy equation of the dyuanies solver (Q in Equation 1d of Menou Rauscher 2009)., The heating rate feeds into the energy equation of the dynamics solver $Q$ in Equation 1d of Menou Rauscher 2009).
 The incident optical fhX altle op of each couuu of the atmosphere. with latitude © aud longitude 0. is F(ó.0)=(1—A)Finecos(@)cos(0) For the day and £=0 on the night side. where Fine is the stellar Παν iucident ou the top of he atiuosphere at the substellar point. aud 4 isthe Boucl albedo. which here we assume to be zero.," The incident optical flux at the top of each column of the atmosphere, with latitude $\phi$ and longitude $\theta$, is $F(\phi,\theta)=(1-A)F_{\mathrm{inc}} \cos(\phi) \cos(\theta)$ for the day and $F=0$ on the night side, where $F_{\mathrm{inc}}$ is the stellar flux incident on the top of the atmosphere at the substellar point and $A$ isthe Bond albedo, which here we assume to be zero."
 We assume a well-inixecl optical absorber, We assume a well-mixed optical absorber
",MMyr! (the median mass of those objects is ~0.7 Msolar), while at ~25 MMyr it drops by an order of magnitude to ~20 ,MMyr! (median mass ~1 Msolar)).","$^{-1}$ (the median mass of those objects is $\sim 0.7$ ), while at $\sim 
25$ Myr it drops by an order of magnitude to $\sim 20$ $^{-1}$ (median mass $\sim 1$ )."
" At face value, the star formation strength for stars in the range 0.4—4.0 thus appears to be much higher in the present burst than in the one while the total Myrintegratedago outputmight (i.e. the total number of stars, as shown by the thick solid histogram) in the two episodes is rather similar."," At face value, the star formation strength for stars in the range $0.4-4.0$ thus appears to be much higher in the present burst than in the one while the total integrated output (i.e. the total number of stars, as shown by the thick solid histogram) in the two episodes is rather similar."
" However, there are selection effects that one must consider."," However, there are important selection effects that one must consider."
" As mentioned importantabove, one effect is the fraction of PMS stars with Ho excess, which most likely varies throughout the PMS phase, making it more difficult to compare the number of younger and older PMS stars to one another."," As mentioned above, one effect is the fraction of PMS stars with $\alpha$ excess, which most likely varies throughout the PMS phase, making it more difficult to compare the number of younger and older PMS stars to one another."
" Another problem is that, while the previous burst has certainly ended, the current one might still continue for a long time, thus making it hard to predict how many stars will eventually be formed."," Another problem is that, while the previous burst has certainly ended, the current one might still continue for a long time, thus making it hard to predict how many stars will eventually be formed."
" Finally, the accuracy on relative ages being at best of order a factor of V2, it is not possible to say exactly how long the previous burst lasted nor how many short bursts took place in the time frame covered by one age bin."," Finally, the accuracy on relative ages being at best of order a factor of $\sqrt{2}$, it is not possible to say exactly how long the previous burst lasted nor how many short bursts took place in the time frame covered by one age bin."
" This carries the that, if there was just one short burst, the star formation implicationstrength might have been comparable to or even higher than that of the current episode."," This carries the implication that, if there was just one short burst, the star formation strength might have been comparable to or even higher than that of the current episode."
" The results of our quantitative age analysis on the star formation efficiency in this field are still tentative, but it will be possible to lift at least some of the necessarilyuncertainties plaguing this picture through a systematic comparison of different star forming clusters in similar states of evolution that we plan to conduct in the future."," The results of our quantitative age analysis on the star formation efficiency in this field are still necessarily tentative, but it will be possible to lift at least some of the uncertainties plaguing this picture through a systematic comparison of different star forming clusters in similar states of evolution that we plan to conduct in the future."
" On the other hand, having established that there are at least two star formation episodes in 33406, separated by ~10 MMyr or more, we still can study whether they are independent of one another or appear to be causally connected."," On the other hand, having established that there are at least two star formation episodes in 346, separated by $\sim 10$ Myr or more, we still can study whether they are independent of one another or appear to be causally connected."
spectrum is too blue.,spectrum is too blue.
 A larger mass transfer rate increases (he flux at all wavelengths. making {he mean residual larger. and produces a svuthetic spectrum that is too blue.," A larger mass transfer rate increases the flux at all wavelengths, making the mean residual larger, and produces a synthetic spectrum that is too blue."
 A smaller mass transfer rate decreases the flux at all wavelengths. again making the mean residual larger. and produces a svnthetic spectrum that is too blue.," A smaller mass transfer rate decreases the flux at all wavelengths, again making the mean residual larger, and produces a synthetic spectrum that is too blue."
 The Fig., The Fig.
 4 residuals suggest a departure from a Standard Model., 4 residuals suggest a departure from a Standard Model.
 Stationary models are possible (Lasota2001) whose £() profiles differ [rom the Standard Model., Stationary models are possible \citep{l2001} whose $T(R)$ profiles differ from the Standard Model.
 Examples include systems in which inradiation bv the WD is important. or svstems in which the mass {rausler stream directly. heats the outer regions of the accretion disk (Visninc&Diamond 1994).. (but see Armitage&Livio 1993)). or there is an increased mass [lux through the outer part of the accretion disk for unspecified reasons 1986).," Examples include systems in which irradiation by the WD is important, or systems in which the mass transfer stream directly heats the outer regions of the accretion disk \citep{vd89,vd92,sm94}, (but see \citealt{al98}) ), or there is an increased mass flux through the outer part of the accretion disk for unspecified reasons \citep*{kh86}."
. In all of our synthetic spectra. the synthetic spectrum has absorption lines where theIUE speciyum shows emission lines.," In all of our synthetic spectra, the synthetic spectrum has absorption lines where the spectrum shows emission lines."
" SPT95 argue that the broad emission lines in (he high"" V—13— 14) state arise [rom the aceretion disk. in contrast to narrow (Fig."," SPT95 argue that the broad emission lines in the ""high"" ${\sim}13-14$ ) state arise from the accretion disk, in contrast to narrow (Fig."
 2) emission lines produced on the irradiated secondary when MV. Lvr is in a low state., 2) emission lines produced on the irradiated secondary when MV Lyr is in a low state.
 Two variations of the Standard Model TUR) relation affect (he corresponding system svithetic spectrum: (1) a change in the functional dependence of Z(H) on R 2003).. and (2) truncation of the accretion disk at some inner radius larger (han the WD radius (Longetal. 1994)..,"	 Two variations of the Standard Model T(R) relation affect the corresponding system synthetic spectrum: (1) a change in the functional dependence of $T(R)$ on $R$ \citep{ow03}, and (2) truncation of the accretion disk at some inner radius larger than the WD radius \citep{L94}. ."
 Physically. (runeation can occur either by evaporation of (he inner disk annuli (Hameuryetal.2000). or by. accretion disk interaction with a magnetic field associated with the WD (IILO2).," Physically, truncation can occur either by evaporation of the inner disk annuli \citep{h2000}
 or by accretion disk interaction with a magnetic field associated with the WD (HL02)."
 We have investigated both options., We have investigated both options.
 The first permits a good fit to theZCE spectra but the models are slishtlv inferior to the second option models., The first permits a good fit to the spectra but the models are slightly inferior to the second option models.
 The second option truncates the accretion disk at an inner radius larger than the WD radius (Longetal.1994) and preserves the Standard Model 7(22) functional dependence on H., The second option truncates the accretion disk at an inner radius larger than the WD radius \citep{L94} and preserves the Standard Model $T(R)$ functional dependence on $R$.
 Since the distance to MV Lyr is a fixed quantity. as the truncation radius of (he accretion disk is increased. producing an initial reduction in flux from the model. and a resulting discrepancy wilh observed speclra. a compensating increase in the adopted mass transfer rate is necessary {ο restore agreement with the observed spectra.," Since the distance to MV Lyr is a fixed quantity, as the truncation radius of the accretion disk is increased, producing an initial reduction in flux from the model, and a resulting discrepancy with observed spectra, a compensating increase in the adopted mass transfer rate is necessary to restore agreement with the observed spectra."
 We started will the mass (ransler rate of Fig., We started with the mass transfer rate of Fig.
 4. then incrementally increased (he (runcation radius and made compensating increases in (he mass transfer rate.," 4, then incrementally increased the truncation radius and made compensating increases in the mass transfer rate."
 The quality of the successive lits improved but remained perceptably discrepant (too blue) even up to a limiting trimeation radius of Γρνα=IS., The quality of the successive fits improved but remained perceptably discrepant (too blue) even up to a limiting truncation radius of $r_b/r_{\rm wd}=18$.
 At this truncation radius. of the tidal cutoff radius. the necessary mass (transfer rate becomes 1.0x10.V.vr.|. close to the one used in L1999. and one dex larger than the value given by IEuneury et al. (," At this truncation radius, of the tidal cutoff radius, the necessary mass transfer rate becomes $1.0{\times}10^{-8}{\cal M}_{\odot}{\rm yr}^{-1}$, close to the one used in L1999, and one dex larger than the value given by Hameury et al. ("
1983) at the long period boundary of the period gap.,1988) at the long period boundary of the period gap.
 We conclude that a Standarcl Mocel TUR) with a truncated accretion disk model produces a svnthetic spectrumthat is too blue and so does notfit the observed MV. Lyr ΠΕ spectra.,	 We conclude that a Standard Model T(R) with a truncated accretion disk model produces a synthetic spectrumthat is too blue and so does notfit the observed MV Lyr spectra.
ave cllicient tracers of cosmic structure across large volumes (Brown et 22003: Zehavi ct 22005a: Eisenstein οἱ 22005b).,are efficient tracers of cosmic structure across large volumes (Brown et 2003; Zehavi et 2005a; Eisenstein et 2005b).
 Moreover. these objects provide particularly reliable photometric redshifts owing to the strong spectral break at s4000. in the galaxy rest frame and the consequent rapid variation with redshift of their observed: colours in the SDSS filter system (Padmanabhan et 22005: Collister et 22007).," Moreover, these objects provide particularly reliable photometric redshifts owing to the strong spectral break at $\approx 4000$ in the galaxy rest frame and the consequent rapid variation with redshift of their observed colours in the SDSS filter system (Padmanabhan et 2005; Collister et 2007)."
 Furthermore. the photometric redshift’ error. distribution of the LRGs can be accurately calibrated owing to the existence of spectroscopic observations of a sub-sample as part of the 2dE-SDSS LRG and Quasar (25LAQ) survey at the Anglo-Australian Telescope (Cannon Atal. 22006).," Furthermore, the photometric redshift error distribution of the LRGs can be accurately calibrated owing to the existence of spectroscopic observations of a sub-sample as part of the 2dF-SDSS LRG and Quasar (2SLAQ) survey at the Anglo-Australian Telescope (Cannon et 2006)."
" The combination of these datasets has allowed: us to construct a large catalogue of more than LO"" LRGs spanning the redshift range 0.4<z0.7 with an r.nis."," The combination of these datasets has allowed us to construct a large catalogue of more than $10^6$ LRGs spanning the redshift range $0.4 <
z < 0.7$ with an r.m.s."
 photometric redshift error σ.=Vs(02)?«Oz52zmO.OB(L|i) where 92—μμnpe," photometric redshift error $\sigma_z \equiv
\sqrt{<(\delta z)^2> - <\delta z>^2} \approx 0.03 (1+z)$ where $\delta
z \equiv z_{\rm phot} - z_{\rm spec}$."
" We have dubbed this catalogue ""MegaZ-LIG (Collister et 22007).", We have dubbed this catalogue “MegaZ-LRG” (Collister et 2007).
 The database covers almost 6.000 dee. or an ellective volume z2.5 5.? CGpc*.," The database covers almost $6{,}000$ $^2$, or an effective volume $\approx 2.5$ $h^{-3}$ $^3$ ."
 We have already.lv used this catalogue to measure the large-scale clustering of the galaxies on linear and. quasi-Iinear scales via a power spectrum analysis ancl thereby. extract measurements of cosmological parameters. (Blake οἱ 22007: see also Pacimanabhan et 22007)., We have already used this catalogue to measure the large-scale clustering of the galaxies on linear and quasi-linear scales via a power spectrum analysis and thereby extract measurements of cosmological parameters (Blake et 2007; see also Padmanabhan et 2007).
 In this study we turn to the smiall-scale clustering properties of the LRGs., In this study we turn to the small-scale clustering properties of the LRGs.
 Our goal is to present new measurements of the small-scale correlation funetion of LRGs at 2~0.5. and to connect these measurements to the physical manner in which these LRGs populate cark matter haloes.," Our goal is to present new measurements of the small-scale correlation function of LRGs at $z \sim 0.5$, and to connect these measurements to the physical manner in which these LRGs populate dark matter haloes."
 lt has long been known that different classes of galaxy possess dillerent clustering properties. in a manner connected. to their small-scale environments. (for. recent observational studies we refer the reader to Norberg et 22002: Dudavari ct 22003: Hogg et 22003: Zehavi et 22005b: and references therein).," It has long been known that different classes of galaxy possess different clustering properties, in a manner connected to their small-scale environments (for recent observational studies we refer the reader to Norberg et 2002; Budavari et 2003; Hogg et 2003; Zehavi et 2005b; and references therein)."
 For many vears these differing clustering properties could. be adequately cleseribecl by fitting a simple power-law function to the two-point galaxy correlation function on small scales. (Pechles 1980)., For many years these differing clustering properties could be adequately described by fitting a simple power-law function to the two-point galaxy correlation function on small scales (Peebles 1980).
 LLowever. recent surveys have measured the clustering pattern accurately enough. to. detect. deviations from. the canonical clustering power-law Lblawkins ct 22003: Zehavi et 22004: Zheng 2004: Eisenstein et 22005a: Zehavict 22005a.b: Phleps et 22006).," However, recent surveys have measured the clustering pattern accurately enough to detect deviations from the canonical clustering power-law Hawkins et 2003; Zehavi et 2004; Zheng 2004; Eisenstein et 2005a; Zehavi et 2005a,b; Phleps et 2006)."
 These deviations provide an important insight into the processes of galaxy formation., These deviations provide an important insight into the processes of galaxy formation.
" The richer structure in the galaxy clustering pattern revealed by recent surveys has been successfully interpreted in terms of the “halo model"" SSeljak 2000: Peacock Smith 2000: Scoccimarro et 22001: Cooray Sheth 2002: Aerlind Weinberg 2002: Ixravtsov et 22004: Zehavi et 22004: Zheng 2004: Zheng et 22005: Zehavi οἱ 22005b: Collister Lahay 2005: Tinker et 22005).", The richer structure in the galaxy clustering pattern revealed by recent surveys has been successfully interpreted in terms of the “halo model” Seljak 2000; Peacock Smith 2000; Scoccimarro et 2001; Cooray Sheth 2002; Berlind Weinberg 2002; Kravtsov et 2004; Zehavi et 2004; Zheng 2004; Zheng et 2005; Zehavi et 2005b; Collister Lahav 2005; Tinker et 2005).
 In this mocdel. the small-scale clustering of a distribution of galaxies is linked to the underlying network of dark matter haloes. whose properties can be measured. using cosmological simulations.," In this model, the small-scale clustering of a distribution of galaxies is linked to the underlying network of dark matter haloes, whose properties can be measured using cosmological simulations."
" A class of galaxies is assumed to populate haloes in accordance with a statistical ""halo occupation distribution"" as a function of the halo mass.", A class of galaxies is assumed to populate haloes in accordance with a statistical “halo occupation distribution” as a function of the halo mass.
 The clustering then naturally separates into two components: the distribution of galaxies within individual haloes. which dominates on small scales (1 Alpe). and the mutual clustering of galaxies inhabiting separate haloes. which dominates on larger scales (21 Alpe).," The clustering then naturally separates into two components: the distribution of galaxies within individual haloes, which dominates on small scales $\la 1$ Mpc), and the mutual clustering of galaxies inhabiting separate haloes, which dominates on larger scales $\ga 1$ Mpc)."
 The combination of these two terms can accurately model the observed: scale-dependent features in the small-scale clustering pattern., The combination of these two terms can accurately model the observed scale-dependent features in the small-scale clustering pattern.
" The outline of this paper is as follows: in Section 2 we briclly describe the ""MegaZ-LIG"" dataset.", The outline of this paper is as follows: in Section \ref{secdata} we briefly describe the “MegaZ-LRG” dataset.
 We then resent measurements of the angular correlation. function in narrow redshift slices together with simple. power-aw fits in Section 3.., We then present measurements of the angular correlation function in narrow redshift slices together with simple power-law fits in Section \ref{secmeas}.
 In Section 4— we introduce. the ramework of the halo model and in Section 5. we re-fit he clustering measurements by parameterizing the halo occupation distribution of the LRGs. comparing our results ο previous work in Section 6..," In Section \ref{sechalo} we introduce the framework of the halo model and in Section \ref{secparfit} we re-fit the clustering measurements by parameterizing the halo occupation distribution of the LRGs, comparing our results to previous work in Section \ref{secprev}."
 Section 7. investigates a range of potential systematic photometric errors in the catalogue hat may bias our results., Section \ref{secsys} investigates a range of potential systematic photometric errors in the catalogue that may bias our results.
 We conclude in Section 8S.., We conclude in Section \ref{secconc}.
" '""hroughout our study we assume a fixed set of large-scale cosmological parameters: fractional matter density Ow=0.3. cosmological constant ον=0.7. curvature Qi=Q0. Hubble. parameter f=0.7. fractional barvon ensity fi,=O,/0,4,0.15. slope of the primordial power μα»ectrum. ης=1. and overall normalization of the power μα»ectrum ox=0.5."," Throughout our study we assume a fixed set of large-scale cosmological parameters: fractional matter density $\Omega_{\rm m} = 0.3$, cosmological constant $\Omega_\Lambda = 0.7$, curvature $\Omega_{\rm
 k} = 0$, Hubble parameter $h = 0.7$, fractional baryon density $f_{\rm b} = \Omega_{\rm b}/\Omega_{\rm m} = 0.15$, slope of the primordial power spectrum $n_{\rm s} = 1$, and overall normalization of the power spectrum $\sigma_8 = 0.8$."
 These values are consistent with fits to 10 large-scale power spectrum of the LRGs (Blake et 2007) and to the latest measurements of the anisotropy μα»ectrum of the Cosmic Microwave Background (Sperecl et 22007)., These values are consistent with fits to the large-scale power spectrum of the LRGs (Blake et 2007) and to the latest measurements of the anisotropy spectrum of the Cosmic Microwave Background (Spergel et 2007).
 We note that there is some cegencracy between je assumed. cosmological parameters and the fitted halo model parameters. ancl we investigate how the best-fitting halo model parameters depend on the values of ax ancl ης.," We note that there is some degeneracy between the assumed cosmological parameters and the fitted halo model parameters, and we investigate how the best-fitting halo model parameters depend on the values of $\sigma_8$ and $n_{\rm s}$."
 We analyze angular clustering in the “AlegaZ-Li” galaxy atabase. a photometric-redshift. catalogue of Luminous lted Galaxies (LRGs) based on the imaging dataset of the SDSS 4th. Data Release.," We analyze angular clustering in the “MegaZ-LRG” galaxy database, a photometric-redshift catalogue of Luminous Red Galaxies (LRGs) based on the imaging dataset of the SDSS 4th Data Release."
 The. construction of this catalogue is described in detail by Collister ct ((2007). ancl cosmological-parameter fitting to the angular power spectrum. is presented. by Blake ct ((2007).," The construction of this catalogue is described in detail by Collister et (2007), and cosmological-parameter fitting to the angular power spectrum is presented by Blake et (2007)."
 We only provide a brief description of the catalogue here. refering the reader to these two papers for more information.," We only provide a brief description of the catalogue here, refering the reader to these two papers for more information."
" AlegaZ-LRG contains over 10"" LRGs spanning the redshift raneo Q4c2<OF with an r.anns."," MegaZ-LRG contains over $10^6$ LRGs spanning the redshift range $0.4 <
z < 0.7$ with an r.m.s."
 redshift error ὃςzz00.0212)., redshift error $\delta z \approx 0.03 (1+z)$.
 The angular selection. function is described. by Blake ct ((2007) ancl encompasses. 5914 deg? (the three southern. SDSS stripes are excluded)., The angular selection function is described by Blake et (2007) and encompasses 5914 $^2$ (the three southern SDSS stripes are excluded).
 The sample was selected from the SDSS imaging database using a series of colour and magnitude cuts (IEisenstein et 22001: Collister et 22007). which amount to a maegnitucle-limited sample of Luminous Red Galaxies with ;-band magnitudes 17.5«i20.," The sample was selected from the SDSS imaging database using a series of colour and magnitude cuts (Eisenstein et 2001; Collister et 2007), which amount to a magnitude-limited sample of Luminous Red Galaxies with $i$ -band magnitudes $17.5 < i < 20$."
 Photometric redshifts were derived for these ealaxies using an Artificial Neural Network method. ANNz (Firth. Lahay Somerville 2003: Collister Lahay 2004). constrained by a spectroscopic sub-samiple of z19.900 galaxies obtained by the 2SLAQ survey (Cannon ct 22006).," Photometric redshifts were derived for these galaxies using an Artificial Neural Network method, ANNz (Firth, Lahav Somerville 2003; Collister Lahav 2004), constrained by a spectroscopic sub-sample of $\approx 13{,}000$ galaxies obtained by the 2SLAQ survey (Cannon et 2006)."
 In this paper we analyze a “conservative” version of the database in which the selection cuts applied to the imaging database are identical to those used. to. produce the great majority. of the spectroscopic sub-sample. tthe faint magnitude limit is briehtened to¢= 18)., In this paper we analyze a “conservative” version of the database in which the selection cuts applied to the imaging database are identical to those used to produce the great majority of the spectroscopic sub-sample the faint magnitude limit is brightened to$i = 19.8$ ).
 Star-galaxy separation cuts were applied both in the initial, Star-galaxy separation cuts were applied both in the initial
brighter than in a model where they emit constantly.,brighter than in a model where they emit constantly.
 Unfortunately. the number of LAINBs that have been discovered actively accreting is too small (~100) to determine which effect dominates for duty cveles less than ten percent.," Unfortunately, the number of LMXBs that have been discovered actively accreting is too small $\sim 100$ ) to determine which effect dominates for duty cycles less than ten percent."
 llowever. if low-mass N-rav binaries are assumed to be the exclusive progenitors οἱ millisecond pulsars (see?.foranalternative).. the demoeraphics of the millisecond pulsars in (he Galaxy aud bevond provides an estimate of the number of sources.," However, if low-mass X-ray binaries are assumed to be the exclusive progenitors of millisecond pulsars \citep[see][ for
an alternative]{1984JApA....5..209V}, the demographics of the millisecond pulsars in the Galaxy and beyond provides an estimate of the number of sources."
 Specifically. if the duration of the epochs when (he neutron star is emitting gravitational radiation is greater than 10.000. vears. several objects in the Galaxy will be above the detection thresholds of LIGO.," Specifically, if the duration of the epochs when the neutron star is emitting gravitational radiation is greater than 10,000 years, several objects in the Galaxy will be above the detection thresholds of LIGO."
" For ty,~ 10.000 vear. these sources would be detectable throughout the Local Group."," For $\tau_\rmscr{on} \sim$ 10,000 year, these sources would be detectable throughout the Local Group."
 Those objects in the Galaxy could also be detected from their X-ray emission whieh would be powered by &ravitational radiation reaction: they would be GW-powered neutron stars., Those objects in the Galaxy could also be detected from their X-ray emission which would be powered by gravitational radiation reaction; they would be GW-powered neutron stars.
" For 7, much less than 10.000. vears and much greater than one vear. no sources are likely to be detectable even with an enhanced LIGO «detector."," For $\tau_\rmscr{on}$ much less than 10,000 years and much greater than one year, no sources are likely to be detectable even with an enhanced LIGO detector."
 However. if the duration of eravitational wave emission per cvcle is less (han several vears. several sources could be detected by an enhanced LIGO.," However, if the duration of gravitational wave emission per cycle is less than several years, several sources could be detected by an enhanced LIGO."
 In this case. the sources will be located at a tvpical distance of 1 Gpe.," In this case, the sources will be located at a typical distance of 1 Gpc."
 Depending on the nature of the r—imocle instability in (he cores of quickly spinning neutron stars. LMXDs may provide eravitational-wave beacons throughout the Galaxy. and the Local Group or al cosmologically signilicant distances.," Depending on the nature of the $r-$ mode instability in the cores of quickly spinning neutron stars, LMXBs may provide gravitational-wave beacons throughout the Galaxy and the Local Group or at cosmologically significant distances."
 I would like to thank Phil Arras for useful discussions. and | was supported by the Chandra Postdoctoral Fellowship Award ZZPPF0-10015 issued by the Chandra X-ray Observatory Center. which is operated by the Smithsonian Astrophysical Observatory for and on behalf ol NASA under contract NASS-39073.," I would like to thank Phil Arras for useful discussions, and I was supported by the Chandra Postdoctoral Fellowship Award PF0-10015 issued by the Chandra X-ray Observatory Center, which is operated by the Smithsonian Astrophysical Observatory for and on behalf of NASA under contract NAS8-39073."
First the surface density profiles are evaluated by taking into account all stars in each sample. and then additionally only the bound stars. to show where and in which way potential escapers inlluence the profile.,"First the surface density profiles are evaluated by taking into account all stars in each sample, and then additionally only the bound stars, to show where and in which way potential escapers influence the profile."
 Even though physically motivated: models. like Wing or Wilson(1975) have often proved. to. more generally represent observations of globular clusters than analytical templates (c.g. AleLaughlin&vancerMarcel 2005)). they have one decisive Law: there is no simple analytical description which can be quickly fitted to star cluster profiles.," Even though physically motivated models like \citet{King66} or \citet{Wilson75} have often proved to more generally represent observations of globular clusters than analytical templates (e.g. \citealt{McLaughlin05}) ), they have one decisive flaw: there is no simple analytical description which can be quickly fitted to star cluster profiles."
 Thus. due to their simplicity. analytical templates are still in use. especially for larger data sets of extra-Galactic Aloreover. analytical templates have the advantage that they can be easily. set up and extended. according to the needs of the corresponding investigation.," Thus, due to their simplicity, analytical templates are still in use, especially for larger data sets of extra-Galactic Moreover, analytical templates have the advantage that they can be easily set up and extended, according to the needs of the corresponding investigation."
 We found that star clusters consist basically of three parts: core. bulk and tidal debris.," We found that star clusters consist basically of three parts: core, bulk and tidal debris."
 “Phe templates which have been set. up in the past were adjusted according to the cdillerent. foci of the corresponding investigations., The templates which have been set up in the past were adjusted according to the different foci of the corresponding investigations.
 For example. Wine(1962) was mainly interested in the bulk of the cluster. and so the template mainly. consists of a bulk and just a lat core. whereas Elson.Fall&Freeman(1987) focused on the tidal debris of voung clusters. and so their template ms a [lat core. no separated. explicitly defined bulk but a »ower-Ilw distribution of tidal debris.," For example, \citet{King62} was mainly interested in the bulk of the cluster, and so the template mainly consists of a bulk and just a flat core, whereas \citet{Elson87}
 focused on the tidal debris of young clusters, and so their template has a flat core, no separated, explicitly defined bulk but a power-law distribution of tidal debris."
 Therefore. we will study some existing templates in a systematic way to show heir advantages and their Haws. and finally design a more general template.," Therefore, we will study some existing templates in a systematic way to show their advantages and their flaws, and finally design a more general template."
 An overview of all templates is given in ‘Tab. 2.., An overview of all templates is given in Tab. \ref{table2}.
 The celassic templates we use. their abbreviations for the rest of the text. anc their fit parameters are. the following:," The `classic' templates we use, their abbreviations for the rest of the text, and their fit parameters are the following:"
dynamical instability. when the Uat-entropy core of. the secondary is being exposed and mass transfer. becomes catastrophic.,dynamical instability when the flat-entropy core of the secondary is being exposed and mass transfer becomes catastrophic.
 This will ultimately lead. to the dvnanical disruption of the remnant core (Gf it is relatively unevolved)., This will ultimately lead to the dynamical disruption of the remnant core (if it is relatively unevolved).
 In this paper we present the first part of a systematic study of the slow merger of massive stars. concentrating on the interaction of the mass stream emanating from the mass-losing immersed star with the core of the giant.," In this paper we present the first part of a systematic study of the slow merger of massive stars, concentrating on the interaction of the mass stream emanating from the mass-losing immersed star with the core of the giant."
 This stream will generally not intersect with itself and form an accretion disc. but. will instead. penetrate deep into the core of the giant before it is stopped. by the increasing pressure inside the core.," This stream will generally not intersect with itself and form an accretion disc, but will instead penetrate deep into the core of the giant before it is stopped by the increasing pressure inside the core."
 As we will show. the stream may be able to reach a radius. as small as ~10 cm.," As we will show, the stream may be able to reach a radius as small as $\sim 10^{10}\,$ cm."
 Since this is deep inside the hvdrogen-exhausted helium core of a star that has already exhausted. helium in the core. this may have important consequences: (1) lt can lead. to the dredge-up of helium from the core which can then be mixed with the rest of the envelope. (," Since this is deep inside the hydrogen-exhausted helium core of a star that has already exhausted helium in the core, this may have important consequences: (1) It can lead to the dredge-up of helium from the core which can then be mixed with the rest of the envelope. ("
2) Since the material in the stream is hydrogen-rich and is heated to temperatures as high as several 10 Ix. in the post-impact region. this niv lead to some unusual nucleosvnthesis. in. particular. s-processing.,"2) Since the material in the stream is hydrogen-rich and is heated to temperatures as high as several $10^8\,$ K in the post-impact region, this may lead to some unusual nucleosynthesis, in particular s-processing."
 How deep the stream can penetrate depends on the initial entropy of the stream material. but also its initial angular momentum and most importantly on the structure near the core and the amount of entropy generated in the interaction with the core.," How deep the stream can penetrate depends on the initial entropy of the stream material, but also its initial angular momentum and most importantly on the structure near the core and the amount of entropy generated in the interaction with the core."
 I ids the purpose of this paper ο systematically explore how the penetration depends on hese various factors and to determine the properties of the »ost-impact material., It is the purpose of this paper to systematically explore how the penetration depends on these various factors and to determine the properties of the post-impact material.
 In. a subsequent. paper. we will use hese results to model the slow merger of two massive stars in the context of the progenitor of SN 1987. whose highly anomalous properties are likely the result of such a merger (sce Pocdsiadlowski 1992. 1997 for discussion and references).," In a subsequent paper, we will use these results to model the slow merger of two massive stars in the context of the progenitor of SN 1987A, whose highly anomalous properties are likely the result of such a merger (see Podsiadlowski 1992, 1997 for discussion and references)."
 While the hyedrodsnamical simulations presented. in this ober were designed with this application in mind. our analysis is sullicientIv. general to allow application to other systems that have been suggested to be post-merger objects (c.g. V Lvedrac Ixahane et 1996]: FIX Comae stars e.g. tucinski 1990: Welty Ramsey 1994].," While the hydrodynamical simulations presented in this paper were designed with this application in mind, our analysis is sufficiently general to allow application to other systems that have been suggested to be post-merger objects (e.g. V Hydrae [Kahane et 1996]; FK Comae stars [e.g. Rucinski 1990; Welty Ramsey 1994])."
 In Section 2. we outline the basic assumptions for modelling the initial properties. of the stream leaving he mass-losing secondary., In Section \ref{streaml1} we outline the basic assumptions for modelling the initial properties of the stream leaving the mass-losing secondary.
 In Section 3 we use this model to obtain a description of the stream at the star of the hydrodynamical simulations., In Section \ref{streambal} we use this model to obtain a description of the stream at the start of the hydrodynamical simulations.
 Section 4— provides a description of the numerical method. used in. the hvedrodsnamical calculations., Section \ref{method} provides a description of the numerical method used in the hydrodynamical calculations.
 In Section 5 we presen the results of detailed hyelrodyvnamical sinaulations auk systematically cliscuss the physics of the interaction between the jet and the ambient matter. in particular the mechanism of the stream dissipation. the elfects of rotation and. of nuclear burning in the stream.," In Section \ref{streamhydro} we present the results of detailed hydrodynamical simulations and systematically discuss the physics of the interaction between the jet and the ambient matter, in particular the mechanism of the stream dissipation, the effects of rotation and of nuclear burning in the stream."
 Finally. in Section 6 we present an semi-empirical prescription which allows to estimate the. penetration. depth. taking into account the entropy. eenerated in the interaction.," Finally, in Section 6 we present an semi-empirical prescription which allows to estimate the penetration depth, taking into account the entropy generated in the interaction."
 Once the immersed. secondary starts to fill its tical lobe. it starts to transfer mass to the core of the giant.," Once the immersed secondary starts to fill its tidal lobe, it starts to transfer mass to the core of the giant."
 The nmiass-transfoer rate is determined by the friction. between the spiraling-in immersed. binary and the envelope which drives the evolution of the system., The mass-transfer rate is determined by the friction between the spiraling-in immersed binary and the envelope which drives the evolution of the system.
 Since the mass-Ioss time-scale is generally much shorter than the nuclear or thermal time-scale for a main-sequence star. the radius responds acliabatically to mass loss. i.e. inOlnHs) Bσα where A5 and Ads areMs theJ radius and the mass ofο the mass-losing secondary. respectively.," Since the mass-loss time-scale is generally much shorter than the nuclear or thermal time-scale for a main-sequence star, the radius responds adiabatically to mass loss, i.e. = ( =, where $R_2$ and $M_2$ are the radius and the mass of the mass-losing secondary, respectively."
" The Roche-lobe radius RpL of the secondary changes according to where the Roche-lobe radius Art is given by Egeleton (1983) = Llere q=Mo/My is the mass ratio. and AZ, is the total mass of the primary core within the binary separation £2."," The Roche-lobe radius $R_{\rm RL}$ of the secondary changes according to where the Roche-lobe radius $R_{\rm RL }$ is given by Eggleton (1983) = Here $q=M_2/M_1$ is the mass ratio, and $M_1$ is the total mass of the primary core within the binary separation $D$."
 Assuming that the radius of the star is equal to the ltoche-Iobe radius. we can equate equations (1) and (2) and obtain an estimate for the mass-Ioss rate from the secondary as where 0InD/Of is the change in the orbital separation due to tidal and viscous friction.," Assuming that the radius of the star is equal to the Roche-lobe radius, we can equate equations (1) and (2) and obtain an estimate for the mass-loss rate from the secondary as = (, where $\partial \ln D/\partial t$ is the change in the orbital separation due to tidal and viscous friction."
 For typical conditions in a ssupergiant. the resulting mass-loss rates are typically in the range of 0.01 to tUvanova. Podsiadlowski Spruit 20001: and lvanova Podsiaclowski 2002 1Η).," For typical conditions in a supergiant, the resulting mass-loss rates are typically in the range of 0.01 to (Ivanova, Podsiadlowski Spruit 20001; and Ivanova Podsiadlowski 2002 [IP])."
 ‘To obtain values for G4. we performed: separate stellar-evolution caleulations where mass was taken olf the assumed secondary on a time-scale much shorter than the thermal time-scale in its outer envelope (see c.g. Podsiacdlowski. Rappaport Pfahl 2001).," To obtain values for $\zeta_{\rm ad }$, we performed separate stellar-evolution calculations where mass was taken off the assumed secondary on a time-scale much shorter than the thermal time-scale in its outer envelope (see e.g. Podsiadlowski, Rappaport Pfahl 2001)."
 In previous treatments of mass outllow from Lj. it is eencrally assumed that the stream is isothermal (see e.g. Lubow Shu 1975: LS: Bisikalo et 11997).," In previous treatments of mass outflow from $L_1$, it is generally assumed that the stream is isothermal (see e.g. Lubow Shu 1975; LS; Bisikalo et 1997)."
" )ecause of the high expected: mass-loss rate in our problem. it is more appropriate to assume that the stream behaves adiabaticallv. since the radiative cillusion time-scale. is ecnerallv much longer than the characteristic How time-scale (for a stream with a racius of 10! em. the dilfusion time is of order 10"" s)."," Because of the high expected mass-loss rate in our problem, it is more appropriate to assume that the stream behaves adiabatically, since the radiative diffusion time-scale is generally much longer than the characteristic flow time-scale (for a stream with a radius of $10^{10}\,$ cm, the diffusion time is of order $10^6\,$ s)."
 In the case where the secondary's spin is aligned. and svnchronous with the orbit. we can determine the density," In the case where the secondary's spin is aligned and synchronous with the orbit, we can determine the density"
The frst observational evidence of he cosinie infrared backgratnd (CITB) was reported by ? and confirmed by 7. and T.,The first observational evidence of the cosmic infrared background (CIB) was reported by \citet{1996A&A...308L...5P} and confirmed by \citet{1998ApJ...508..123F} and \citet{1998ApJ...508...25H}.
 The €ID is composed of the relic cussion at infrared waveleασ» of the formation and evolution of galaxies axl «κ.usists of contributions from infrared stilirst ealaxies and to a lesser degree from active ealactic clei., The CIB is composed of the relic emission at infrared wavelengths of the formation and evolution of galaxies and consists of contributions from infrared starburst galaxies and to a lesser degree from active galactic nuclei.
 Deep cosimoloegica survevs of this backgratnd have been carried out with ISO (see??.foryeviews) lialli voat 15 jun with ISOCAN (ce...λα at 90 and |TO ju wit1 ISOPIOT (eQe ο] wit1 Spitzer at 21. TO. alic 160 pu in (e.g..2?) al with eronid-based instrunneuts SCUBA (c.g.2). LABOCA (ce.T) aux MAMBO (c.¢..? a SHO. STU. and 1200 sau res)eetively.," Deep cosmological surveys of this background have been carried out with ISO \citep [see] [for reviews] {2000ARA&A..38..761G,2005SSRv..119...93E}
 mainly at 15 $\mu$ m with ISOCAM \citep [e.g.,] [] {2002A&A...384..848E}; at 90 and 170 $\mu$ m with ISOPHOT \citep [e.g.,] [] {2001A&A...372..364D}; ; with Spitzer at 24, 70, and 160 $\mu$ m \citep [e.g.,] []
{2004ApJS..154...70P,2004ApJS..154...87D} and with ground-based instruments SCUBA \citep [e.g.,] [] {2002PhR...369..111B}, LABOCA \citep [e.g.,][]{2008A&A...485..645B} ,and MAMBO \citep [e.g.,] []{2000astro.ph.10553B} at 850, 870, and 1300 $\mu$ m respectively."
" The xilloon-bo1e experiment BLAST perform he firs deep extragalactiο Ναον at wavelereths 250-5004 capade of micasig large mulxu of star-fornud ealaxies. and t10]"" contributions to the CIB (?).. "," The balloon-borne experiment BLAST performed the first deep extragalactic surveys at wavelengths $\mu$ m capable of measuring large numbers of star-forming galaxies, and their contributions to the CIB \cite[][]{2009Natur.458..737D}. ."
These survevs allowec us to obtain a far clea‘er maderstanudiES1g of the CIB aud is sotrces (see7.oragenerareview) but many qiCslons renan WaΙΟ such as the evolution of their spatial «listributio1 with redshift.," These surveys allowed us to obtain a far clearer understanding of the CIB and its sources \citep [see][for a general review]
{2005ARA&A..43..727L} but many questions remain unanswered such as the evolution of their spatial distribution with redshift."
 The spatia distribuion of inlare ealaxi as functio1 of redsuft is a LON COIL)Cnent of the seenario of galaxv formation arc evolution., The spatial distribution of infrared galaxies as a function of redshift is a key component of the scenario of galaxy formation and evolution.
 Tlowever. ifs study has been lamved bv hieh confusion and instruucutal noise andor by he πια] size of the fields of observation.," However, its study has been hampered by high confusion and instrumental noise and/or by the small size of the fields of observation."
 Tentative stiles. with a s«nall nuuber of sources at SHO jan (?).. oud ovkence of a relationship between stbuulnueter ealaxies and the kDynation of Ηνη. ealaxies iu dense environnents.," Tentative studies, with a small number of sources at 850 $\mu$ m \citep{2004ApJ...611..725B}, , found evidence of a relationship between submillimeter galaxies and the formation of massive galaxies in dense environments."
 VVος by 7? and 7? niensured a stong cluscringe of ultra lines mfrared ealaxies (ULIRG) deteced wih Spitzer a lügl redshifts.," Works by \citet{2006ApJ...641L..17F} and \citet{2008MNRAS.383.1131M}
 measured a strong clustering of ultra luminous infrared galaxies (ULIRG) detected with Spitzer at high redshifts."
" Alternatively. he infrared backeroun anisotropics could also provide informatio1 about the correlatio1j between the sources of the CIB and «ark iiater (222), ancl its redshift evolution."," Alternatively, the infrared background anisotropies could also provide information about the correlation between the sources of the CIB and dark matter \citep{2000ApJ...530..124H,2001ApJ...550....7K,2007ApJ...670..903A}, and its redshift evolution."
 7T and ? reported tιο detection of a correated couponent in the background anisotropies Wine Svitzer3IPS (160 jn) and BLAST (250. 350. and 500 ;nu ) data.," \citet{2007ApJ...665L..89L} and \citet{2009arXiv0904.1200V} reported the detection of a correlated component in the background anisotropies using Spitzer/MIPS (160 $\mu$ m) and BLAST (250, 350, and 500 $\mu$ m) data."
 These auhors found that star formation is lüehv] uased at z20.8., These authors found that star formation is highly biased at $>$ 0.8.
 The strong evolution of the bias paraneer with redshift. caused by the shifting of star formaticni fo niore lassive halos with increasing redsliift. infers fmat environnmenal effects influence the vigorous star To inprove onr muerstandiuge of the formation and evohtio1 of galaxies 1sing CIB anisotropies. we need more information about tjo. redsuft of the sources coutribitine to the CTD.," The strong evolution of the bias parameter with redshift, caused by the shifting of star formation to more massive halos with increasing redshift, infers that environmental effects influence the vigorous star To improve our understanding of the formation and evolution of galaxies using CIB anisotropies, we need more information about the redshift of the sources contributing to the CIB."
 We also weed a niethiod. that allows to go deeper than fi6 Confusion nolse level., We also need a method that allows to go deeper than the confusion noise level.
 Iu this conext. an invaluae tevol is the stacking echuique. which alows a statistic:d stidy of groups of sources that cannot f detected imuivichally at a eiven waveleneth.," In this context, an invaluable tool is the stacking technique, which allows a statistical study of groups of sources that cannot be detected individually at a given wavelength."
 Its reqires the knowlec5ο ο the positions of f10 Νοον Cle v'sacked™ as afered from their iudividual detection at alo wavelonetl., Its requires the knowledge of the positions of the sources being “stacked” as inferred from their individual detection at another wavelength.
 This knowledge is then used o stack tie signal of the sources at the wavelength at whic1 they cannot be detected individually., This knowledge is then used to stack the signal of the sources at the wavelength at which they cannot be detected individually.
 Since hne sigwal of the somees dnercascs with the muiuber of SOTLECCS N aud thenoise (if Ciaussian) mereases with VN. he sigjial-to-noise ratk) will increase with WN.," Since the signal of the sources increases with the number of sources N and thenoise (if Gaussian) increases with $\sqrt{N}$, the signal-to-noise ratio will increase with $\sqrt{N}$."
 For an adclitioal description of the basics of stacking echniques we refe| to for example ? and ?.., For an additional description of the basics of stacking techniques we refer to for example \citet[][]{2006A&A...451..417D} and \citet[][]{2009arXiv0904.1205M}.
 Stacking was used oO ineasure the contribution of 2| jon galaxies to the backgronud at 70 aid 160 pan using MOPS data (?).., Stacking was used to measure the contribution of 24 $\mu$ m galaxies to the background at 70 and 160 $\mu$ m using MIPS data \citep{2006A&A...451..417D}.
 Contribution from galaxies down to 60 dy at 21 4n is at leas of the 21 ju. and of he 70 and 160 jan backerounds. respecively.," Contribution from galaxies down to 60 $\mu Jy$ at 24 $\mu$ m is at least of the 24 $\mu$ m, and of the 70 and 160 $\mu$ m backgrounds, respectively."
 At longer wavelengths studies used this techuiqiο to deteruiue he contribution of poolatious selected in the Liear- and midàufrared to the FIRB (far-intraed backgronnud) vackeronnd: 3.60 pan selected sources ο the sho jad vackeronnd (7) and 8 san aud 21 san seected sourcesto he 850 yan and 150 san backgrounds (77)..," At longer wavelengths studies used this technique to determine the contribution of populations selected in the near- and mid-infrared to the FIRB (far-infrared background) background: 3.6 $\mu m$ selected sources to the 850 $\mu$ m background \citep{2006ApJ...647...74W}
 and 8 $\mu$ m and 24 $\mu$ m selected sourcesto the 850 $\mu$ m and 450$\mu$ m backgrounds \citep{2006ApJ...644..769D, 2008MNRAS.386.1907S}. ."
 Finally. ? neasnred total biuillineter mtensifies associatedwitl," Finally, \citet[][]{2009arXiv0904.1205M} measured total submillimeter intensities associatedwith"
 Finally. ? neasnred total biuillineter mtensifies associatedwitli," Finally, \citet[][]{2009arXiv0904.1205M} measured total submillimeter intensities associatedwith"
6397 they found only upper limits. higher (therefore consistent) than our measurements.,"6397 they found only upper limits, higher (therefore consistent) than our measurements."
 As far as the NGC 6752 dwarfs are concerned. they quote a similar value to ours for star 200613. the value 10 times higher quoted for star 4428 in that paper. should instead be considered as an upper limit (E. Carretta private communication. see also Pasquini et al.," As far as the NGC 6752 dwarfs are concerned, they quote a similar value to ours for star 200613, the value 10 times higher quoted for star 4428 in that paper, should instead be considered as an upper limit (E. Carretta private communication, see also Pasquini et al."
 2007)., 2007).
 Our N abundances for GC stars can be consistently compared to the results for field stars of Ecuvillon et al. (, Our N abundances for GC stars can be consistently compared to the results for field stars of Ecuvillon et al. (
2004) and Israelian et al. (,2004) and Israelian et al. (
2004).,2004).
 Figure 3. shows the [N/H| abundance vs. [Fe/H] for field and cluster stars., Figure \ref{fignfe} shows the [N/H] abundance vs. [Fe/H] for field and cluster stars.
 Most field stars closely follow a ΕΙ relationship. with a scatter compatible with the error measurements (quoted as ~+ 0.15 by Israelian et al.," Most field stars closely follow a 1:1 relationship, with a scatter compatible with the error measurements (quoted as $\sim \pm$ 0.15 by Israelian et al."
 2004 and given as a reference in the upper left corner of the Figure)., 2004 and given as a reference in the upper left corner of the Figure).
 The [N/O] ratio is shown in Fig. 4..," The [N/O] ratio is shown in Fig. \ref{figno},"
 for the present discussion we ignore the possible complications arising from the fact that the O abundances displayed have been derived in an inhomogeneous way and using different O indicators. and take all measurements at face value.," for the present discussion we ignore the possible complications arising from the fact that the O abundances displayed have been derived in an inhomogeneous way and using different O indicators, and take all measurements at face value."
 With the that any systematic differences between the different analysis will tend to blur any existing pattern., With the that any systematic differences between the different analysis will tend to blur any existing pattern.
 Two points emerge from Figure 2: The first point is somewhat expected. since we have known for long time of the presence of chemical anomalies in GC stars. which includes the presence of N rich stars.," Two points emerge from Figure 2: The first point is somewhat expected, since we have known for long time of the presence of chemical anomalies in GC stars, which includes the presence of N rich stars."
 N rich stars. though know? in the field. are extremely rare. and a typical sample. like the one of Israelian et al. (," N rich stars, though known in the field, are extremely rare, and a typical sample, like the one of Israelian et al. ("
2004). contains at most one such object.,"2004), contains at most one such object."
" The two NGC 6752 stars show a large spread in N (almost a factor of 20) and the cluster stars. including all stars observed in NGC 6397, have a very high N abundance. clearly departing from the field stars relationship at high significance."," The two NGC 6752 stars show a large spread in N (almost a factor of 20) and the cluster stars, including all stars observed in NGC 6397, have a very high N abundance, clearly departing from the field stars relationship at high significance."
 This is clearly related to the peculiar formation mechanism of GCs. and we will come back on this point in the next section.," This is clearly related to the peculiar formation mechanism of GCs, and we will come back on this point in the next section."
 As far as the second point is concerned we believe that this aspect should be considered with some attention. since it has not been explored too much so far.," As far as the second point is concerned we believe that this aspect should be considered with some attention, since it has not been explored too much so far."
 For this reason we look at the two clusters separately., For this reason we look at the two clusters separately.
 NGC 6752 is a cluster showing all possible chemical inhomogeneities and anticorrelations., NGC 6752 is a cluster showing all possible chemical inhomogeneities and anticorrelations.
 But. in its anomalies. the observed trends are extremely consistent.," But, in its anomalies, the observed trends are extremely consistent."
 The difference of a factor - 20 in N among the two stars is very high. but it ts in line with the other large differences recorded for them in Na (factor 10) and Li (factor 2.5 ) in these stars and in oxygen (factor 4) in other stars having abundance pattern similar to star 200613 (in 200613 O has not been measured).," The difference of a factor $\sim$ 20 in N among the two stars is very high, but it is in line with the other large differences recorded for them in Na (factor 10) and Li (factor 2.5 ) in these stars and in oxygen (factor 4) in other stars having abundance pattern similar to star 200613 (in 200613 O has not been measured)."
 With Our Iew value. the N of star 4428 is higher. but formally compatible with the field stars having similar metallicity.," With our new value, the N of star 4428 is higher, but formally compatible with the field stars having similar metallicity."
" Even if we do not know how this GC formed. the observations are compatible with the idea that star 4428 15 the prototype of an ""unpolluted' star. perfectly similar in all aspects to its field counterparts of similar [Fe/H]. while star 200613 is. at the opposite. the prototype of highly polluted star. with high N. high Na. low Li and low O. showing therefore the signature of gas which experienced high temperatures where nuclear cycles such as CNO. Ne-Na. and the consequent Li destrtction. have occurred."," Even if we do not know how this GC formed, the observations are compatible with the idea that star 4428 is the prototype of an 'unpolluted' star, perfectly similar in all aspects to its field counterparts of similar [Fe/H], while star 200613 is, at the opposite, the prototype of highly polluted star, with high N, high Na, low Li and low O, showing therefore the signature of gas which experienced high temperatures where nuclear cycles such as CNO, Ne-Na, and the consequent Li destruction, have occurred."
 When plotting N vs. O (Fig. 4)).," When plotting N vs. O (Fig. \ref{figno}) ),"
 the situation is similar. unfortunately we have O abundances only for 3 out of five stars.," the situation is similar, unfortunately we have O abundances only for 3 out of five stars."
 The impression is that the cluster stars (including star 4428) stand definitely out of the field relationship., The impression is that the cluster stars (including star 4428) stand definitely out of the field relationship.
 This may be related to the fact that star 4428 has an O/Fe ratio which ts lower. by roughly a factor of 2. than field stars of the same [Fe/H]. when N ts plotted against [Fe/H] star 4428 follows the trend of field stars. but when plotted against oxygen 1t appears to have too much nitrogen. with respect to field stars with the same [O/H].," This may be related to the fact that star 4428 has an O/Fe ratio which is lower, by roughly a factor of 2, than field stars of the same [Fe/H], when N is plotted against [Fe/H] star 4428 follows the trend of field stars, but when plotted against oxygen it appears to have too much nitrogen, with respect to field stars with the same [O/H]."
" Clearly in a cluster like NGC 6752. which has a very strong Na- anticorrelation. it is difficult to tell which is the ""unpolluted"" oxygen content of the cluster."," Clearly in a cluster like NGC 6752, which has a very strong Na-O anticorrelation, it is difficult to tell which is the “unpolluted” oxygen content of the cluster."
 Among the stars analysed by Gratton et al. (, Among the stars analysed by Gratton et al. (
2001) there is one dwarf and two subgiants which have higher [O/H] than star 4428.,2001) there is one dwarf and two subgiants which have higher [O/H] than star 4428.
 We have chosen star, We have chosen star
a function of the real-space separation along the s direction. re.,"a function of the real-space separation along the $\pi$ direction, $r_z$."
 phis can be derived. from the equation of conservation of particle pairs. which are within a comoving separation r from a mass particle. i.e. (Peebles 1980) where (0) is the scale factor.," This can be derived from the equation of conservation of particle pairs, which are within a comoving separation $r$ from a mass particle, i.e. (Peebles 1980) where $a(t)$ is the scale factor."
 Assuming that £(r) is described by a power law model. we solve the above equation to find the infall velocity of the particles We now modify equation 23. to include the cllects of the bulk motions ‘This is the model Ll £(o.x): we then follow the same implementation of the 7XIcock-DPaczvnski οσοι and fitting procedure as for model LAs in model L. we keep 5; constant. using the same value as above (loss et al.," Assuming that $\xi (r)$ is described by a power law model, we solve the above equation to find the infall velocity of the particles We now modify equation \ref{eqn:sigma2} to include the effects of the bulk motions This is the model II $\xi (\sigma ,\pi )$; we then follow the same implementation of the “Alcock-Paczynski” effect and fitting procedure as for model I. As in model I, we keep $b_L$ constant, using the same value as above (Ross et al."
 2007). Le. τε0.05 and bp=1.66d:0.35.," 2007), i.e. $\beta _L=0.45\pm0.05$ and $b_L=1.66\pm0.35$."
 We now use our £(o.x) measurements from the previous Section and the sy and * values from the €(s) fits shown in Table 4. to put constraints on 49 and bo and G2)77. for each of the QSO samples.," We now use our $\xi (\sigma, \pi)$ measurements from the previous Section and the $s_0$ and $\gamma$ values from the $\xi (s)$ fits shown in Table \ref{table:spectroscopic} to put constraints on $\beta_Q$ and $b_Q$ and $\langle \omega_z^2\rangle ^{1/2}$ for each of the QSO samples."
 The results are shown in Table and in Figures 1H4--19.., The results are shown in Table \ref{Table:beta_b_redz} and in Figures \ref{fig:2slaq_1}- \ref{fig:sdss_2}.
 Comparing the results from both models. from. the three different QSO samples associated with the same spectroscopic 2SLAQ LAG sample. we note that mocel ] gives slightly. lower QSO-LRG velocity dispersions than model LH. but. both are consistent. (~620kms+ and ~TOYkms *) with the expected whe eq2skns|. produced by adding in quadrupole QSO and LRG velocity dispersions from previous QSO0-OSO and LRG-LAG studies. which gave  ancl 7. respectively.," Comparing the results from both models from the three different QSO samples associated with the same spectroscopic 2SLAQ LRG sample, we note that model I gives slightly lower QSO-LRG velocity dispersions than model II, but both are consistent $\sim$ $^{-1}$ and $\sim727$ $^{-1}$ ) with the expected $\langle \omega_z^2\rangle ^{1/2}\simeq$ $^{-1}$, produced by adding in quadrupole QSO and LRG velocity dispersions from previous QSO-QSO and LRG-LRG studies, which gave $^{-1}$ and $^{-1}$, respectively."
 Comparison of bo between the cdillerent samples (for both moclels) shows that the results are in good statistical agreement (‘Table 7))., Comparison of $b_Q$ between the different samples (for both models) shows that the results are in good statistical agreement (Table \ref{Table:beta_b_redz}) ).
 Phe best way to summoarise the results is to average them., The best way to summarise the results is to average them.
 All six measurements (the three samples and the two models) give an average of 39=04550.10 and bo=1.420.2 at z—0.55. which is consistent with the values found by da -lagela et al. (," All six measurements (the three samples and the two models) give an average of $\beta _Q=0.55\pm0.10$ and $b_Q=1.4\pm0.2$ at z=0.55, which is consistent with the values found by da $\hat{A}$ ngela et al. ("
2008). 39=0.60.tit and bg=1540.2 at z214 and bo=LIτὸ at 20.6 (see their Fig.,"2008), $\beta _Q=0.60_{-0.11}^{+0.14}$ and $b_Q=1.5\pm0.2$ at z=1.4 and $b_Q=1.1\pm0.2$ at $z\simeq0.6$ (see their Fig."
 13)., 13).
 Previous attempts to study the dependence of clustering on luminosity. were not successful because of the recshilt-luminosity degeneracy. as higher. luminosity QSOs resice at higher redshifts.," Previous attempts to study the dependence of clustering on luminosity, were not successful because of the redshift-luminosity degeneracy, as higher luminosity QSOs reside at higher redshifts."
 Here. combining QSOs with large LRG samples we get the statistical power to break this clegeneracy.," Here, combining QSOs with large LRG samples we get the statistical power to break this degeneracy."
 Lathis Section. we shall trey to examine if and how the QSO-LRG clustering depends on QSO luminosity. at fixed redshift.," Inthis Section, we shall try to examine if and how the QSO-LRG clustering depends on QSO luminosity, at fixed redshift."
" We first estimate the average absolute magnitude. of each of our QSO samples. as follows: where Mi, is the absolute magnitude of each QSO. by (or g) is its apparent magnitude. d; is the luminosity distance (Alpe) that corresponds to the redshift z and the last term is the k-correction. where we have assumed a QSO spectral index a=1.0."," We first estimate the average absolute magnitude of each of our QSO samples, as follows: where $M_{b_J}$ is the absolute magnitude of each QSO, $b_J$ (or $g$ ) is its apparent magnitude, $d_L$ is the luminosity distance (Mpc) that corresponds to the redshift $z$ and the last term is the k-correction, where we have assumed a QSO spectral index $\alpha^\prime=-1.0$ ."
 We should note here that we treat the 56; band (2QZ) as being equivalent to the g band (PSLAQ. SDSS: Richares ct al.," We should note here that we treat the $b_J$ band (2QZ) as being equivalent to the $g$ band (2SLAQ, SDSS; Richards et al."
 2005)., 2005).
 To check if there is a dependence of QSO-LRG clustering on luminosity we use the integrated correlation function. as a more robust statistical tool (see cla Angela et al.," To check if there is a dependence of QSO-LRG clustering on luminosity we use the integrated correlation function, as a more robust statistical tool (see da $\hat{A}$ ngela et al."
 2008)., 2008).
 We calculate it up to scales of Alpe and normalise the result. to the volume contained in a sphere with radius of !Mpe: The choice of the radius has been mace for two reasons., We calculate it up to scales of $^{-1}$ Mpc and normalise the result to the volume contained in a sphere with radius of $^{-1}$ Mpc: The choice of the radius has been made for two reasons.
 The ! Mpe scale is large enough to apply linear theory (Croom ct al., The $^{-1}$ Mpc scale is large enough to apply linear theory (Croom et al.
 2005). ancl redshift-space distortions (finger of god. or redshift errors) do not significantly allect the measurements.," 2005), and redshift-space distortions (finger of god or redshift errors) do not significantly affect the measurements."
 In the case of the QSOs-25LAQ (spectroscopic) LRGs we have estimated £25 via equation 28 using the ές] measurements shown in bie. 7.., In the case of the QSOs-2SLAQ (spectroscopic) LRGs we have estimated $\xi _{20}$ via equation \ref{eqn:integ_xis} using the $\xi (s)$ measurements shown in Fig. \ref{fig:xis_all1}.
 For the QSO-28LAQ (photometric) and ANOmega LRGs we have substituted the £(5s) in equation 28. with the £r) fits shown in Fig. S.., For the QSO-2SLAQ (photometric) and AAOmega LRGs we have substituted the $\xi (s)$ in equation \ref{eqn:integ_xis} with the $\xi (r)$ fits shown in Fig. \ref{fig:Limber_all}.
 The results using 25LAQ LRGs (spectroscopic and photometric) are shown in Fig., The results using 2SLAQ LRGs (spectroscopic and photometric) are shown in Fig.
 20. and using AANOmeea LCs are shown in Fig. 21.., \ref{fig:xi_bar_2slaq} and using AAOmega LRGs are shown in Fig. \ref{fig:xi_bar_aao}.
 No conclusion can be drawn about the recdshift evolution of the QSO-LI clustering as the average redshift of the samples is too restricted., No conclusion can be drawn about the redshift evolution of the QSO-LRG clustering as the average redshift of the samples is too restricted.
 Although the £oo. results using the AAOmeea LRGs. show some indications that bright QSOs (SDSS) cluster less with LRGs than faint QSOs (2PSLAQ). the results using the 25LAQ LAG samples (both spectroscopic and photometric) stay statistically constant. thus confirming the results in the previous Sections. that the QSO-LRG clustering is independent of QSO luminosity.," Although the $\xi _{20}$ results using the AAOmega LRGs, show some indications that bright QSOs (SDSS) cluster less with LRGs than faint QSOs (2SLAQ), the results using the 2SLAQ LRG samples (both spectroscopic and photometric) stay statistically constant, thus confirming the results in the previous Sections, that the QSO-LRG clustering is independent of QSO luminosity."
 Lavine tested. the luminositw dependence of the QSO-LRG clustering. we now investigate the dependence of the QSO bias on the [uminosity.," Having tested the luminosity dependence of the QSO-LRG clustering, we now investigate the dependence of the QSO bias on the luminosity."
" Assuming that this bias is independent of scale. we calculate the QSO bias using: where £, is the matter real-space correlation function. averaged in !Mpe spheres."," Assuming that this bias is independent of scale, we calculate the QSO bias using: where $\xi _{\rho}$ is the matter real-space correlation function, averaged in $^{-1}$ Mpc spheres."
 In the. case. of QSO-Whotometric LRGs we use our é£(r(r.20) measurements as estimated before.," In the case of QSO-photometric LRGs we use our $\xi (r,20)$ measurements as estimated before."
" For consistency. for the QSO-28LAQ (spectroscopic) LRGs. we use our dwe,(o0) measurements (shown in Table 9)) since they are [ess noisy than those or £(r)."," For consistency, for the QSO-2SLAQ (spectroscopic) LRGs, we use our $w_p(\sigma)$ measurements (shown in Table \ref{table:biases_qso}) ) since they are less noisy than those for $\xi (r)$ ."
" To estimate £,,, we use the valuesas estimated w da «ποσα et al.", To estimate $\xi _{mm}$ we use the valuesas estimated by da $\hat{A}$ ngela et al.
 2008., 2008.
" So in the case of QSO-AAOmega LRGs £,,,=QO.1l at εςτε 0.7). and in the case of QSO- LAGS £,,=0.12 (22 0.55)."," So in the case of QSO-AAOmega LRGs $\xi _{mm}=0.11$ at $z\approx0.7$ ), and in the case of QSO-2SLAQ LRGs $\xi _{mm}=0.12$ $z\approx0.55$ )."
 Finally. we calculate he LRG bias for the ANOmega and 25LAQ. based on the," Finally, we calculate the LRG bias for the AAOmega and 2SLAQ, based on the"
jenettin ct al (1976)).,Benettin et al \citeyear{Bennetin}) ).
" Thus by fixing the motion integrals as f=1245 and (= 0.2. choosing A,c10) and |= 10. we estimate No corresponding to a typical chaotic disc-crossing orbit for the cases m=1.2.3.4 (Vable 1))."," Thus by fixing the motion integrals as $E=-1.245$ and $\ell=0.2$ , choosing $\Delta_{o}\simeq 10^{-9}$ and $t=10^{7}$ , we estimate $N$ corresponding to a typical chaotic disc-crossing orbit for the cases $m=1,2,3,4$ (Table \ref{Table}) )."
 We ound that the degree of instability increases modestly with m., We found that the degree of instability increases modestly with $m$.
 An interesting phenomenon occurs when the effective »otential has saddle points outside the disc., An interesting phenomenon occurs when the effective potential has saddle points outside the disc.
 This happens for he m=1 and m=2 disces. in whose case there are critical »»unuts in the ranges 1«A.LOS and 1«A.x1015. respectively.," This happens for the $m=1$ and $m=2$ discs, in whose case there are critical points in the ranges $1< R_{c}\leq 1.198$ and $1< R_{c}\leq
1.075$, respectively."
 Such equilibrium: points are outside the disc out near its edge. so for certain values of IZ the contourof," Such equilibrium points are outside the disc but near its edge, so for certain values of $E$ the contourof"
values for [τν from Bloomctal.(2001).,values for $E_{iso}$ from \citet{bfs01}.
. The fit elves (Fig.⋅ ⋅ 6)), The fit gives (Fig. \ref{Eisovtj}) )
 with. qz5=51/1. aud a gooducss-of-üt ο~1019.," with $\chi^2_r = 51/4$ and a goodness-of-fit $Q
\sim 10^{-10}$."
 As is also the case for the relation shown in Equ. (2))," As is also the case for the relation shown in Eqn. \ref{Leqn}) ),"
 the quality of the fit is poor. but the trend that bursts of higher ΟΠΟΙΟΥ have shorter jet-break timescales is plainly evident.," the quality of the fit is poor, but the trend that bursts of higher energy have shorter jet-break timescales is plainly evident."
 This relation allows one to test the sugeestion by Frailetal.(2001) that there is a reservoir for gamunarav energv. £.. that is constant for all GRBs.," This relation allows one to test the suggestion by \citet{fksd01} that there is a reservoir for gamma-ray energy, $E_\gamma$, that is constant for all GRBs."
" For bursts with a large jet-openig angle. 0;. this eucrgy would be spread out. resultingin lower fiuxes and lower interred isotropic cuereies, Lj."," For bursts with a large jet-opening angle, $\theta_j$, this energy would be spread out, resultingin lower fluxes and lower inferred isotropic energies, $E_{iso}$."
 This suggests E.& constant = EUxUEsurjY," This suggests $E_\gamma \approx$ constant = $E_{iso}
\theta_j^2 \propto (E_{iso} \tau_j)^{3/4}$."
 ‘The assumption that energev in ganunivcrays is constant predicts DL;x/To. Which is not consistent with the data.," The assumption that energy in gamma-rays is constant predicts $E_{iso} \propto \tau_j^{-1}$, which is not consistent with the data."
 Thus. while the range of observed £. ids sienificautly sunaller than that of τω it is not exactly coustaut.," Thus, while the range of observed $E_\gamma$ is significantly smaller than that of $E_{iso}$, it is not exactly constant."
" The existence of these correlatious. in particular he relation between jet-break times auc spectral aes, point to a surprisinely direct relationship vetween the GRB phase (represeuted by the pulse ag and the luminosity) aud the afterglow phase (represented by the jet break time)."," The existence of these correlations, in particular the relation between jet-break times and spectral lags, point to a surprisingly direct relationship between the GRB phase (represented by the pulse lag and the luminosity) and the afterglow phase (represented by the jet break time)."
 In the next section we will argue that these relationships are sinenmatic dn origin., In the next section we will argue that these relationships are kinematic in origin.
 The effectively linear relatiouship between he two timescales: pulse lags At aud jet-break iues 7; (Equ. 1)), The effectively linear relationship between the two timescales; pulse lags $\Delta t$ and jet-break times $\tau_j$ (Eqn. \ref{t_jeqn}) )
 suggests a surprisingly direct connection between the CRB phase. from which he pulse lags ave derived. and the afterglow phase. roni which jet-break times derive.," suggests a surprisingly direct connection between the GRB phase, from which the pulse lags are derived, and the afterglow phase, from which jet-break times derive."
 An initial conclusion that cau be drawn is that. since the GRD phase is generally thought not to depend ou he density of the iuterstellar τοςαι (SAD). this relation disfavors explanations of the jet break deriving from external causes such as sudden drops in ISAL density (INunuar&Panaitescu 2000)..," An initial conclusion that can be drawn is that, since the GRB phase is generally thought not to depend on the density of the interstellar medium (ISM), this relation disfavors explanations of the jet break deriving from external causes such as sudden drops in ISM density \citep{kp00b}."
 But what could be the meaning of such a connection between the GRD aud the afterelow phase. expressed in the relatious preseuted iu Fies. ," But what could be the meaning of such a connection between the GRB and the afterglow phase, expressed in the relations presented in Figs. \ref{tjvtlag},"
L. 5 and 677, \ref{Lvtlag} and \ref{Eisovtj}?
 It was argued by Salimonsou(2000)— that the spectral lag isotropic sganuna-cray peak huninosity relationship (Norrisetal.2000:Salinouson ταν have a simple kinematic explanation.," It was argued by \citet{jay00} that the spectral lag isotropic gamma-ray peak luminosity relationship \citep{nmb00,jay00} may have a simple kinematic explanation."
 Defining a Doppler factor of the ejecta moving with anele 0 from the line of sight. with a velocity §=c/c at a redshift z..a proper GRD timescale r iu the frame of the clutter will be observed as To calculate the lab-frame Iuniuositv we assinaie hat the jet opening angle is greater than the relativistic Deaming angle (0;>1/25).," Defining a Doppler factor of the ejecta moving with angle $\theta$ from the line of sight, with a velocity $\beta \equiv v/c$ at a redshift $z$, a proper GRB timescale $\tau$ in the frame of the emitter will be observed as To calculate the lab-frame luminosity we assume that the jet opening angle is greater than the relativistic beaming angle $(\theta_j >
1/\gamma)$."
 A given xàiotoun spectrmm of£) (photous keVbou ly which is relativistically invariant. will vield a peak uninosity⋅⋅ Ljxfey7Eo(E/DME. (eres 1) or ∙∕↙ ⋅⋅ ⇁↓−∩⋖∑∕1j⋉∖⋜∐↘↽∐⋯⊔↴⋝↸∖↥⋅↕↿∐∐∐∪↴∖↴↕↑⋅↖↽⊀∖∣↗∕⊽⇁∖ (photons when viewed through a lab-frame vandpass [01.62] (e.g..Fenimorectal.1992).," A given photon spectrum $\phi(E)$ (photons $^{-1}$ $^{-1}$ ), which is relativistically invariant, will yield a peak luminosity $L_{pk} \propto \int^{e_2}_{e_1} E \phi(E/{\cal D}) dE$ (ergs $^{-1}$ ) or peak number luminosity $N_{pk} \propto
\int^{e_2}_{e_1}\phi(E/{\cal D}) dE$ (photons $^{-1}$ ) when viewed through a lab-frame bandpass $e_1,e_2$ ] \citep[e.g.,][]{fehk92}."
. An acceptable analytical ft to off) is given as a woken power-law (Bandetal.1993)., An acceptable analytical fit to $\phi(E)$ is given as a broken power-law \citep{bmf+93}.
". However. since most of the photon flux is frou low enereies. we use. for the sake of arguineut. oulv the low eud of the power-law spectruu o(£)xE"". keepiug in nünd that there will be a correction of order log(CEy). which is s2all for our purposes: here. Ey is the break energy in the energy spectiun (Baudetal. 1993)."," However, since most of the photon flux is from low energies, we use, for the sake of argument, only the low end of the power-law spectrum $\phi(E) \propto E^{-\alpha}$, keeping in mind that there will be a correction of order $\log (E_0)$, which is small for our purposes; here $E_0$ is the break energy in the energy spectrum \citep{bmf+93}."
 As such we find that the Iuminosities vary as: Tf spectral pulse-lag is due to some proper decay timescale Ar. this time scale will be At=Az/D iu the lab-frame.," As such we find that the luminosities vary as: If spectral pulse-lag is due to some proper decay timescale $\Delta
\tau$, this time scale will be $\Delta t = \Delta \tau/{\cal D}$ in the lab-frame."
 Observations find az1 (Preece 1998).., Observations find $\alpha \approx 1$ \citep{ppbm+98}. .
 Thus. combining Equs.," Thus, combining Eqns."
 (5 6)). we arene that the inverse relationship of the pulse lag - huninosity correlation discovered by Norris and that discussed by Salinounsou(2000) result from a purely kinematic effect.," \ref{tdeltaeqn} \ref{lumvsdeltaeqn}) ), we argue that the inverse relationship of the pulse lag - luminosity correlation discovered by \citet{nmb00}
 and that discussed by \citet{jay00}
 result from a purely kinematic effect."
 That is. faster expanding bursts will appear to be brighter and to evolve more quickly than slower bursts.," That is, faster expanding bursts will appear to be brighter and to evolve more quickly than slower bursts."
"sl),",$^{-1}$ ).
 The reconstruction satisfies various observational constraints., The reconstruction satisfies various observational constraints.
" Linear polarization disappears at Ar=120 s, when the apex of the loop abandons the formation region of the 630 nm lines."," Linear polarization disappears at $\Delta t=120$ s, when the apex of the loop abandons the formation region of the 630 nm lines."
 Two minutes later (At=210—240 s) some faint signals are observed in the Mg b magnetogram at the same position of the footpoints in the photosphere., Two minutes later $\Delta t=210-240$ s) some faint signals are observed in the Mg b magnetogram at the same position of the footpoints in the photosphere.
" Since longitudinal magnetograms are blind to horizontal fields, we cannot determine when the apex reaches the temperature minimum."," Since longitudinal magnetograms are blind to horizontal fields, we cannot determine when the apex reaches the temperature minimum."
" However, at Ar=180—240 s, a larger portion of the reconstructed loop is predominantly vertical at the temperature minimum height."," However, at $\Delta t=180-240$ s, a larger portion of the reconstructed loop is predominantly vertical at the temperature minimum height."
" At that moment, the observed separation between footpoints in the Mg b magnetogram is 1350 km while the distance between the loop footpoints is between 1000 and 1730 km."," At that moment, the observed separation between footpoints in the Mg b magnetogram is 1350 km while the distance between the loop footpoints is between 1000 and 1730 km."
" At At=750 s, brightenings in the Ca H-line, corresponding to the low chromosphere, are identified co-spatially with the loop footpoints."," At $\Delta t=750$ s, brightenings in the Ca H-line, corresponding to the low chromosphere, are identified co-spatially with the loop footpoints."
" Although brightenings in this line are widely used as a proxy for magnetic activity, we cannot associate their appearance with the loop reaching a given height."," Although brightenings in this line are widely used as a proxy for magnetic activity, we cannot associate their appearance with the loop reaching a given height."
" If brightenings in are produced only when the magnetic field is concentrated, it is possible that the loop could have reached the low chromosphere before At=750 s but we have been unable to detect it."," If brightenings in are produced only when the magnetic field is concentrated, it is possible that the loop could have reached the low chromosphere before $\Delta t= 750$ s but we have been unable to detect it."
 The ascent to the corona remains unconstrained by our observations., The ascent to the corona remains unconstrained by our observations.
 The dynamics of the emergence process can be more complicated., The dynamics of the emergence process can be more complicated.
 The example shown in Figs., The example shown in Figs.
" 2 and 3 is an extended event, rapidly expanding over the whole granule."," \ref{loop2_1} and \ref{loop2_2} is an extended event, rapidly expanding over the whole granule."
 The flux emerges in a preexisting granule as a structure showing a simple bipolar loop with a clear preferred azimuth before developing a full three-dimensional structure and dynamics (Fig. 3))., The flux emerges in a preexisting granule as a structure showing a simple bipolar loop with a clear preferred azimuth before developing a full three-dimensional structure and dynamics (Fig. \ref{loop2_2}) ).
" Then, one of the footpoints suffers strong shear when reaching the intergranular lane and the loop fans out, with the azimuth of individual field lines spanning 60°."," Then, one of the footpoints suffers strong shear when reaching the intergranular lane and the loop fans out, with the azimuth of individual field lines spanning $^\circ$."
 At Ar=330 s the footpoints are detected in the magnetogram coexisting with a (faint) linear polarization signal at the lines from photospheric layers., At $\Delta t=330$ s the footpoints are detected in the magnetogram coexisting with a (faint) linear polarization signal at the lines from photospheric layers.
" This means that, although many field lines have already reached the upper atmosphere. some of them are still in the photosphere (Fig. 2)."," This means that, although many field lines have already reached the upper atmosphere, some of them are still in the photosphere (Fig. \ref{loop2_1}) )."
 At Ar=270 s a dark, At $\Delta t=270$ s a dark
"One of the key results of this work is that the low metal accretion rates onto the WD Mz«Mag, are possible only if the debris disk hosted by the WD has small mass, below Mo given by equation (31)), and is everywhere optically thin to incident stellar radiation.","One of the key results of this work is that the low metal accretion rates onto the WD $\dot M_Z\ll\dot M_\infty$ are possible only if the debris disk hosted by the WD has small mass, below $M_0$ given by equation \ref{eq:M_0}) ), and is everywhere optically thin to incident stellar radiation."
 is demonstrated in 83 where we find Mz5ThistheAGτιία=<Mass (see equation (25))) throughout evolution1) of the low mass disk., This is demonstrated in \ref{sect:low} where we find $\dot M_Z\approx \dot M_\infty\tpar(x=1)\ll \dot M_\infty$ (see equation \ref{eq:sol}) )) throughout the whole evolution of the low mass disk.
" On the contrary, if the disk is massive enough to contain optically thick regions, which typically requires belowMa=107°—10?! g, then Mz is guaranteed to not fall (Rin/Rout)Ma=©0.2M, as we have shown in §4.."," On the contrary, if the disk is massive enough to contain optically thick regions, which typically requires $M_d\gtrsim 10^{20}-10^{21}$ g, then $\dot M_Z$ is guaranteed to not fall below $(R_{in}/R_{out})\dot M_\infty\approx 0.2\dot M_\infty$, as we have shown in \ref{sect:high}. ."
 This is in agreement with the statement in R11 that Mz should not deviate significantly from Mass for optically thick disks., This is in agreement with the statement in R11 that $\dot M_Z$ should not deviate significantly from $\dot M_\infty$ for optically thick disks.
" This lower bound on Mz for massive disks is very robust, in particular it does not depend on the disk mass."," This lower bound on $\dot M_Z$ for massive disks is very robust, in particular it does not depend on the disk mass."
 This point is illustrated in Figure 8 where we show the time evolution of Mz for three disks with the same initial Gaussian density distribution (33)) but with different initial masses Mq set to be in the ratio 0.5:12., This point is illustrated in Figure \ref{fig:mdot3} where we show the time evolution of $\dot M_Z$ for three disks with the same initial Gaussian density distribution \ref{eq:prof}) ) but with different initial masses $M_d$ set to be in the ratio $0.5:1:2$.
" One can easily see that the only thing, which is different between these three cases is the disk lifetime (which is always proportional to the disk mass), while Mz is essentially the same."," One can easily see that the only thing, which is different between these three cases is the disk lifetime (which is always proportional to the disk mass), while $\dot M_Z$ is essentially the same."
 'The time evolution of Mz in the optically thick case is found to depend predominantly on the shape of the initial density distribution., The time evolution of $\dot M_Z$ in the optically thick case is found to depend predominantly on the shape of the initial density distribution.
 Debris distributions in the form of 8 narrow ring tend to produceMz only weakly varying with time., Debris distributions in the form of a narrow ring tend to produce$\dot M_Z$ only weakly varying with time.
" Extended initial distributions of the debris with radial width comparable to mean radius generally lead to Mz varying in time and decaying by a factor of several by the end of the disk lifetime, see Figure 5.."," Extended initial distributions of the debris with radial width comparable to mean radius generally lead to $\dot M_Z$ varying in time and decaying by a factor of several by the end of the disk lifetime, see Figure \ref{fig:mdot2}."
" Based on these results one should expect WDs with strong IR excesses, signifying massive, optically thick debris disk around them,current to exhibit M.zsuggest:0.2M.."," Based on these results one should expect WDs with strong IR excesses, signifying massive, optically thick debris disk around them, to exhibit $\dot M_Z\gtrsim 0.2\dot M_\infty$."
" This is exactly what the data seem to no WD with a debris disk detected via its IR signature is inferred to have Mz below Ms, (R11).", This is exactly what the current data seem to suggest: no WD with a debris disk detected via its IR signature is inferred to have $\dot M_Z$ below $\dot M_\infty$ (R11).
" However, if the systems exhibiting IR excesses were found to show Mz«Moo, the debris disks around them should be optically thin."," However, if the systems exhibiting IR excesses were found to show $\dot M_Z\ll \dot M_\infty$ , the debris disks around them should be optically thin."
" It is then natural to expect a positive correlation between the strength of the IR excess and the Mz in such systems, since both are linearly proportional to 7j<1 in the optically thin regime."," It is then natural to expect a positive correlation between the strength of the IR excess and the $\dot M_Z$ in such systems, since both are linearly proportional to $\tpar\ll 1$ in the optically thin regime."
" Indeed, Mzxτι according to (25)) while the amount of the WD luminosity that the optically thin disk intercepts and reradiates in the near-IR. is ((r)dr, which also scales positively with 7)."," Indeed, $\dot M_Z\propto\tpar$ according to \ref{eq:sol}) ) while the amount of the WD luminosity that the optically thin disk intercepts and reradiates in the near-IR is (r)dr, which also scales positively with $\tpar$ ."
 The timescale on which debris disks get depleted can be quite different in the optically thick and thin cases., The timescale on which debris disks get depleted can be quite different in the optically thick and thin cases.
" Characteristic lifetime of an optically thin disk composed of ~1 cm particles tinin is several 10? yr, as equations (17)) and (26)) demonstrate."," Characteristic lifetime of an optically thin disk composed of $\sim 1$ cm particles $t_{thin}$ is several $10^5$ yr, as equations \ref{eq:t_0_sub}) ) and \ref{eq:T_thin}) ) demonstrate."
" This timescale depends only on the debris particle properties — size and density, being longer for bigger and denser particles — and the outer extent of the disk (for Rout/Rin7:5 one finds disk lifetime approaching a Myr)."," This timescale depends only on the debris particle properties — size and density, being longer for bigger and denser particles — and the outer extent of the disk (for $R_{out}/R_{in}\approx 5$ one finds disk lifetime approaching a Myr)."
" In the case of an optically thick disk the lifetime can be considerably longer since f;4;4; scales linearly with the disk mass, see equation (40))."," In the case of an optically thick disk the lifetime can be considerably longer since $t_{thick}$ scales linearly with the disk mass, see equation \ref{eq:T_thick}) )."
" Thus, massive disks require long time to get exhausted by the PR drag alone."," Thus, massive disks require long time to get exhausted by the PR drag alone."
" For example, an optically thick disk created by the tidal disruption of an asteroid with. the diameter of 300 km would have a mass of 4x107? g (assuming bulk density of 3 g cm?)."," For example, an optically thick disk created by the tidal disruption of an asteroid with the diameter of 300 km would have a mass of $4\times 10^{22}$ g (assuming bulk density of 3 g $^{-3}$ )."
 Accretion at the steady rate Mz—3x10? g s-! (e0.5Mo given by the numerical estimate in equation (9))) would deplete this disk only within 40 Myr., Accretion at the steady rate $\dot M_Z=3\times 10^7$ g $^{-1}$ $\approx 0.5M_\infty$ given by the numerical estimate in equation \ref{eq:mdot_PR}) )) would deplete this disk only within 40 Myr.
" Such massive disks, however, are likely to be affected by the interaction with the gas that is produced by sublimation of solid debris at the inner edge of the disk."," Such massive disks, however, are likely to be affected by the interaction with the gas that is produced by sublimation of solid debris at the inner edge of the disk."
 They may then evolve in a runaway fashion as described in Rafikov (2011b) and get exhausted much earlier., They may then evolve in a runaway fashion as described in Rafikov (2011b) and get exhausted much earlier.
" An interesting finding of this work is the universal appearance of the sharp outer edge in optically thick disks, caused by the faster inward migration of the disk annuli having smaller optical depth."," An interesting finding of this work is the universal appearance of the sharp outer edge in optically thick disks, caused by the faster inward migration of the disk annuli having smaller optical depth."
" This outer edge steadily moves inwards which, combined with the outward expansion of the optically thin tail at small radii, leads to the gradual reduction of the radial width of the optically thick part of the debris disk."," This outer edge steadily moves inwards which, combined with the outward expansion of the optically thin tail at small radii, leads to the gradual reduction of the radial width of the optically thick part of the debris disk."
" Thus, even an extended disk can turn into a narrow ring with timeas can be easily seen in Figure 7.."," Thus, even an extended disk can turn into a narrow ring with timeas can be easily seen in Figure \ref{fig:tophat}. ."
 Observational evidence for sharp outeredges would strongly support the pictureof disk evolution due to the PR drag outlined in this work., Observational evidence for sharp outeredges would strongly support the pictureof disk evolution due to the PR drag outlined in this work.
" Interestingly, Jura (2007b, 2009) have previously claimed the need for"," Interestingly, Jura (2007b, 2009) have previously claimed the need for"
where the umuber of galaxies aud SNe for cach sample is indicated in the first two rows.,where the number of galaxies and SNe for each sample is indicated in the first two rows.
 We aust stress that. since we adopt the same galaxy catalog. mput parameters. bias corrections and πιοΊσα recipe. the differences between the columns in Table 2. are only due to the different logs of observatious.," We must stress that, since we adopt the same galaxy catalog, input parameters, bias corrections and numerical recipe, the differences between the columns in Table \ref{evares} are only due to the different logs of observations."
 The new Evans rates are consistent with the earlier results mit. at least with respect to the average value. in much better agreement with the rate from plotoeraphic searches.," The new Evans rates are consistent with the earlier results but, at least with respect to the average value, in much better agreement with the rate from photographic searches."
 Loosine more in detail. it apyears that for early type galaxies (E-SO and σα) the updated Evans value is iu better agrecment with the photographic scarcl rate. whereas for late spirals (Sbce-Sd) the old value was closer.," Looking more in detail, it appears that for early type galaxies (E-S0 and S0a-Sb) the updated Evans value is in better agreement with the photographic search rate, whereas for late spirals (Sbc-Sd) the old value was closer."
 We attribute these fluctuations to the small statistics of individual SN. searches. which become wider when the saluple is divided iuto bins.," We attribute these fluctuations to the small statistics of individual SN searches, which become wider when the sample is divided into bins."
 As inentioned before. corrections for searcli biases are he most controversial step in the caleulation of SN rates.," As mentioned before, corrections for search biases are the most controversial step in the calculation of SN rates."
 Tn our approach (cf., In our approach (cf.
 C97). the correction factors are tuned o cancel the sign of the biases from the caleulated. SN rates regardless of their physical causes.," C97), the correction factors are tuned to cancel the sign of the biases from the calculated SN rates regardless of their physical causes."
 The ideal would ve to build a model for the dust aud SN distribution which is consistent with the present data on galaxies aud SN xogenitor populatious aud derive from it estiuuates of the ouases., The ideal would be to build a model for the dust and SN distribution which is consistent with the present data on galaxies and SN progenitor populations and derive from it estimates of the biases.
 Iu a receut paper. Hatano e al. (1998))," In a recent paper, Hatano et al. \cite{hatano}) )"
 described a siuple model based on an asstumed distribution of the dust aud SN populations which predicts that. because of extinction. SNe in inclined spirals appear ou average cinuuer aud shoe a much wider magnitude scatter than those iu face-on spirals.," described a simple model based on an assumed distribution of the dust and SN populations which predicts that, because of extinction, SNe in inclined spirals appear on average dimmer and shoe a much wider magnitude scatter than those in face-on spirals."
 Moreover. because the dust distribution peaks iu the central regions of galaxies. this effect is more pronounced for SNe occurring iu those roeglous.," Moreover, because the dust distribution peaks in the central regions of galaxies, this effect is more pronounced for SNe occurring in those regions."
 This provides an alternative to the classical explanation of the selection effect iu the central region of galaxies which would thus derive frou the eulianced extinction of SNe instead of the reduced huuinositv contrast., This provides an alternative to the classical explanation of the selection effect in the central region of galaxies which would thus derive from the enhanced extinction of SNe instead of the reduced luminosity contrast.
 According to this scenario. the bias would be most severe in dust. inclined. spirals aud. because of the «ια scale height. for core collapse SNe.," According to this scenario, the bias would be most severe in dust, inclined spirals and, because of the small scale height, for core collapse SNe."
 Iudoeed. all these features were found in the observed SN sample (Table 1 of Cappellaro Turatto 1997)).," Indeed, all these features were found in the observed SN sample (Table 1 of Cappellaro Turatto \cite{ct:97}) )."
 A special characteristic of the Hatauo et al., A special characteristic of the Hatano et al.
 model is that core collapse SNe do not occur within 3 Ixpc of the ceuter of the galaxy., model is that core collapse SNe do not occur within 3 Kpc of the center of the galaxy.
 Therefore SN II or Toc do not appear in the central reeious of facc-on galaxies although an increasing uuuber of core collapse SNe appears to be projected ou the centers of the more inclined spirals due to projection effects., Therefore SN II or Ib/c do not appear in the central regions of face-on galaxies although an increasing number of core collapse SNe appears to be projected on the centers of the more inclined spirals due to projection effects.
 In any case. type Ia SNe are more lughly concentrated than ype IL or Ib/c. Though ITatauo et al.," In any case, type Ia SNe are more highly concentrated than type II or Ib/c. Though Hatano et al."
 claim that the observations coufiiii their mode Lowe should mention that vau den Bergh (1997)) and Wane ct al. (1997))," claim that the observations confirm their model, we should mention that van den Bergh \cite{vdb97}) ) and Wang et al. \cite{wang}) )"
 reached the opposite conclusion on the basis of simular data., reached the opposite conclusion on the basis of similar data.
 Even if the Hatano ct al., Even if the Hatano et al.
 model should be considered as exploratory given the above controversy. it is of interest to test row adopting it cau change the SN rate estimates.," model should be considered as exploratory given the above controversy, it is of interest to test how adopting it can change the SN rate estimates."
 We thus have replaced he empirical bias corrections mentioned m Section ?? with the observed SN Iuuinositv distribution for cach SN ype in spirals of different inchnation clerivec from the IIatano et al. (, We thus have replaced the empirical bias corrections mentioned in Section \ref{evans} with the observed SN luminosity distribution for each SN type in spirals of different inclination derived from the Hatano et al. (
1998) inodel.,1998) model.
 Then we computed the control times for cach bin of the huninosity function for cach galaxy anc SN type., Then we computed the control times for each bin of the luminosity function for each galaxy and SN type.
 The total control time was obtained as the weighted average according to the observed Iuninositv cistribution., The total control time was obtained as the weighted average according to the observed luminosity distribution.
 The results of this calculation are shown in Table 3.., The results of this calculation are shown in Table \ref{hattab}.
 Taken at face value and compared with the empirical bias corrections (Tab. Ον ," Taken at face value and compared with the empirical bias corrections (Tab. \ref{final}) ),"
we found that by using the Tatano ct al., we found that by using the Hatano et al.
 inodel. the SN Ia rate iu the cutire sample," model, the SN Ia rate in the entire sample"
"what was found by Haehnelt,Steinmetz,&Rauch(2000).",what was found by \citet{2000ApJ...534..594H}.
". The new compilation of velocity width data by Wolfe,Ga-wiser,&Prochaska(2005) extends to significantly larger velocities and appears to require a somewhat larger value of veo=70kms!."," The new compilation of velocity width data by \citet{2005ARA&A..43..861W} extends to significantly larger velocities and appears to require a somewhat larger value of $v_\ro{c,0}=
70 \kmsec$."
 The apparent lack of neutral gas in small dark matter halos can be plausibly attributed to the feedback effects of star formation and/or photo-heating due to the meta-galactic UV background., The apparent lack of neutral gas in small dark matter halos can be plausibly attributed to the feedback effects of star formation and/or photo-heating due to the meta-galactic UV background.
" In most numerical simulations and models of galaxy formation, these feedback mainly affects halos with somewhat smaller virial velocities than this."," In most numerical simulations and models of galaxy formation, these feedback mainly affects halos with somewhat smaller virial velocities than this."
" So either halos with small virial velocities have larger velocity widths than we have assumed here (ie. p(v«|ve) is different), or else stellar feedback in halos with virial velocities of vc=50kms! is more efficient than generally assumed."," So either halos with small virial velocities have larger velocity widths than we have assumed here (i.e. $p(v_\ro{w} | v_{\ro{c}})$ is different), or else stellar feedback in halos with virial velocities of $v_\ro{c,0}= 50 \kmsec$ is more efficient than generally assumed."
 The cumulative velocity width distribution is not very sensitive to the shape of the off for a>2., The cumulative velocity width distribution is not very sensitive to the shape of the cut-off for $\alpha \ge 2$.
" The high velocity tail of the velocity distribution as compiled by Wolfe,Gawiser,&Prochaska(2005) has proven difficult to reproduce with numerical simulations, which attempt to model the spatial distribution of the neutral gas in DLAs self-consistently rather than assume a scaling of the absorption cross section with the virial velocity of DM halos."," The high velocity tail of the velocity distribution as compiled by \citet{2005ARA&A..43..861W} has proven difficult to reproduce with numerical simulations, which attempt to model the spatial distribution of the neutral gas in DLAs self-consistently rather than assume a scaling of the absorption cross section with the virial velocity of DM halos."
" Generally, the simulations fail to produce a sufficient number of absorption systems with velocity widths as wide as observed."," Generally, the simulations fail to produce a sufficient number of absorption systems with velocity widths as wide as observed."
" This is normally attributed by the authors of these studies to the fact that momentum and energy input into the gas due to star formation may not have been modelled with sufficient sophistication (Nagamine,Springel,2006,2007;Pontzenetal. 2008).."," This is normally attributed by the authors of these studies to the fact that momentum and energy input into the gas due to star formation may not have been modelled with sufficient sophistication \citep{2004MNRAS.348..421N,2007ApJ...660..945N,
2006ApJ...645...55R,2007arXiv0710.4137R,2008arXiv0804.4474P}."
" We now ask whether our model, which successfully fits the velocity width distribution of the associated low-ionization metal absorption of DLAs, can also reproduce the cumulative size distribution of the emission regions of the new population of faint emitters, as shown by black solid curve in Figure 2 (Rauchetal.2008,Figure19,Hicorrected).."," We now ask whether our model, which successfully fits the velocity width distribution of the associated low-ionization metal absorption of DLAs, can also reproduce the cumulative size distribution of the emission regions of the new population of faint emitters, as shown by black solid curve in Figure \ref{fig:plot1} \citep[][ Figure 19, \hi corrected]{2008ApJ...681..856R}."
" It asymptotes to dN/dz=0.23, which is very similar to the observed incidence rate of DLAs at the same redshift."," It asymptotes to $\dd \cN / \dd z = 0.23$, which is very similar to the observed incidence rate of DLAs at the same redshift."
" To calculate the cumulative incidence rate for the DLA host galaxies in our model, we must relate the size of the emission region to the mass of the corresponding DM halo."," To calculate the cumulative incidence rate for the DLA host galaxies in our model, we must relate the size of the emission region to the mass of the corresponding DM halo."
 We do this simply by assuming that the absorption cross-section is related to radius as if it was a sphere., We do this simply by assuming that the absorption cross-section is related to radius as if it was a sphere.
 The cumulative incidence rate is then given by The result is in the left panel of Figure 2.., The cumulative incidence rate is then given by The result is in the left panel of Figure \ref{fig:plot1}.
" The predicted sizes are consistently smaller than the observed sizes by a factor of between 1.5 to 3, suggesting that the emission regions of the emitters are larger than the cross section for damped absorption."," The predicted sizes are consistently smaller than the observed sizes by a factor of between 1.5 to 3, suggesting that the emission regions of the emitters are larger than the cross section for damped absorption."
" Even though the integrated inferred incidence rate is similar to that of DLAs, the emission regions of the emitters can thus not be identical to the regions responsible for damped absorption."," Even though the integrated inferred incidence rate is similar to that of DLAs, the emission regions of the emitters can thus not be identical to the regions responsible for damped absorption."
" Nevertheless, the two can be closely related."," Nevertheless, the two can be closely related."
 To demonstrate this we now explore a simple model where only a fraction fa of the halos are emitting radiation above the detection threshold at any one time., To demonstrate this we now explore a simple model where only a fraction $f_\ro{d}$ of the halos are emitting radiation above the detection threshold at any one time.
" To keep the total inferred incidence rate fixed, we allow the cross section of the individual emission regions to be larger than those for damped absorption by a corresponding factor fjl|"," To keep the total inferred incidence rate fixed, we allow the cross section of the individual emission regions to be larger than those for damped absorption by a corresponding factor $f_\ro{d}^{-1}$."
" The emitters show typical signs of radiative transfer effects, like large velocity widths and asymmetric line profiles (Rauchetal.2008,Section 6.3),, and should thus be optically thick {ο radiation."," The emitters show typical signs of radiative transfer effects, like large velocity widths and asymmetric line profiles \citep[][Section 6.3]{2008ApJ...681..856R}, and should thus be optically thick to radiation."
" However, we have no good handle here on the actual column density required to scatter emitted at large radii effectively."," However, we have no good handle here on the actual column density required to scatter emitted at large radii effectively."
" It may well be lower than that required for damped absorption, in which case it is very plausible that the region for emission extends beyond"," It may well be lower than that required for damped absorption, in which case it is very plausible that the region for emission extends beyond"
during the time periods following the flare. (,during the time periods following the flare. (
"We mark the beginning of the and flares respectively with ""A' and X in the light curves shown in Figs.",We mark the beginning of the and flares respectively with `A' and `X' in the light curves shown in Figs.
 2- 4. , \ref{fig2-hratio}$ $-$ \ref{fig4-hratio}. .)
Mnorderlof uctherquanti f glheamount/absenecof variablilg. wuwinipiyalata test to the hardness ratio trends as shown in the left panels of Figs.," In order to further quantify the amount / absence of variablity, we apply a $\chi^2$ test to the hardness ratio trends as shown in the left panels of Figs."
 2- a, \ref{fig2-hratio}$ $-$ \ref{fig4-hratio}.
ble Idelailstheseresullsforlasdegreesof freedom., Table \ref{tab-hardnessvar} details these results for 138 degrees of freedom.
 We mark the regions corresponding to the and flare events. in Figs., We mark the regions corresponding to the and flare events in Figs.
 2- 4., \ref{fig2-hratio}$ $-$ \ref{fig4-hratio}.
0 helimespancorrespondinglotlhe ASCAand RNTEflaresweree d, The time span corresponding to the and flares were chosen in order that the times surrounding the flares and minimum (following said flare) can be compared.
ifferent ος variability hal io , The spectral behaviour during these periods is thoroughly investigated in Section 4.3.
states:css defined μία sdtha," For present purposes, we draw the reader's attention to the deep minimum immediately following the bright flare. ("
t ALTuml followings,The general location of this minimum in the light curve is labeled with `DM' in Figs.
a , \ref{fig2-hratio}$ $-$ \ref{fig4-hratio}. .)
Despite the hardness ratio results discussed previously (e.g. no obvious trends seen in Its. lt. Re). the spectrum during this time veriod always hardens in all colours. whereas this is not necessarily ye case with the other minima events (e.g. the deep minimum ollowing the soft flare).," Despite the hardness ratio results discussed previously (e.g. no obvious trends seen in $\rm R_3$, $\rm R_5$, $\rm R_6$ ), the spectrum during this time period always hardens in all colours, whereas this is not necessarily the case with the other minima events (e.g. the deep minimum following the soft flare)."
 In addition. it is not clear whether je spectrum reaches its lowest minimum slightly before or after an observed hardening.," In addition, it is not clear whether the spectrum reaches its lowest minimum slightly before or after an observed hardening."
 We also bring attention (by means of a horizontal line iough the hardness ratio versus. time plots in Figs., We also bring attention (by means of a horizontal line through the hardness ratio versus time plots in Figs.
 2- othegeneraltrend forharderspectrainthe- 260 ks time oeriod that begins with the flare., \ref{fig2-hratio}$ $-$ \ref{fig4-hratio}) ) to the general trend for harder spectra in the $\sim$ 260 ks time period that begins with the flare.
 Tany complicated processes contribute to the overall temporal and spectral behaviour of an AGN., Many complicated processes contribute to the overall temporal and spectral behaviour of an AGN.
 It is necessary to disentangle hese components in order to assess the physical processes that are responsible for the observed variability phenomenon., It is necessary to disentangle these components in order to assess the physical processes that are responsible for the observed variability phenomenon.
 Effects due ο reprocessing are significant only at the higher energies. the X keV) band being the most sensitive probe of any subtle changes in the reflection component.," Effects due to reprocessing are significant only at the higher energies, the (10-20 keV) band being the most sensitive probe of any subtle changes in the reflection component."
 Martocchia Matt (1996) suggested hat gravitational bending/focusing as the X-ray source gets closer o the black hole will enhance the amount of reflection., Martocchia Matt (1996) suggested that gravitational bending/focusing as the X-ray source gets closer to the black hole will enhance the amount of reflection.
 Therefore. it is possible that flares closer to the hole during the minima ean enhance the amount of reflection. through increased beaming of he emission towards the disk.," Therefore, it is possible that flares closer to the hole during the minima can enhance the amount of reflection, through increased beaming of the emission towards the disk."
 We conclude however from the combined hardness ratio findings described thus far that the spectral jardening during periods of diminished intensity point largely to changes in the intrinsic photon index I being the main culprit or the observed spectral variability., We conclude however from the combined hardness ratio findings described thus far that the spectral hardening during periods of diminished intensity point largely to changes in the intrinsic photon index $\Gamma$ being the main culprit for the observed spectral variability.
 We are led to this conclusion because spectral changes are seen in all bands. which is most likely ο be due to changes in the spectral index.," We are led to this conclusion because spectral changes are seen in all bands, which is most likely to be due to changes in the spectral index."
 In other words. we may be seeing the effects of changes in the coronal temperature affecting the intrinsic power law slope. or that changes in DL are due to coverage of flares such that we are observing the effects of flares occurring at different heights (e.g. Poutanen Fabian 1999).," In other words, we may be seeing the effects of changes in the coronal temperature affecting the intrinsic power law slope, or that changes in $\Gamma$ are due to coverage of flares such that we are observing the effects of flares occurring at different heights (e.g. Poutanen Fabian 1999)."
 However. as we shall now show. direct spectral analysis reveals a more complex situation.," However, as we shall now show, direct spectral analysis reveals a more complex situation."
 Having established that variability effects are many. we next investigate spectral changes in time sequence in order to obtain a better understanding of variability phenomena between the different flux states with particular emphasis on the flares and deep a," Having established that variability effects are many, we next investigate spectral changes in time sequence in order to obtain a better understanding of variability phenomena between the different flux states with particular emphasis on the flares and deep minima."
nalysis is restricted to the 3 to 20 keV PCA energy band., Data analysis is restricted to the 3 to 20 keV PCA energy band.
 The lower energy restriction at 3 keV is selected in order that the necessity for modelling photoelectric absorption due to Galactic ISM material. or the warm absorber that are known to be present in this object is removed. (," The lower energy restriction at 3 keV is selected in order that the necessity for modelling photoelectric absorption due to Galactic ISM material, or the warm absorber that are known to be present in this object is removed. ("
Lee et al.,Lee et al.
 1999 have shown that effects due to the warm absorber at 3 keV are negligible for this data set.), 1999 have shown that effects due to the warm absorber at 3 keV are negligible for this data set.)
 We choose the time intervals (Fig. 1) , We choose the time intervals (Fig. \ref{fig1-ascaxteltc}) )
in order to contrast and study μοδαthé Theyrroundine, in order to contrast and study the different variability states.
 are such int and correspond to the two deep minima. and intervals and for the relatively calm periods following these minima. and a flare event for comparison.," They are defined such that intervals and correspond to the two deep minima, and intervals and for the relatively calm periods following these minima, and a flare event for comparison."
 We investigate in detail in Section 4.3 the times surrounding the counterpart of the soft flare. and the hard flare observed by (unfortunately missed by ). as depicted in Fig.," We investigate in detail in Section 4.3 the times surrounding the counterpart of the soft flare, and the hard flare observed by (unfortunately missed by ), as depicted in Fig."
 |. and subsequently in Figs 9.. and |1..," \ref{fig1-ascaxteltc}
 and subsequently in Figs \ref{fig10-ascaflare}, and \ref{fig12-xteflare}."
 Both the and light curves for the ull observation are shown in Fig. I.., Both the and light curves for the full observation are shown in Fig. \ref{fig1-ascaxteltc}.
 We also separate the data according to flux in order to assess whether a clear picture can be developed of the dominant processes hat may be at work for a given flux state., We also separate the data according to flux in order to assess whether a clear picture can be developed of the dominant processes that may be at work for a given flux state.
 To do so. we separate the 400 ks observation by flux. with and being the intermediate states between the lowest£142.84. 1044 and ughestf4 (4.7910 ~)) fluxes.," To do so, we separate the 400 ks observation by flux, with and being the intermediate states between the lowest $\rm 2.84 \times 10^{-11}$ ), and highest $\rm 4.79 \times 10^{-11}$ ) fluxes."
 Properties: of. hese flux levels (in the 2-60 keV. and 3-20 keV energy bands). and intervals are detailed respectively in Table 2.. and Table 7..," Properties of these flux levels (in the 2-60 keV, and 3-20 keV energy bands), and $-$ intervals are detailed respectively in Table \ref{tab1-pl}, and Table \ref{tab5-intproperties}."
 A nominal fit to the entire data set Ge..," A nominal fit to the entire data set (ie.,"
 PCA and HEXTE) demonstrated the clear existence of a redshifted broad iron Ka line at ~ 6.0 keV and reflection hump between 20 and 30 keV as shown in Lee et al. (, PCA and HEXTE) demonstrated the clear existence of a redshifted broad iron $\alpha$ line at $\sim$ 6.0 keV and reflection hump between 20 and 30 keV as shown in Lee et al. (
1998. 1999).,"1998, 1999)."
 In general. fits to the solar abundance models of George Fabian 1991 (hereafter GF91) are better than simple power law fits and further reinforce the preference for a reflection component.," In general, fits to the solar abundance models of George Fabian 1991 (hereafter GF91) are better than simple power law fits and further reinforce the preference for a reflection component."
 Since the reflected continuum does not contribute significant flux to the observed spectrum below 10 keV. a simple power law and iron line fit to the data below 10 keV will reveal changes in the intrinsic power-law of the X-ray source.," Since the reflected continuum does not contribute significant flux to the observed spectrum below 10 keV, a simple power law and iron line fit to the data below 10 keV will reveal changes in the intrinsic power-law of the X-ray source."
" If indeed spectral variability is due to changes in the intrinsic power law slope which would implicate changes in the conditions of the X-ray emitting corona. we should see noticeable changes in the values of Εν5, between the different temporal states."," If indeed spectral variability is due to changes in the intrinsic power law slope which would implicate changes in the conditions of the X-ray emitting corona, we should see noticeable changes in the values of $\Gamma_{(3-10)}$ between the different temporal states."
" However. because reflection only significantly affects energies above 10 keV. changes seen in [ντο20, would suggest that variability may be due to the amount of reprocessing. ("," However, because reflection only significantly affects energies above 10 keV, changes seen in $\Gamma_{(10-20)}$ would suggest that variability may be due to the amount of reprocessing. ("
We have tested this premise by comparing simulated spectra in using the model and find that the overall flux contribution from reflection alone is > 60 per cent above 10 keV.) This would have strong implications for geometry Ge.,We have tested this premise by comparing simulated spectra in using the model and find that the overall flux contribution from reflection alone is $>$ 60 per cent above 10 keV.) This would have strong implications for geometry (ie.
 where the direct X-ray flares are partially obscured). motion of the source (eg.," where the direct X-ray flares are partially obscured), motion of the source (eg."
 Reynolds Fabian 1997: Beloborodov 1998). or gravity," Reynolds Fabian 1997; Beloborodov 1998), or gravity"
τ,.
"ς Moreover, on noting that (e,.....εφ} Is exchangeable by we σοι As in earlier schemes we obtain representation with Moreover, the condition is fulfilled since νου. VarS,, and L, are all of order »."," Moreover, on noting that $(\eps_1,\ldots,\eps_n)$ is exchangeable by we get As in earlier schemes we obtain representation with Moreover, the condition is fulfilled since $\E\,S_n$ , $\var\,S_n$ and $L_n$ are all of order $n$."
 See again Remark 4.2 to conclude that holds true., See again Remark \ref{orden} to conclude that holds true.
" Let 5,, be a random variable with negative hypergeometric distribution of the first kind, that is, with e,=no and 9%,=2."," Let $S_n$ be a random variable with negative hypergeometric distribution of the first kind, that is, with $\alpha_n=n\alpha$ and $\beta_n=n\beta$."
" The CLT for $, states that for η—x", The CLT for $S_n$ states that for $n\to\infty$.
" To quickly derive it [rom Theorem 4.1., let £,,—qna; and note that Further, the 7th factorial moment of ον is easily derived as Thus.again due to we concludethat representation holds with The final result follows by Remark 4.2.."," To quickly derive it from Theorem \ref{gen}, let $L_n=\frac{n\alpha}{\alpha+\beta}$ and note that Further, the $i$ th factorial moment of $S_n$ is easily derived as Thus,again due to we concludethat representation holds with The final result follows by Remark \ref{orden}. ."
GRB 021004 triggered HETE II on 2002 October 4 at 12:06:13 UT.,GRB 021004 triggered HETE II on 2002 October 4 at 12:06:13 UT.
 The burst lasted T~100 seconds (Shirasaki et al., The burst lasted $T\sim 100$ seconds (Shirasaki et al.
 2002) and the 7—400 keV {lnence was ~3.2x10° eres 7 (Lamb et al., 2002) and the $7-400$ keV fluence was $\sim 3.2\times10^{-6}$ ergs $^{-2}$ (Lamb et al.
 2002)., 2002).
 The prompt localization of GRB 021004 by ETE II allowed the follow-up of the alterelow at very early time., The prompt localization of GRB 021004 by HETE II allowed the follow-up of the afterglow at very early time.
 Fox (2002) detected an optical transient ~9 mins after the trigger at the level of 2~15.56 mag., Fox (2002) detected an optical transient $\sim 9$ mins after the trigger at the level of $R\sim 15.56$ mag.
 Torii. Kato and Yamaoka (2002) observed the error box of the burst ~3.5 mins after the trigger. vielding upper limit around 2~13.6 mag at the position of the transient discovered bv Fox (2002).," Torii, Kato and Yamaoka (2002) observed the error box of the burst $\sim 3.5$ mins after the trigger, yielding upper limit around $R\sim 13.6$ mag at the position of the transient discovered by Fox (2002)."
 The spectroscopic observations of the optical alterglow revealed an emission line interpreted as Ly-a at 2=2.328 (Mirabal el al., The spectroscopic observations of the optical afterglow revealed an emission line interpreted as $\alpha$ at $z=2.328$ (Mirabal et al.
 2002)., 2002).
" Assuming ο=0.3.A,0.7 and h=0.6. the isotropic gamma-ray energy is es5.6xLO”? eres."," Assuming $\Omega_0=0.3,\lambda_0=0.7$ and $h=0.6$, the isotropic gamma-ray energy is $\sim 5.6\times10^{52}$ ergs."
 The dense sampling of the alterglow lisht curve at early time revealed (he peculiarity. a major re-brightening around ~0.1 day after the trigger and a short time scale variability around ~I day.," The dense sampling of the afterglow light curve at early time revealed the peculiarity, a major re-brightening around $\sim 0.1$ day after the trigger and a short time scale variability around $\sim 1$ day."
" The slope of the afterglow is ~/"""" for the earliest three observations (9-17 min alter trigger) bv Fox (2002). and after the luminosity increases around. ~0.1 davs. it decaved with ~1 as usual alterglows do."," The slope of the afterglow is $\sim t^{-0.7}$ for the earliest three observations (9-17 min after trigger) by Fox (2002), and after the luminosity increases around $\sim 0.1$ days, it decayed with $\sim t^{-1.05}$ as usual afterglows do."
 These temporal features might be modeled by, These temporal features might be modeled by
We solve the RT equation on a uniform Cartesian grid of 128% cells (i.e.. the root grid ol the AMIR. simulation).,"We solve the RT equation on a uniform Cartesian grid of $128^3$ cells (i.e., the root grid of the AMR simulation)."
 The total radiation field is divided into direct. photons coming from point sources and the diffuse fIux originating from recombination and bremsstralilung processes in the volume., The total radiation field is divided into direct photons coming from point sources and the diffuse flux originating from recombination and bremsstrahlung processes in the volume.
 From the computational perspective. the only dillerence between these (wo components is (he number of active sources.," From the computational perspective, the only difference between these two components is the number of active sources."
" For example. as discussed in details below. by the end of the run at z=4 our entire volume featiwes ~2x10* star forming (SF) cells and ~2xLO"" gas cells."," For example, as discussed in details below, by the end of the run at $z=4$ our entire volume features $\sim 2\times 10^3$ star forming (SF) cells and $\sim 2\times 10^6$ gas cells."
 The light crossing time across the computational volume is much shorter than the timescales for opacity changes., The light crossing time across the computational volume is much shorter than the timescales for opacity changes.
 This allows us to omit (he term in the RT equation (see also Abel et al., This allows us to omit the time-dependent term in the RT equation (see also Abel et al.
 1999 and Norman. Paschos. Abel 1993).," 1999 and Norman, Paschos, Abel 1998)."
 The only exception is a logical switeh which does not allow anv significant. changes in the Iuminositv (greater than 1054) of individual stellar sources to propagate faster than ie speed of light., The only exception is a logical switch which does not allow any significant changes in the luminosity (greater than $10\%$ ) of individual stellar sources to propagate faster than the speed of light.
" At each (timestep we compute the luminosity ulZi, of every source and ‘compare it to its luminosity Zt4 al (he previous timestep.", At each timestep we compute the luminosity $L_{\rm UV}^n$ of every source and compare it to its luminosity $L_{\rm UV}^{n-1}$ at the previous timestep.
" If it changes by more than 104 en Only those cells whieh are within eM"" from the source will be affected. and as we cross ve CAL” yadius along each ray. we multiply the photon count by £7\/ulLi."," If it changes by more than $10\%$ then only those cells which are within $c\Delta
t^n$ from the source will be affected, and as we cross the $c\Delta
t^n$ radius along each ray, we multiply the photon count by $L_{\rm
UV}^{n-1}/L_{\rm UV}^n$."
" We include this V.witch because the box light crossing (ime Των (86 Myr) is not negligible compared to the imestep (1 Myr). although in practice it is not always important because the Iuminosity of individual sources rises ancl declines slowly on Jag, (Myr to tens of Myr) which is large compared to 7j."," We include this switch because the box light crossing time $\tau_{\rm box}$ (3–6 Myr) is not negligible compared to the timestep (1 Myr), although in practice it is not always important because the luminosity of individual sources rises and declines slowly on $\tdyn$ (Myr to tens of Myr) which is large compared to $\tau_{\rm box}$."
 The radiation field is reconstructed from the 3D source fincton at each timestep., The radiation field is reconstructed from the 3D source function at each timestep.
 Lu this simulation. for both the direct ancl diffuse components. we adopt a photon conserving scheme similar to the one suggested by Abel et al. (," In this simulation, for both the direct and diffuse components, we adopt a photon conserving scheme similar to the one suggested by Abel et al. ("
1999).,1999).
 In our present model. for both components we asstune a power law (with index a= 5) spectrum of the photon munber density.," In our present model, for both components we assume a power law (with index $\alpha=5$ ) spectrum of the photon number density."
" All radiative (transfer is currently done in a single frequency. group above the hydrogen Lyman lint 14,."," All radiative transfer is currently done in a single frequency group above the hydrogen Lyman limit $\nu_{\rm
Ly}$."
 Transfer of direct stellar photons is computed explicitly using the equad-tree technique ol Abel&Wandelt(2001)., Transfer of direct stellar photons is computed explicitly using the quad-tree technique of \citet{abel01}.
. We refer the reader to this paper for the details. here we just briefly outline the algorithm.," We refer the reader to this paper for the details, here we just briefly outline the algorithm."
 Around each source we build a (ree of ravs. starting with 12 ravs ab level /=1. with individual ταν segments splitting recursively into four child ravs each as thev move further away [rom the source.," Around each source we build a tree of rays, starting with 12 rays at level $l=1$, with individual ray segments splitting recursively into four child rays each as they move further away from the source."
 The total number of ravs al anv given level /=1.2.... of the hierarchy is 12x41|. with each ταν corresponding to the same equal area of (he sphere.," The total number of rays at any given level $l=1,2,...$ of the hierarchy is $12\times 4^{l-1}$, with each ray corresponding to the same equal area of the sphere."
 The length of individual ray seements is chosen such that on one hand, The length of individual ray segments is chosen such that on one hand
As with the results of the torque dUfds. the main clilferences between two disc models (MMSN vs 807) are the position of the dips and the steepness of the scaled torque as a function of z.,"As with the results of the torque $d \Gamma^L/ dz$, the main differences between two disc models (MMSN vs S07) are the position of the dips and the steepness of the scaled torque as a function of $z$."
 These dillerences arises from the change in the temperature structure discussed in 8:. that is. the steeper surface density. slope provides lower temperatures around the planet and the mid-plane region.," These differences arises from the change in the temperature structure discussed in $\S$ 3.2, that is, the steeper surface density slope provides lower temperatures around the planet and the mid-plane region."
 Except for these. the basic features are the same as above.," Except for these, the basic features are the same as above."
 In both disc models. massive planets migrate more quickly than low mass planets.," In both disc models, massive planets migrate more quickly than low mass planets."
 This trend is clearly shown in Fig., This trend is clearly shown in Fig.
 T which shows the migration time and rate as a function of orbital radius for planets with various masses (top and bottom panels. respectively).," \ref{fig7} which shows the migration time and rate as a function of orbital radius for planets with various masses (top and bottom panels, respectively)."
 The migration time increases for planets with large orbital radii. as expected.," The migration time increases for planets with large orbital radii, as expected."
 Also. migration in the MAMSN. disc models is more rapid relative to the S07 disce mocels. as discussed above.," Also, migration in the MMSN disc models is more rapid relative to the S07 disc models, as discussed above."
 We found that migration in the MMSN clise models is up to 4.0 times [Faster at Lauthan in the S07 disc moclels (see the left column of Table 2))., We found that migration in the MMSN disc models is up to 4.0 times faster at 1 au than in the S07 disc models (see the left column of Table \ref{table2}) ).
 lt is interesting that he migration time for planets in disces around M. stars is comparable to that of classical T ‘Pauri svestems., It is interesting that the migration time for planets in discs around M stars is comparable to that of classical T Tauri systems.
 This can be shown by making an order of magnitude estimate., This can be shown by making an order of magnitude estimate.
 The tidal torque ds. sealed by Mo(ruhy) CEEWO2:T)., The tidal torque is scaled by $\Sigma_p (r_p/h_p)^2$ \citep[TTW02;][]{pbck09}.
" At 1 au. dises around. classical To ‘Lauri systems typically have Mtrhy~10 ο ?(1/0.1D?=10"" (?)."," At 1 au, discs around classical T Tauri systems typically have $\Sigma_p (r_p/h_p)^2 \simeq 10^4$ g $^{-2} \times (1/0.1)^2 = 10^6$ \citep{cg97}."
" For disces around. AL stars. Mo(ru/huyo407 g em7?«(1/0.01)?=10"" (807)."," For discs around M stars, $\Sigma_p (r_p/h_p)^2 \simeq 10^2$ g $^{-2} \times (1/0.01)^2 = 10^6$ (S07)."
 Thus. the planetary migration time for discs around both stars is similar.," Thus, the planetary migration time for discs around both stars is similar."
 Note that. in these above simple caleulations. we used characteristic values of Γρ to estimate an order of the timescale.," Note that, in these above simple calculations, we used characteristic values of $r_p/h_p$ to estimate an order of the timescale."
 In order to further elucidate the ellects. of the mass of he planet. we replot the migration time as a function of planetary mass at. two cilferent orbital radii (Fig. 8)).," In order to further elucidate the effects of the mass of the planet, we replot the migration time as a function of planetary mass at two different orbital radii (Fig. \ref{fig8}) )."
 If these elfects. arising from the eravitational force. of a λαοί were neglected. the migration time would scale as an inverse function of their masses denoted by the dotted lines (also see equations (3.5)). CX4)) and. (ALL).," If these effects arising from the gravitational force of a planet were neglected, the migration time would scale as an inverse function of their masses denoted by the dotted lines (also see equations \ref{t_mig}) ), \ref{torque_density}) ) and \ref{torque_tot}) ))."
 Obviously. our calculations show some deviation from this function or both dise models (although it is not so large).," Obviously, our calculations show some deviation from this function for both disc models (although it is not so large)."
 The deviation increases with planetary masses. resulting in faster migration.," The deviation increases with planetary masses, resulting in faster migration."
 This indicates that the increment of the torque around the mid-plane is larger than the decrement by the dip formation around z—rg. and that the dilference Increases with planetary mass.," This indicates that the increment of the torque around the mid-plane is larger than the decrement by the dip formation around $z=r_H$, and that the difference increases with planetary mass."
 JS05 found that the deviation increases with planetary mass. but i results in slower migration.," JS05 found that the deviation increases with planetary mass, but it results in slower migration."
 This indicates that. in their disc mocels. the decrement is larger than the increment.," This indicates that, in their disc models, the decrement is larger than the increment."
 Lt is well known that the dip formation strongly depends on t1e erazing angle at which photons emitted from the star strike the disc surface. (2).., It is well known that the dip formation strongly depends on the grazing angle at which photons emitted from the star strike the disc surface \citep{js04}.
 For the discs around the classical T Tauri stars J805 examined. the scale height is laree. and consequently the grazing angle is also large. compared wih the disces around. M stars we examined.," For the discs around the classical T Tauri stars JS05 examined, the scale height is large, and consequently the grazing angle is also large, compared with the discs around M stars we examined."
 As a result. the decrement by the dip formation is enhanced.," As a result, the decrement by the dip formation is enhanced."
 In addition. the deviation becomes. smaller with increasing orbital racius of the planet.," In addition, the deviation becomes smaller with increasing orbital radius of the planet."
 This is understood by the power-law function of the surface density., This is understood by the power-law function of the surface density.
 Since the ise densitv decreases with the distance from the central star. the density and temperature distortions caused by the eravitational force of the planet. are diminished. at larger disc radii.," Since the disc density decreases with the distance from the central star, the density and temperature distortions caused by the gravitational force of the planet are diminished at larger disc radii."
 Consequently. the migration time approaches the simple analvtical solution. which is dillerent. from. 805.," Consequently, the migration time approaches the simple analytical solution, which is different from JS05."
 _Ve sugeest that the main reason for the dillerence is that they targeted massive dises around. classical “Po Tauri stars including viscous heating (although our target is low mass discs around an M stars withou viscous. heating)., We suggest that the main reason for the difference is that they targeted massive discs around classical T Tauri stars including viscous heating (although our target is low mass discs around an M stars without viscous heating).
 In addition. the orbital radit of the planet they considered are 0.5 to 4 au.," In addition, the orbital radii of the planet they considered are 0.5 to 4 au."
 For such a short range. the cllects arising from the planet on the disc are not diminished.," For such a short range, the effects arising from the planet on the disc are not diminished."
 Our results. bv contrast. pertain to geometrically thin. low mass clises ancl the outer part. of massive clises for which viscous heating is safely neglected while the results of JSO5 are for geometrically thick. (the inner part of) massive cises in which viscous heating is not negligible.," Our results, by contrast, pertain to geometrically thin, low mass discs and the outer part of massive discs for which viscous heating is safely neglected while the results of JS05 are for geometrically thick, (the inner part of) massive discs in which viscous heating is not negligible."
" lt is interesting that the deviation from 1/Ad, is almost. identical for both clise mocels if the migration times are appropriately scaled.", It is interesting that the deviation from $1/M_p$ is almost identical for both disc models if the migration times are appropriately scaled.
 Also. the cdilference of the mügration times between two cise models becomes. small with increasing the orbital radius of the planets.," Also, the difference of the migration times between two disc models becomes small with increasing the orbital radius of the planets."
 In this section. we examine the elfects of dust settling in the presence of planets with various masses.," In this section, we examine the effects of dust settling in the presence of planets with various masses."
 We also discuss the dilference between the well mixed and dust settling cases., We also discuss the difference between the well mixed and dust settling cases.
 We again [focus on the behaviour of the torque. dL’(ry.z)/dz. exerted by a laver of gas at height = (sec equation (3.2))).," We again focus on the behaviour of the torque, $ d \Gamma^L (r_p,z)/ dz$, exerted by a layer of gas at height $z$ (see equation \ref{layer_torque}) ))."
 As is the case of well mixed dust. the results of the two disc models are similar. so that we only. present the results of the the S07 dise model (Xx&. 1).," As is the case of well mixed dust, the results of the two disc models are similar, so that we only present the results of the the S07 disc model $\Sigma \propto r^{-1}$ )."
 We discuss the dillerences between them later., We discuss the differences between them later.
 Fie., Fig.
 9 shows the torque dV(n.ο)ας as a function of 2 dor the planets with various masses.," \ref{fig9} shows the torque $d \Gamma^L (r_p,z)/ dz $ as a function of $z$ for the planets with various masses."
 Compared. with the well mixed case (Fig. 5)).," Compared with the well mixed case (Fig. \ref{fig5}) ),"
 dust settling provides a stronger torque around the mid-plane region. as expected. (see § 3)).," dust settling provides a stronger torque around the mid-plane region, as expected (see $\S$ \ref{torque}) )."
 ‘This arises from lower temperatures in the mid-plane region induced mainly by dust settling., This arises from lower temperatures in the mid-plane region induced mainly by dust settling.
 Since a small change at such lower temperatures strongly allects the resonant positions. the distribution of the torque slightly Ductuates. especially around the mid-plane region.," Since a small change at such lower temperatures strongly affects the resonant positions, the distribution of the torque slightly fluctuates, especially around the mid-plane region."
 Although one might worry that such Ductuations cause spurious ellects. we checked," Although one might worry that such fluctuations cause spurious effects, we checked"
We first investigated: whether our samples of low-recshilt 151X galaxies are located predominantly. in the Ποιά. as found in the seminal study of Zablucloll et (1996). or in the proximity of rich clusters.,"We first investigated whether our samples of low-redshift E+A galaxies are located predominantly in the field, as found in the seminal study of Zabludoff et (1996), or in the proximity of rich clusters."
 A catalogue of galaxy clusters within the 2dEGIUS observed sky areas was compiled: by de Propris ct ((2002). sourced. from. the Abell APAL ancl Edinburgh-Durham Cluster (IEEDCX) catalogues.," A catalogue of galaxy clusters within the 2dFGRS observed sky areas was compiled by de Propris et (2002), sourced from the Abell, APM and Edinburgh-Durham Cluster (EDCC) catalogues."
 The 2PdPGRS was utilized by ce Propris et tto. measure precise redshifts. velocity clispersions and centroids for these clusters.," The 2dFGRS was utilized by de Propris et to measure precise redshifts, velocity dispersions and centroids for these clusters."
 We simply measured the vector separation of each E|A galaxy and the catalogued clusters., We simply measured the vector separation of each E+A galaxy and the catalogued clusters.
" Two components of separation were measured: the transverse pprojected) physical distance D,=(r;;N8)/(1|z). where r is the racial co-moving distance to the I2|A galaxy at recshilt z and A@ is the angular separation of the galaxy and. eluster centroid: and the racial (redshift-space) physical distance D,—(ολα|2). where 4 is the Hubble constant at the I[X galaxy redshift and As is the redshift separation of the galaxy ancl cluster centroid."," Two components of separation were measured: the transverse projected) physical distance $D_t = (r \times
\Delta\theta)/(1+z)$, where $r$ is the radial co-moving distance to the E+A galaxy at redshift $z$ and $\Delta\theta$ is the angular separation of the galaxy and cluster centroid; and the radial (redshift-space) physical distance $D_r = (c \, \Delta z/H)/(1+z)$, where $H$ is the Hubble constant at the E+A galaxy redshift and $\Delta z$ is the redshift separation of the galaxy and cluster centroid."
 “Phe radial separation was considered independently because its apparent value can be significantly enhanced by line-of-sight peculiar velocities., The radial separation was considered independently because its apparent value can be significantly enhanced by line-of-sight peculiar velocities.
" An EVA galaxy was considered a ""cluster object if both D,<ny anc D,Yoo?|(GofH). where σ is the velocity dispersion of the cluster in question. typically σ1000 km ?."," An E+A galaxy was considered a `cluster' object if both $D_t < r_0$ and $D_r < \sqrt{(r_0)^2 + (2\sigma/H)^2}$, where $\sigma$ is the velocity dispersion of the cluster in question, typically $\sigma \sim 1000$ km $^{-1}$."
 The selection. volume is thus broadened in the radial direction to allow for peculiar velocities., The selection volume is thus broadened in the radial direction to allow for peculiar velocities.
 The critical physical scale ry was taken as 5 Alpe (chosen to exceed comfortably the virial radius of a typical rich cluster), The critical physical scale $r_0$ was taken as 5 Mpc (chosen to exceed comfortably the virial radius of a typical rich cluster).
 We restricted. our cluster analysis to the SGP region, We restricted our cluster analysis to the SGP region
exact numbers will chaiee for objects of comparable mass. bit we expect the qualitative effect to be similar: the effective spin will not be near extremal for σοιiparable-miass objects. so clyuamica instability may still set in prior to tidal stripping.,"exact numbers will change for objects of comparable mass, but we expect the qualitative effect to be similar: the effective spin will not be near extremal for comparable-mass objects, so dynamical instability may still set in prior to tidal stripping."
 C'ouuteriutuitively. therefore. somewhat.Aiigher mass ratios could be more likely to result i accretion disks and stabe mass trausler for a rapidly spiuniug black hole. because the effective spit could be closer to maxiual.," Counterintuitively, therefore, somewhat mass ratios could be more likely to result in accretion disks and stable mass transfer for a rapidly spinning black hole, because the effective spin could be closer to maximal."
 It iight. however. still be cdifficult to generate stable mass traus," It might, however, still be difficult to generate stable mass transfer."
 For exaimple. consider a extreme Ixerr black hole of mass mpi=YM. orbited by a neutron star of nass mys=LEAL. and radius Hoμήνα.," For example, consider a extreme Kerr black hole of mass $m_{\rm
BH}=7\,M_\odot$ orbited by a neutron star of mass $m_{\rm
NS}=1.4\,M_\odot$ and radius $R_0=5m_{\rm NS}$."
 The higher mass ratio meais that the full spacetime is likely to be reasonably. close to err. makiug extrapolatious nore reliaje.," The higher mass ratio means that the full spacetime is likely to be reasonably close to Kerr, making extrapolations more reliable."
 We have Ry=171 aud Fide=2.3., We have ${\hat R}_0=1.71$ and ${\hat r}_{\rm tide}=2.3$.
 Then r/ipy=3.93. so r/AM=3.28.," Then $r/m_{\rm
BH}=3.93$, so $r/M=3.28$."
 Using tle Damour (2001) approximatOl. dey0.50. so in the Ixerr spacetime rjsco/M=2.91.," Using the Damour (2001) approximation, ${\hat a}_{\rm eff}=0.80$, so in the Kerr spacetime $r_{\rm
ISCO}/M=2.91$."
 Taking hese numbers at [ace value. even in this case gravitational racjatiou losses would cause the neurou star to plunge before it was stripped.," Taking these numbers at face value, even in this case gravitational radiation losses would cause the neutron star to plunge before it was stripped."
 However. this is clOne erough to the tidal radius that deails are iuportant: lor example. if in reality the ellective spi1 parameter is (a0.90. then risco/M2.32 and stable mass trausfer might eusue.," However, this is close enough to the tidal radius that details are important: for example, if in reality the effective spin parameter is ${\hat a}_{\rm eff}=0.90$, then $r_{\rm ISCO}/M=2.32$ and stable mass transfer might ensue."
 We are therefore in a domain where existiig approximatious are inadequate to determine the [ate of the inmergiug neutrou star., We are therefore in a domain where existing approximations are inadequate to determine the fate of the merging neutron star.
 Full uo-approximaion general relativistic calculations will be necessary., Full no-approximation general relativistic calculations will be necessary.
 Nonetheless. it appears possible that pausible combination of ueutron star aid black bole masses. spius. aud orbital iuclinatious will resilt. in an accretion disk or stable mass trausfer.," Nonetheless, it appears possible that plausible combination of neutron star and black hole masses, spins, and orbital inclinations will result in an accretion disk or stable mass transfer."
 Mergers of ieutron stars with black holes are of great interest [or uodels of short. gamuina-ray bursts aud as sources of ueh-frequency gravitational radiation., Mergers of neutron stars with black holes are of great interest for models of short gamma-ray bursts and as sources of high-frequency gravitational radiation.
 The likely siguatu‘es of such eveuts depeud strougly ou wheher he star is swallowed whole iu a direct me‘eer or Whether it is tidally stripped far enowel away [ron the black hole that an accreion disk [odus Or stable mass trai occurs., The likely signatures of such events depend strongly on whether the star is swallowed whole in a direct merger or whether it is tidally stripped far enough away from the black hole that an accretion disk forms or stable mass transfer occurs.
 We have shown hat lor slowly rotating black holes. current calculations of tida stripy aud of the spacetimes of binary compact objects suggest tha clirect mereer will almost always take place.," We have shown that for slowly rotating black holes, current calculations of tidal stripping and of the spacetimes of binary compact objects suggest that direct merger will almost always take place."
 This is in agreemeu with recent general relativistic uunerica simulations. aud could actually enhance the likeihood o ‘a baryon-[ree environment for a sliOrt gamuna-ray burst (e.g.. Mésszárros Rees 1992).," This is in agreement with recent general relativistic numerical simulations, and could actually enhance the likelihood of a baryon-free environment for a short gamma-ray burst (e.g., Mésszárros Rees 1992)."
 A prograde encounter of a neutron= star witl a rapidly 'otatiug black hole is more likely to result i1 a disk. but the current uiMrstaudiug of such close mergers is uncertain enough that it could be thaty NS-BH merger wil ye direct.," A prograde encounter of a neutron star with a rapidly rotating black hole is more likely to result in a disk, but the current understanding of such close mergers is uncertain enough that it could be that NS-BH merger will be direct."
 Future fuly general relativistic numerical simulations will be required to resolve tliis issue., Future fully general relativistic numerical simulations will be required to resolve this issue.
discarding also the term @&2(f) leads to the imequality (23)).,discarding also the term $\Phi_2(t)$ leads to the inequality \ref{eqcap1}) ).
 This is crucial for proving: Let i!. for j=0.1.2. be defined as in Lemna S.L.," This is crucial for proving: Let $\mu^{(j)}$, for $j=0,1,2$, be defined as in Lemma \ref{lem-cap}."
 Let also y=suplby(d1076)|dq(ót):fC[1.831V.," Let also $\eta = - \sup \{\Phi_0(\delta^{-1} \bar{\delta}^2 t)
 + \Phi_1(\bar{\delta} t); \; t \in
[1,\delta^{-1}] \} $."
 Then. if yp>ilogs13 the measure p? is absolutely continuous with an L? density. and if ij>loe(é+) the density ipie) is continuous.," Then, if $\eta > \frac{1}{2} \log(\delta^{-1})$ the measure $\mu^{(2)}$ is absolutely continuous with an $L^2$ density, and if $\eta > \log(\delta^{-1})$ the density $\rho(\mu;x)$ is continuous."
" Frou, Loenuua S.l itf follows that for all NU>0. and for all the function ο(1) is less than ΝΑ."," From Lemma \ref{lem-cap} it follows that for all $N>0$, and for all $t \in [\delta^{-N},\delta^{-N-1}]$ the function $\Phi_2(t)$ is less than $-\eta N$."
 One can therefore apply the Sobolev criterion described in Sect. 6.., One can therefore apply the Sobolev criterion described in Sect. \ref{sec-sobo}.
 This theorem quickly turns iuto a computer.assisted proof of absolute coutiuuity., This theorem quickly turns into a computer–assisted proof of absolute continuity.
 Iu fact. notice that the supreunmui required to compute iis mdeed the maxi ofa continuous function defined ou a finite interval.," In fact, notice that the supremum required to compute $\eta$ is indeed the maximum of a continuous function defined on a finite interval."
 As such. it is easily computable with sufficient precision to establish whether the inequalities in Thin.," As such, it is easily computable with sufficient precision to establish whether the inequalities in Thm."
 3. are satisfied., \ref{gio2} are satisfied.
 For the systems we have been cousideriug m this paper this is casily accomplished. by computing the Fourier transforms via algorithm 2.," For the systems we have been considering in this paper this is easily accomplished, by computing the Fourier transforms via algorithm 2."
 As just au instance of this procedure. we can prove Let y'!=4(A1|Ay). and let pl). j=1.2 defined as in Lemma ;S.L.," As just an instance of this procedure, we can prove Let $\mu^{(0)}=\frac{1}{2}(\Delta_{-1}+\Delta_1)$, and let $\mu^{(j)}$, $j=1,2$ defined as in Lemma \ref{lem-cap}."
 Then. if tpsà=2 the measure ii?) 2)ds absolutely continuous. with. a continuous. density. and if 0=10E itis absolutelg continuous with an L density.," Then, if $\delta=\frac{2}{5}$ the measure $\mu^{(2)}$ is absolutely continuous with a continuous density, and if $\delta=\frac{3}{10}$ it is absolutely continuous with an $L^2$ density."
 Tu the two cases quoted. one can easily compute that y/loe(é1)~1.5 aud flog1)296838. respectively.," In the two cases quoted, one can easily compute that $\eta/\log(\delta^{-1}) \simeq 1.5035$ and $\eta/\log(\delta^{-1}) \simeq
.96838$, respectively."
 It appears nuucricallv that iu addition to beime a rigorous inequality. formula (23)) is a rather precise estimate of the asviuptotic behavior of the Fourier trausform of 7).," It appears numerically that in addition to being a rigorous inequality, formula \ref{eqcap1}) ) is a rather precise estimate of the asymptotic behavior of the Fourier transform of $\mu^{(2)}$."
 So. by letting now à= we find that the value y/loe(é1)~ implics that the corresponding measure gf?! is singular continous.," So, by letting now $\delta=\frac{2}{10}$ we find that the value $\eta/\log(\delta^{-1})
\simeq 0.2218$ implies that the corresponding measure $\mu^{(2)}$ is singular continuous."
 This conclusion is also confirmed by the iterative analysis detailed iu Sect. 6.., This conclusion is also confirmed by the iterative analysis detailed in Sect. \ref{sec-sobo}.
 We have described nuuerical techniques to determine the continuity type of measures eenerated by affine. homogencous LE.S. with uucouutably many maps.," We have described numerical techniques to determine the continuity type of measures generated by affine, homogeneous I.F.S. with uncountably many maps."
 These techniques are inspired bv the theory of ανασα. systems., These techniques are inspired by the theory of dynamical systems.
 Iu turn. they sugeest methods to obtain rigorous. and computerassisted proofs one should focus ou the existence (or nou.existence) of an attractive fixed point iu a suitable space of densities. most conveniently uuder the bounded: variation distance.," In turn, they suggest methods to obtain rigorous, and computer–assisted proofs: one should focus on the existence (or non–existence) of an attractive fixed point in a suitable space of densities, most conveniently under the bounded variation distance."
 Alternatively one must consider the decay rate of Fourier cocticieuts. also obtained via fixed point procedures.," Alternatively, one must consider the decay rate of Fourier coefficients, also obtained via fixed point procedures."
 We have sketched examples of both techuiques. that have helped us to clarity the fact that such LE.S. measures can be of all continuity types.," We have sketched examples of both techniques, that have helped us to clarify the fact that such I.F.S. measures can be of all continuity types."
 According to the language adopted in this paper. we have described the implications of the continuity of & on those of p in what we believe to be a rather detailed picture.," According to the language adopted in this paper, we have described the implications of the continuity of $\sigma$ on those of $\mu$ in what we believe to be a rather detailed picture."
 We have also shown applications of our methods to the classical problem of, We have also shown applications of our methods to the classical problem of
As illustrated in Fig.,As illustrated in Fig.
 1. the preseut dav solar Li abundance = is nuch lower than he meteoritic Li abundance =," 1, the present day solar Li abundance = is much lower than the meteoritic Li abundance =."
 This large depletion iu the observed solar Li abundance by a factor of 160 relative to the primordial solar system conrposition remains one of the most serious challenges of standart solar models. which. as illustrated iun Fig," This large depletion in the observed solar Li abundance by a factor of 160 relative to the primordial solar system composition remains one of the most serious challenges of standard solar models, which, as illustrated in Fig."
 l. destrov only a nüuor amount (0.06 dex) of Li 1981).," 1, destroy only a minor amount (0.06 dex) of Li ."
. It is important to note that modern models using updated OPAL opacities predict too much lithium destruction during the preauain sequence 2003).. at varlance with the observations of Li in solar-nass stars iu vouus open clusters (sec e.g.àY Fie.," It is important to note that modern models using updated OPAL opacities predict too much lithium destruction during the pre-main sequence , at variance with the observations of Li in solar-mass stars in young open clusters (see e.g. Fig."
 6)., 6).
 When the new low solar abundances are used the probloi is reduced 2006).. but in eeucral classical models have problems dealing with pre-main sequence convection. so other parameters are invoked to reduce the efficient lithium destruction during the preauaiu sequence 1998).," When the new low solar abundances are used the problem is reduced , but in general classical models have problems dealing with pre-main sequence convection, so other parameters are invoked to reduce the efficient lithium destruction during the pre-main sequence ."
. A comparison between the Stu aud solar analogs of one solar mass and solar metallicity shows the Sun to be ΠΠρου by a factor of 10., A comparison between the Sun and solar analogs of one solar mass and solar metallicity shows the Sun to be “lithium-poor” by a factor of 10.
 This apparent peculiaritv in the solu Li abundance. led to sugecst that the Sun iav be of dubious value for calibrating non-staudard models of Li depletion.," This apparent peculiarity in the solar Li abundance, led to suggest that the Sun may be of dubious value for calibrating non-standard models of Li depletion."
 However. Pasquini ct al. (," However, Pasquini et al. ("
1991) have shown that there are solar-tvpe stars that have a Li abundance as low as in the Sun.,1994) have shown that there are solar-type stars that have a Li abundance as low as in the Sun.
 Nevertheless.," Nevertheless,"
the scatter iu these values is still ouly a lew kilometers aud well below the resolutious of the ünages (compare to Table 1) . providing further confirmation that the fittiug procedures euiployec bere are robust.,"the scatter in these values is still only a few kilometers and well below the resolutions of the images (compare to Table \ref{obstab1}) ), providing further confirmation that the fitting procedures employed here are robust."
 Note that all the observatious give a values within a few silometers of 119.9I0 km. which is very ¢ose to exactly halfway between the inner edge of the Laplace Gap at 119.8[5 Kin aud the inier edge of the Laplace riuglet at 120.036 kin 2010).," Note that all the observations give $a$ values within a few kilometers of 119,940 km, which is very close to exactly halfway between the inner edge of the Laplace Gap at 119,845 km and the inner edge of the Laplace ringlet at 120,036 km \citep{Hedman10}."
. Turning to the eccentriciy and pericenter. we may uote tiat while the periceuter is always around 1505 [rom he Sun. both the pericenter location aud the eccentricity vary significantly [ron one obse'vatiou to another.," Turning to the eccentricity and pericenter, we may note that while the pericenter is always around $180^\circ$ from the Sun, both the pericenter location and the eccentricity vary significantly from one observation to another."
 The values of ae range from 9.5 kin to 25 kin. aud the periceuer locations deviate [rom the auti-5un clirection by up to 35°.," The values of $ae$ range from 9.5 km to 25 km, and the pericenter locations deviate from the anti-Sun direction by up to $35^\circ$."
 The rinelet therefore does 1οἱ 1nalutal La perfectly fixed orientation relative to the $LLL., The ringlet therefore does not maintain a perfectly fixed orientation relative to the Sun.
 Fiially. coisider the iiclination and the node estimates.," Finally, consider the inclination and the node estimates."
 Six of the seven estimates for a? [αἱ 1 a relatively narrow rauge of 2.3-3.2 kim., Six of the seven estimates for $ai$ fall in a relatively narrow range of 2.3-3.2 km.
 The one outlier is the 5.1 an estimate [from he Orbi 1.96 cdaa., The one outlier is the 5.1 km estimate from the Orbit 96 data.
 Howeve ais observation was mace while the spacecra[t was well above he riuD>oOjane. aud the ring opening angle chauged more over the course o“(hus observation han it auy ol the otliers (see Table 1)).," However, this observation was made while the spacecraft was well above the ringplane, and the ring opening angle changed more over the course of this observation than in any of the others (see Table \ref{obstab1}) )."
 Thus this measurement of the iclination may be 'egarcde« as suspect., Thus this measurement of the inclination may be regarded as suspect.
 Looking at the remaining data. the relatively simall scatter in a7 may imply |at the inclination of this ringlet does not vary much with time.," Looking at the remaining data, the relatively small scatter in $ai$ may imply that the inclination of this ringlet does not vary much with time."
 However. we also ind tha the node positions are very widely scattered.," However, we also find that the node positions are very widely scattered."
 This implies that this rinelet’s liue of nodes does not have a fixed orientation relative to the Sun., This implies that this ringlet's line of nodes does not have a fixed orientation relative to the Sun.
 The above data show that the riuglet's pericenter is on average anti-aligued. with the Sun. sugeMODesting that a force like solar radiation pressure is ii[Iueuciug the shape ancl orientation of this rinelet.," The above data show that the ringlet's pericenter is on average anti-aligned with the Sun, suggesting that a force like solar radiation pressure is influencing the shape and orientation of this ringlet."
 However. the eccentricity aud aliguimen ol this riuglet also vary significantly over time. and this indicates that the riuglet's dynamics are more complex than we might have expected.," However, the eccentricity and alignment of this ringlet also vary significantly over time, and this indicates that the ringlet's dynamics are more complex than we might have expected."
 Iu order to facilitate the interpretation of these data. we will," In order to facilitate the interpretation of these data, we will"
The chemically peculiar (CI) stars lying in the temperaure range νε< IxIx. are identified. by the presence of anomalously strong (and/or weak) absorpjon lines of certain heavy and rare carth clements in. t1011’ spectra.,The chemically peculiar (CP) stars lying in the temperature range $\le T_{eff} \le$ K are identified by the presence of anomalously strong (and/or weak) absorption lines of certain heavy and rare earth elements in their spectra.
 Two main classes of peculiarities are recognised among CP stars in the spectral range B5 to F5., Two main classes of peculiarities are recognised among CP stars in the spectral range B5 to F5.
 The inst class comprises the peculiar A tvpe stars. designate as “Ap. whose spectra are characterisecl by over-abuncdances of Si. Mn. Cr. Sr and rare earth elements.," The first class comprises the peculiar A type stars, designated as `Ap', whose spectra are characterised by over-abundances of Si, Mn, Cr, Sr and rare earth elements."
 The other class comprises the metalic line stars. designated as “Am stars. whose spectra are ¢jwacterised by an uncderabuncdance of Ca (and/or Se) andor an overabundance of the Fe eroup and heavier elements with respect to normal stars of same colour.," The other class comprises the metallic line stars, designated as `Am' stars, whose spectra are characterised by an underabundance of Ca (and/or Sc) and/or an overabundance of the Fe group and heavier elements with respect to normal stars of same colour."
 Am stars wil1 spectral types based on metallic lines and the Call Ix-line cillering bv five or more sub-classes are called classical Xam sars: otherwise they are called. marginal Am stars (clesignated as Am:)., Am stars with spectral types based on metallic lines and the II K-line differing by five or more sub-classes are called classical Am stars; otherwise they are called marginal Am stars (designated as `Am:').
 About of the stars in the lower classical instability strip with luminosities ranging from the zero age main, About of the stars in the lower classical instability strip with luminosities ranging from the zero age main
A prominent iron Ίνα line is present around 6.4 keV. llowever. its EW with respect to the Compton reflection component. being 610.2) eV. falls short of the expected one. which should be larger than 1 keV (seee.g.2)..,"A prominent iron $\alpha$ line is present around 6.4 keV. However, its EW with respect to the Compton reflection component, being $610^{+30}_{-50}$ eV, falls short of the expected one, which should be larger than 1 keV \citep[see e.g.][]{mbf96}."
 Two factors can contribute in reconciling the observed. value with the theoretical one., Two factors can contribute in reconciling the observed value with the theoretical one.
 One is a small inclination angle. since the iron line EW increases with this parameter.," One is a small inclination angle, since the iron line EW increases with this parameter."
" The other is iron underabundance: the dependence of EW with this parameter is almost linear for cle,«1 (7).", The other is iron underabundance: the dependence of EW with this parameter is almost linear for $A_{Fe}<1$ \citep{mbf96}.
 Both these requirements are satisfied bv the Compton reflection component properties. which arises in the same material: the value of £2 is consistent with low inclination angles (see Sect. 4.1.1))," Both these requirements are satisfied by the Compton reflection component properties, which arises in the same material: the value of $R$ is consistent with low inclination angles (see Sect. \ref{CR}) )"
 and the iron edge depth suggests an iron underabundance of a factor &0.82 (see Sect. 3))., and the iron edge depth suggests an iron underabundance of a factor $\simeq0.82$ (see Sect. \ref{analysis}) ).
 The iron Wea width is actually resolved in the EPIC pn spectrum. with e=32.11 eV. Even ifthis measure should be taken with care because it is well below the instrument resolution. it is interesting to note that it is fully compatible with the value found byChandra. whic 11s 0—258i eV. The observed iron line width can be in principle explained by several means.," The iron $\alpha$ width is actually resolved in the EPIC pn spectrum, with $\sigma=32^{+13}_{-14}$ eV. Even if this measure should be taken with care because it is well below the instrument resolution, it is interesting to note that it is fully compatible with the value found by, which is $\sigma=28^{+11}_{-7}$ eV. The observed iron line width can be in principle explained by several means."
 One clleet that should be in principle taken into account is that the neutral iron line is indeed composed of a doublet. Way at 6.404 keV and Kae at 6.391 keV. with a flux ratio of 2:1 (7).," One effect that should be in principle taken into account is that the neutral iron line is indeed composed of a doublet, $\alpha_1$ at 6.404 keV and $\alpha_2$ at 6.391 keV, with a flux ratio of 2:1 \citep{bea67}."
. The dillerence between the two lines. being lower than the spectral resolution of present X-ray detectors. is generally. ignored anc a weighted mean of 6.400. keV is Commonly used.," The difference between the two lines, being lower than the spectral resolution of present X-ray detectors, is generally ignored and a weighted mean of 6.400 keV is commonly used."
 ? already. noted that a fit with two Gaussian lines does not vary significantly the measure of the total line width in gratings data., \citet{yaq01} already noted that a fit with two Gaussian lines does not vary significantly the measure of the total line width in gratings data.
" We performed the same test with our high statistics EPIC pn spectrum. fitting the neutral iron line with à Ka, and à Kae. whose energies and relative intensities were freezed to the atomic values. while their width was constrained to vary together."," We performed the same test with our high statistics EPIC pn spectrum, fitting the neutral iron line with a $\alpha_1$ and a $\alpha_2$, whose energies and relative intensities were freezed to the atomic values, while their width was constrained to vary together."
 The resulting. common width is =3511 eV. in full agreement with the value found with a single Gaussianline.," The resulting, common width is $\sigma=35^{+14}_{-16}$ eV, in full agreement with the value found with a single Gaussian."
 On the other hand. since the line is) produced. by reprocessing from a Conipton-thick material. a Compton Shoulder (CS) is expected on theoretical grounds (7). as a result of Compton scattering of the line photons on the same material where they originate.," On the other hand, since the line is produced by reprocessing from a Compton-thick material, a Compton Shoulder (CS) is expected on theoretical grounds \citep{matt02}, as a result of Compton scattering of the line photons on the same material where they originate."
 Indeed. the iron line profiles of Circinus and NGC LOGS are successfully modelled with a narrow core and a CS both in and. spectra (22?272).. with a better statistical significance with respect to a single Gaussian line with a resolved width.," Indeed, the iron line profiles of Circinus and NGC 1068 are successfully modelled with a narrow core and a CS both in and spectra \citep{bianchi02,mbm03,matt04,ogle03}, with a better statistical significance with respect to a single Gaussian line with a resolved width."
 Therefore. another possibility is that the observed. ENIM is the result of a bad modelling of the overall profile. which does not include the CS.," Therefore, another possibility is that the observed FWHM is the result of a bad modelling of the overall profile, which does not include the CS."
 To paranietrise this component. we addec a further Gaussian line. with centroid energy at 6.3 keV anca=40 ον) (7)..," To parametrise this component, we added a further Gaussian line, with centroid energy at 6.3 keV and $\sigma=40$ eV \citep{matt02}."
 Phe CS is not required on statistical erouncls (Αν~O0 for 1 doo., The CS is not required on statistical grounds $\Delta\chi^2\simeq0$ for 1 d.o.f.
" less). but the CS Εικ. even if low. is consistent with production from a Compton-thick material. being 10.ορ,; that of the line core. and the iron line width is now unresolved (the upper limit being 50 eV)."," less), but the CS flux, even if low, is consistent with production from a Compton-thick material, being $10^{+9}_{-6}$ that of the line core, and the iron line width is now unresolved (the upper limit being 50 eV)."
 We hen tested if this model applies well also to the LETC data., We then tested if this model applies well also to the HETG data.
 However. this is not the case.," However, this is not the case."
 An unresolved ine core plus a CS gives a significantly wese fit with respect oa single. resolved Gaussian (N47=13 or the same d.o.f.).," An unresolved line core plus a CS gives a significantly worse fit with respect to a single, resolved Gaussian $\Delta\chi^2=13$ for the same d.o.f.)."
 On the other hand. if the line core width is left free to vary (with a best fit value ofa=29. 17). we get an upper limit to he CS flux of4.. still compatible with the expectations. oit not explaining the FWHLIAL of the line.," On the other hand, if the line core width is left free to vary (with a best fit value of $\sigma=29^{+10}_{-8}$ ), we get an upper limit to the CS flux of, still compatible with the expectations, but not explaining the FWHM of the line."
 A more likely possibility is that we are not observing a line from purely neutral iron., A more likely possibility is that we are not observing a line from purely neutral iron.
 Lo this is the case. the iron We line is expected to be actually the blend of a number of emission features from cilTerent ionic species. cach composed. by a doublet. whose transitions are separated. by energies of tens eV. Indeed. the best fit value of the centroid energy for the line in the EPIC pn spectrum suggests that iron is ionised in the range (2)..," If this is the case, the iron $\alpha$ line is expected to be actually the blend of a number of emission features from different ionic species, each composed by a doublet, whose transitions are separated by energies of tens eV. Indeed, the best fit value of the centroid energy for the line in the EPIC pn spectrum suggests that iron is ionised in the range \citep{house69}."
 However. the W3/ive Dux ratio is 0.12(07.," However, the $K\beta/K\alpha$ flux ratio is $0.12^{+0.06}_{-0.05}$."
 This value does not depend only on the ratio between the fluorescence vields of the two transitions (1:8). but also on the cüllerent total [ux of the continuum. ancl photoabsorption cross sections at. the energies of the two lines.," This value does not depend only on the ratio between the fluorescence yields of the two transitions (1:8), but also on the different total flux of the continuum and photoabsorption cross sections at the energies of the two lines."
 To account for all these ellects. ? used the ?. aand obtained. for neutral iron. a ratio of 0.155-0.160. depending on the inclination angle.," To account for all these effects, \citet{mbm03} used the \cite{basko78} and obtained for neutral iron a ratio of 0.155-0.160, depending on the inclination angle."
 Pherefore. the measured," Therefore, the measured"
((bound-Iree threshokl). producing a local wasxiniun in (he energy distribution of hydrogen models.,"(bound-free threshold), producing a local maximum in the energy distribution of hydrogen models."
 Furthermore. the predicted Ilime profile assuming a pure hydrogen atmosphere for LIES 1447 clearly rules out (his solution.," Furthermore, the predicted line profile assuming a pure hydrogen atmosphere for LHS 1447 clearly rules out this solution."
 The other two objects. E351—50 and WD 0227-444. are (oo cool to showHa... even if we assume a pure hyverogen composition.," The other two objects, $-$ 50 and WD $-$ 444, are too cool to show, even if we assume a pure hydrogen composition."
 Hence we must rely solely on (he fits to the energy distributions., Hence we must rely solely on the fits to the energy distributions.
" F351—50 represents an excellent example of a cool. pure hydrogen atmosphere white ναί,"," $-$ 50 represents an excellent example of a cool, pure hydrogen atmosphere white dwarf."
 The differences between the hydrogen and helium solutions are extreme in this case (A?=3.6 for the hydrogen fit as opposed to ~150 for the helium fit)., The differences between the hydrogen and helium solutions are extreme in this case $\chi^2=3.6$ for the hydrogen fit as opposed to $\sim 150$ for the helium fit).
 Our pure hydrogen fit. however. [ails to reproduce the observed flux at D within the uncertainties. and also at V to a lesser extent.," Our pure hydrogen fit, however, fails to reproduce the observed flux at $B$ within the uncertainties, and also at $V$ to a lesser extent."
 This discrepancy has been explained by BRL in terms of a missing opacity source in the ultraviolet of the pure hydrogen models. most likely due to a pseudo-continuum opacily originating from the Lyman edege (see 55.2.2 of BRL for a complete description). although this explanation has been challenged by WollEetal.(2002).," This discrepancy has been explained by BRL in terms of a missing opacity source in the ultraviolet of the pure hydrogen models, most likely due to a pseudo-continuum opacity originating from the Lyman edge (see 5.2.2 of BRL for a complete description), although this explanation has been challenged by \citet{wolff02}."
. Note also that the failure of the pure hydrogen models to match the observed [Iuxes in this particular region of the energv distribution is most likely al the origin of the peculiar solution obtained for F351—50 by Oppenheimeretal.(2001b)— Toy=2844 Ix. logg=6.5. and N(Ile)/N(11)=0 (see their Fig.," Note also that the failure of the pure hydrogen models to match the observed fluxes in this particular region of the energy distribution is most likely at the origin of the peculiar solution obtained for $-$ 50 by \citet{opp01b} – $\Te=2844$ K, $\logg=6.5$, and $\nhe=0$ (see their Fig."
 9). based solely on a spectrum covering the region between 0.4 and 1. jam. Finally. WD 0227—444. shown at the bottom of Figure 3.. represents a good example of a cool white dwarf with a pure helium atmosphere (47=8.9 as opposed to 25.1 for the hvdrogen fit).," 9) – based solely on a spectrum covering the region between 0.4 and 1 $\mu$ m. Finally, WD $-$ 444, shown at the bottom of Figure \ref{fg:f3}, represents a good example of a cool white dwarf with a pure helium atmosphere $\chi^2=8.9$ as opposed to 25.1 for the hydrogen fit)."
 In this case. only the observed [αν at {1 is not matched by the helium model. within (he uncertainties.," In this case, only the observed flux at $H$ is not matched by the helium model, within the uncertainties."
 In contrast. six out (he seven bands used in our filling procedure are not matched properly by the hydrogen model.," In contrast, six out the seven bands used in our fitting procedure are not matched properly by the hydrogen model."
 The atmospheric parameters Zar. logg. ancl atmospheric composition (1l or He) for the 32 objects listed in Table 1 are given in Table 2 together with the caleulated stellar mass. absolute visual magnitude. Iuminositv. distance. and white dwarf cooling age.," The atmospheric parameters $\Te$, $\logg$, and atmospheric composition (H or He) for the 32 objects listed in Table 1 are given in Table 2 together with the calculated stellar mass, absolute visual magnitude, luminosity, distance, and white dwarf cooling age."
 The latter is obtained from the theoretical cooling sequences described above., The latter is obtained from the theoretical cooling sequences described above.
 A value of logg=8.0 was assumed [or all stars expect where noted: photometric distances are given [or (hese objects., A value of $\logg=8.0$ was assumed for all stars expect where noted; photometric distances are given for these objects.
 Three objects in the OILDIIS sample stood out in our analvsis. WD 2356-209 and LIIS 1402 labeled in Figure 2.. which had to be analvzed in greater detail.," Three objects in the OHDHS sample stood out in our analysis, WD $-$ 209 and LHS 1402 labeled in Figure \ref{fg:f2}, which had to be analyzed in greater detail."
 Wediscuss them in the next two sections., Wediscuss them in the next two sections.
 The third object is the extremely massive DA white dwarl LIIS 4033 analvzed in detail by Dahnetal. (2004).., The third object is the extremely massive DA white dwarf LHS 4033 analyzed in detail by \citet{dahn04}. .
For each trial we compute the power spectrum of a helt curve constructed with the same observation times as the original but with the fux values for cach time drawn randomly (with replacement) frou the original flux values.,For each trial we compute the power spectrum of a light curve constructed with the same observation times as the original but with the flux values for each time drawn randomly (with replacement) from the original flux values.
 A fit to the lieh probability cud of the resulting distribution of maxi powers indicates that the muuber of independent frequcucics is ~ 9 times the munhber of data points (94961)., A fit to the high probability end of the resulting distribution of maximum powers indicates that the number of independent frequencies is $\sim$ 9 times the number of data points $\times$ 961).
 This results in a false alarin probability of ~ 3410 7. which is a factor of 10 stnaller than an extrapolation of the Moute Carlo results.," This results in a false alarm probability of $\sim$ $\times$ $^{-8}$, which is a factor of 10 smaller than an extrapolation of the Monte Carlo results."
 Siuularly. the probability that a peak as large asthe one neax X» pz (~ 16) will occur there by chance is  X10 7. 29," Similarly, the probability that a peak as large asthe one near 2950 $\mu$ Hz $\sim$ 16) will occur there by chance is $\sim$ $\times$ $^{-8}$."
5The probability of a chance occurrence of a oeak at least this hieh somewhere iu the frequency rauge is ~ according to the analytic calculation. OY  according to the Monte Carlo results.," The probability of a chance occurrence of a peak at least this high somewhere in the frequency range is $\sim$ according to the analytic calculation, or $\sim$ according to the Monte Carlo results."
 To characterize these variations. we model them as a suni of sine waves aud determine the best-fit amplitude. requencev. and plase by uou-linear least-squares fits to he data. which we performed using both Period1 (Lenz&Breeer2005) aud AIPFIT (Alarkwardt2009).," To characterize these variations, we model them as a sum of sine waves and determine the best-fit amplitude, frequency, and phase by non-linear least-squares fits to the data, which we performed using both Period04 \citep{len05} and MPFIT \citep{mar09}."
.. The argest peak in the 2008 July amplitude spectrum is at approximately half the 2950 pz frequency but is rot significantly preseut on the subsequent nights. so we include this extra. low-frequency component in the fit for tha uieht ouly and note that fitting without it does not vield a significant difference.," The largest peak in the 2008 July amplitude spectrum is at approximately half the 2950 $\mu$ Hz frequency but is not significantly present on the subsequent nights, so we include this extra, low-frequency component in the fit for that night only and note that fitting without it does not yield a significant difference."
 We list the best-fit parameters aud their formal errors iu Table 2 , We list the best-fit parameters and their formal errors in Table \ref{table:fit1}. .
Doth aud hhave nou-sinusoidal light curves because of the presence, Both and have non-sinusoidal light curves because of the presence
"derived jointly with a y minimization algorithm: therefore vsini and the angles 5, and p, enter as new parameters.",derived jointly with a ${\chi}^2$ minimization algorithm; therefore $v sin i$ and the angles $\beta_{p}$ and $\beta_{s}$ enter as new parameters.
 5 is the angle between the stellar spi axis projected onto the plane of the sky and the orbital spi axis projected onto the plane of the sky., $\beta$ is the angle between the stellar spin axis projected onto the plane of the sky and the orbital spin axis projected onto the plane of the sky.
 B 00* would indicate a projected rotation axis perpendicular to the orbital plane. whereas B = 90° would indicate that the rotation axis lies in. the orbital plane.," $\beta$ $^{\circ}$ would indicate a projected rotation axis perpendicular to the orbital plane, whereas $\beta$ = $^{\circ}$ would indicate that the rotation axis lies in the orbital plane."
 The longitude of the ascending node of the orbit (Q) is not known. resulting in an ambiguity in the sign of the angle 5.," The longitude of the ascending node of the orbit $\Omega$ ) is not known, resulting in an ambiguity in the sign of the angle $\beta$."
 In our definition. a positive B indicates that the RM effect. integrated over the complete eclipse. would give a positive residual 1 radial-velocity.," In our definition, a positive $\beta$ indicates that the RM effect, integrated over the complete eclipse, would give a positive residual in radial-velocity."
 In that case. the companion spends a longer time in front of the stellar surface moving towards us. tha that part of the stellar surface that is moving away from us.," In that case, the companion spends a longer time in front of the stellar surface moving towards us, than that part of the stellar surface that is moving away from us."
 The linear limb-darkening coefficients of both stars are fixed during the fits to 0.6 as they are only weakly constrained by our fits., The linear limb-darkening coefficients of both stars are fixed during the fits to $0.6$ as they are only weakly constrained by our fits.
 In the y7 fit. all data points out of eclipse and all data points during the primary and secondary eclipses are fitted simultaneously.," In the ${\chi}^2$ fit, all data points out of eclipse and all data points during the primary and secondary eclipses are fitted simultaneously."
 Our radial-velocity data points and the best fit can be seen in Fig., Our radial-velocity data points and the best fit can be seen in Fig.
 4. for the complete orbit including the two eclipses., \ref{fig:orbit} for the complete orbit including the two eclipses.
 Fig., Fig.
 5 shows the orbit from a point above the orbital plane., \ref{fig:orbit_plane} shows the orbit from a point above the orbital plane.
 Figs., Figs.
 6 and 7. display the radial-velocities of the primary and secondary stars during and around their eclipses., \ref{fig:beta_primary} and \ref{fig:beta_secondary} display the radial-velocities of the primary and secondary stars during and around their eclipses.
 The best fit parameters are listed in the third column of Table 2. with their 1-c uncertainties. as derived from the y fit.," The best fit parameters are listed in the third column of Table \ref{tab:fit} with their $\sigma$ uncertainties, as derived from the ${\chi}^2$ fit."
 All radial-velocity measurements are given with their epochs in Table Al in the Appendix., All radial-velocity measurements are given with their epochs in Table \ref{tab:data} in the Appendix.
 The radial-velocities directly after the primary eclipse. and the data point at a phase of 0.7. have been extracted using a Gaussian fit to the BF of one star after the subtraction of the other star in the spectrum: however. there is no significant change in the derived parameters if these data points are omitted.," The radial-velocities directly after the primary eclipse, and the data point at a phase of 0.7, have been extracted using a Gaussian fit to the BF of one star after the subtraction of the other star in the spectrum; however, there is no significant change in the derived parameters if these data points are omitted."
 The average uncertainties 1n the radial-velocities. out of the eclipses. are kkm/s for the primary and kkm/s for the secondary.," The average uncertainties in the radial-velocities, out of the eclipses, are km/s for the primary and km/s for the secondary."
 The uncertainties of those radial-velocity measurements that were determined from the centers of the lines. we estimated to be kkm/s for the primary and kkm/s for the secondary.," The uncertainties of those radial-velocity measurements that were determined from the centers of the lines, we estimated to be km/s for the primary and km/s for the secondary."
 This was calculated by comparing them with the velocities obtained from a Gaussian fit outside the eclipses., This was calculated by comparing them with the velocities obtained from a Gaussian fit outside the eclipses.
" In our second approach. to derive information about the binary orbit and the orientation of the rotation axes. we simulated the shape of the BF of the two stars. as governed by the orbital motion, stellar rotation. orientation of the stellar spin axes. velocity fields on the stellar surface. limb-darkening and. in case of a measurement taken during an eclipse. the fraction of the stellar disk which is covered by the companion at the time of the measurement."," In our second approach, to derive information about the binary orbit and the orientation of the rotation axes, we simulated the shape of the BF of the two stars, as governed by the orbital motion, stellar rotation, orientation of the stellar spin axes, velocity fields on the stellar surface, limb-darkening and, in case of a measurement taken during an eclipse, the fraction of the stellar disk which is covered by the companion at the time of the measurement."
 These simulated BFs were subsequently compared to the measured BFs. and the relevant parameters determined.," These simulated BFs were subsequently compared to the measured BFs, and the relevant parameters determined."
 With this method. we not only use the first moment of the stellar absorption lines during the analysis. but we utilize the complete BF and its change over the course of the eclipse as a diagnostic tool for determining the orbital and. stellar parameters.," With this method, we not only use the first moment of the stellar absorption lines during the analysis, but we utilize the complete BF and its change over the course of the eclipse as a diagnostic tool for determining the orbital and stellar parameters."
 We simulated the rotation profile of the two stars 1n the VII43CCyg system with a few thousand elements across the visible half-spheres with equal surface brightness., We simulated the rotation profile of the two stars in the Cyg system with a few thousand elements across the visible half-spheres with equal surface brightness.
 We included linear limb-darkening. asymmetric macro-turbulence and solar-Itke differential rotation (See Gray 2005).," We included linear limb-darkening, asymmetric macro-turbulence and solar-like differential rotation (See \citeauthor{Gray2005} \citeyear{Gray2005}) )."
 Accordingly. besides the parameters which were fitted in Section. 3.3.1.. another parameter. the Gaussian width of the macro-turbulence. is required.," Accordingly, besides the parameters which were fitted in Section \ref{sect:center}, another parameter, the Gaussian width of the macro-turbulence, is required."
 We kept the σ of the Gaussia1 for the tangential and radial-velocity fields fg; and i covered surface fraction equal. as the quality of the fit did not improve by including them as free parameters.," We kept the $\sigma$ of the Gaussian for the tangential and radial-velocity fields $\zeta_{RT}$ and their covered surface fraction equal, as the quality of the fit did not improve by including them as free parameters."
 The inclination of the orbit and the sizes of the (vo stars can be varied in our fits., The inclination of the orbit and the sizes of the two stars can be varied in our fits.
 However. in the final fits presented here. we used a linear limb-darkening coefficient of 0.6. solid body rotation. and the inclination and radii from the literature (see also Section 4)).," However, in the final fits presented here, we used a linear limb-darkening coefficient of 0.6, solid body rotation, and the inclination and radii from the literature (see also Section \ref{sect:discussion}) )."
 Each exposure of VII43CCyg had a duration of = mmin., Each exposure of Cyg had a duration of $\approx$ min.
 Provided that one takes the photon midpoint of the observation. this poses no serious problem for the analysis of the center of gravity of the lines.," Provided that one takes the photon midpoint of the observation, this poses no serious problem for the analysis of the center of gravity of the lines."
 However. looking at the line profile. one has to take two effects into account.," However, looking at the line profile, one has to take two effects into account."
 First. the radial-velocities of the two stars change during a mmin exposure.," First, the radial-velocities of the two stars change during a min exposure."
 This effect is strongest in our data set at the time of the primary eclipse. where the change in radial-velocity ts x 3kkm/s during one exposure.," This effect is strongest in our data set at the time of the primary eclipse, where the change in radial-velocity is $\approx$ $/$ s during one exposure."
 This artificially widens a BF taken during a mmin exposure relative to à BF which would have been taken instantaneously., This artificially widens a BF taken during a min exposure relative to a BF which would have been taken instantaneously.
 During the eclipses. a second effect becomes important: the coverage of the eclipsed stars can change considerably duri5 20-min.," During the eclipses, a second effect becomes important; the coverage of the eclipsed stars can change considerably during 20-min."
 Again. this effect is stronger during the primary eclipse than during the secondary eclipse. as the primary eclipse is deeper and shorter.," Again, this effect is stronger during the primary eclipse than during the secondary eclipse, as the primary eclipse is deeper and shorter."
 Sine we need an exposure over several ninutes to receive a sufficiently, Sine we need an exposure over several minutes to receive a sufficiently
Horizontal branch (IB) stars are core helium-burning stars hat have evolved olf the red giant. branch (ROB).,Horizontal branch (HB) stars are core helium-burning stars that have evolved off the red giant branch (RGB).
 In elobular clusters (GC's) they are so named because of their appearance on optical colow-magnituce diagrams. where hey trace a horizontal path blueward of the RGD.," In globular clusters (GCs) they are so named because of their appearance on optical colour-magnitude diagrams, where they trace a horizontal path blueward of the RGB."
 These stars are thought to have similar core masses of ~O.5AL.., These stars are thought to have similar core masses of $\sim$ $M_{\odot}$.
 Their location alone the LB is therefore. determined wimarilv by the mass ancl opacity of their thin. stellar envelopes., Their location along the HB is therefore determined primarily by the mass and opacity of their thin stellar envelopes.
 To explain the observed LLB stars in GC's. it is thought that substantial mass loss must occur during he star's ROB phase (e.g.2)..," To explain the observed HB stars in GCs, it is thought that substantial mass loss must occur during the star's RGB phase \citep[e.g.][]{Rood73}."
 The HB morphologies of cdilferent Galactic GC's are known to vary significantly., The HB morphologies of different Galactic GCs are known to vary significantly.
 This suggests that the cluster properties have a significant effect on the evolution of stars onto the LIB., This suggests that the cluster properties have a significant effect on the evolution of stars onto the HB.
" The. metallicity ofa GC has long been proposed. as the ""first. parameter related to the morphology of a GC's LB. with metal poor clusters often having blucr LIB stars than metal rich CC's (c.g.77)."," The metallicity of a GC has long been proposed as the `first parameter' related to the morphology of a GC's HB, with metal poor clusters often having bluer HB stars than metal rich GCs \citep[e.g.][]{Sandage60,Dorman95}."
 However. the metallicity of a cluster alone is insullicient to explain the HB morphology in all Galactic GCs.," However, the metallicity of a cluster alone is insufficient to explain the HB morphology in all Galactic GCs."
 Firstly. the clusters which host the bluest LIBs actually lave intermediate metallicities (e.g.2)..," Firstly, the clusters which host the bluest HBs actually have intermediate metallicities \citep[e.g.][]{OConnell99}."
 Also. some clusters are observed to have very. dilferent. HB. morphologies. but similar metallicities (e.g.NGC288and362:?)..," Also, some clusters are observed to have very different HB morphologies, but similar metallicities \citep[e.g. NGC~288 and NGC~362:][]{Bellazzini01}."
" This jas led to the notorious search for à ""second. parameter to describe HB. morphology.", This has led to the notorious search for a `second parameter' to describe HB morphology.
 Many. clilferent parameters have oen proposed as à second parameter. all of which may xav some role.," Many different parameters have been proposed as a second parameter, all of which may play some role."
" Vhese include the clusters age (e.g.οτι. its wlium abundance (e.g.οτι, the stellar core rotation 77). mass loss on the RGB (e.g.2). and cluster core density (?).."," These include the cluster's age \citep[e.g.][]{Demarque88,Lee94}, its helium abundance \citep[e.g.][]{Sweigart97,Lee05}, the stellar core rotation \citep[e.g.][]{Mengel76,Peterson85}, , mass loss on the RGB \citep[e.g.][]{Rood73} and cluster core density \citep{Buonanno97}."
 Despite decades of research. the relative effects of these different parameters are vet to be Lully understood. (fordiscussionofproposedsecond.parameters.seee.g. 2?)..," Despite decades of research, the relative effects of these different parameters are yet to be fully understood \citep[for a discussion of proposed second parameters, see e.g. ][]{Fusi_Pecci97,Catelan09}."
 The bluest LIB stars observed in GC's are the blue tail’ or 'extreme-HD (IIB) stars., The bluest HB stars observed in GCs are the `blue tail' or `extreme-HB' (EHB) stars.
" These are analogues to the subedwart D. (sd) stars observed in the Galactic Ποιά and are considered here to to be those stars with Ty, 220.0001 (?).."," These are analogues to the subdwarf B (sdB) stars observed in the Galactic field and are considered here to to be those stars with $_{eff}>$ 20,000K \citep{Brown01}."
" These stars have become increasingly important. as the leading contenders in explaining the ""ultraviolet. excess? (UVX) observed in elliptical galaxies (c.g.222)..."," These stars have become increasingly important as the leading contenders in explaining the `ultraviolet excess' (UVX) observed in elliptical galaxies \citep[e.g.][]{Code69,Dorman95,OConnell99}."
 Because of their high clleetive temperatures. they are relatively faint at optical wavelengths but make an important. often dominating. contribution to the FUY emission. [rom old stellar populations.," Because of their high effective temperatures, they are relatively faint at optical wavelengths but make an important, often dominating, contribution to the FUV emission from old stellar populations."
 ELLB stars are observed in both metal rich and metal poorclusters (e.g.27). and it is likely that," EHB stars are observed in both metal rich and metal poorclusters \citep[e.g.][]{DCruz96,Rich97} and it is likely that"
"isophotal radius. velocity dispersion and luminosity of7x0.5 kkpc (for⋅ Lg,c1.5.10⊥↴""""eergss 1 72 2 7). az4565kms⋅1. and Lg,~−2.101T cress1 42 with uncertainties from the lens model ancl emission. line fits of 40% in radius. in dispersion and in Iuminositv.","isophotal radius, velocity dispersion and luminosity of$R\simeq0.5$ kpc (for $L_{H\alpha}>1.5\times10^{-16}$ $^{-1}$ $^{-2}$ $^{-2}$ ), $\sigma\simeq 45-65 \kms$, and $L_{H\alpha}\sim2\times10^{-17}$ $^{-1}$ $^{-2}$, with uncertainties from the lens model and emission line fits of $\sim 40$ in radius, in dispersion and in luminosity."
" This implies a dynamical mass of the Lla cbumps [rom the velocity dispersion of Aj,=2+110""M. assuming uniform density. and a La star-formation rate of SER =2040.7Move |."," This implies a dynamical mass of the $\alpha$ clumps from the velocity dispersion of $M_{dyn}=2\pm1\times10^9\Msol$ assuming uniform density, and a $\alpha$ star-formation rate of SFR $=2.0\pm0.7\Msolyr$ ."
 The velocity Ποια shows a smooth velocity gradient around the ring with peak-to-peak amplitude 2V;sin;=1202-10kms|.," The velocity field shows a smooth velocity gradient around the ring with peak-to-peak amplitude $2\,V_c\,\sin{i}=120\pm10\,\kms$."
 Using the axial ratio from the optical imaging.oὃν we estimate cos?=0.050.1 which suggests à rotationally supported mass A;=RVO/C22140.8107AL. within a radius of Okkpe.," Using the axial ratio from the optical imaging, we estimate $\cos{i}=0.6\pm0.1$ which suggests a rotationally supported mass $M_{dyn}=R\,V_c^2/G=2.1\pm0.8\times10^{10}\,M_{\odot}$ within a radius of kpc."
 Lf the galaxy has a uniform spherical mass distribution. then the observed clispersion suggests a larger dvnanical mass of 5«LOMM.," If the galaxy has a uniform spherical mass distribution, then the observed dispersion suggests a larger dynamical mass of $5\times10^{10}\Msol$."
 The resolved. star forming regions in the ring thus account for about of the integrated Lo Bux and of the mass interior to the ring., The resolved star forming regions in the ring thus account for about of the integrated $\alpha$ flux and of the mass interior to the ring.
 The integrated La emission. line Luxsuggests a star- rate of SER=7£2Aleve te with the two," The integrated $\alpha$ emission line fluxsuggests a star-formation rate of $7\pm2 \Msolyr$ , with the two"
field so that we could get the maximum difference in c.,field so that we could get the maximum difference in $c$.
" We keep the same magnetic field configuration for Model Set 3, but here we exchange a and c using the measured relation 2))."," We keep the same magnetic field configuration for Model Set 3, but here we exchange $a$ and $c$ using the measured relation )."
 In this Model Set we fix the quantity a+2c., In this Model Set we fix the quantity $a+2c$.
 This allows us to hold the magnetic field fixed while going from a hadronic-like CR model (c= to a reacceleration model (a— 0)., This allows us to hold the magnetic field fixed while going from a hadronic-like CR model $c=0$ ) to a reacceleration model $a=0$ ).
" We designed this Model0) Set to verify that our results are robust to even exchanges of a and c using the measured relation, which they should be."," We designed this Model Set to verify that our results are robust to even exchanges of $a$ and $c$ using the measured relation, which they should be."
" InModel Set 4 we fix b=1.0, a=0.0, and c=0.7 and examine the extreme allowed average magnetic field."," InModel Set 4 we fix $b=1.0$, $a=0.0$, and $c=0.7$ and examine the extreme allowed average magnetic field."
" We chose these values of b, a, and c such that we could get the maximum change in (B)."," We chose these values of $b$ , $a$, and $c$ such that we could get the maximum change in $\aveb$ ."
 We repeat this test, We repeat this test
 As discussed in Nolta (2009). the WMAPS estimated value of 24) using only V and W data is 1056—40% (depending on the choice of Galactic mask} lower than when using Q.V and W data. although the uncertainty in the former case is ~3 times the uncertainty in the latter.," As discussed in \cite{2009ApJS..180..296N}, the WMAP5 estimated value of $A_0$ using only V and W data is $\sim 10\% -40\%$ (depending on the choice of Galactic mask) lower than when using Q,V and W data, although the uncertainty in the former case is $\sim 3$ times the uncertainty in the latter."
 Therefore. using information from Q band may lead to a slight overestimate of ely at 60.90 GHz.," Therefore, using information from Q band may lead to a slight overestimate of $A_0$ at $60 -90$ GHz."
 As pointed out by Huffenbergeretal.(2006.2008) the point sources correction and the way it is treated in the likelihood function affects the determination of the cosmological parameters. in particular of the slope of the power spectrum of scalar perturbations. ns.," As pointed out by \cite{2006ApJ...651L..81H,2008ApJ...688....1H} the point sources correction and the way it is treated in the likelihood function affects the determination of the cosmological parameters, in particular of the slope of the power spectrum of scalar perturbations, $n_s$."
 To test the effect. if any. of our approach," To test the effect, if any, of our approach"
C(=2mnghrHp (see Fig. 39. ,"$\ell_c=2 \pi r_{\rm cz} / \alpha H_P$ (see Fig. \ref{fig:zc2I}) ),"
which temporarily favors the generation of waves (see models 0 to 2)., which temporarily favors the generation of waves (see models 0 to 2).
 However. with further retraction of the envelope. the thermal diffusivity Ay at the convective boundary increases. which favors the disappearance of low-frequency and/or large degree waves (see Fig.," However, with further retraction of the envelope, the thermal diffusivity $K_T$ at the convective boundary increases, which favors the disappearance of low-frequency and/or large degree waves (see Fig."
 5. and Eq. 7)., \ref{fig:spectreI} and Eq. \ref{optdepth}) ).
 This explains the reduction of the total energy flux at the end of the PMS (models 3 to 6) and on the MS., This explains the reduction of the total energy flux at the end of the PMS (models 3 to 6) and on the MS.
 During the pre-main sequence. an SLO exists only for a very short period of time.," During the pre-main sequence, an SLO exists only for a very short period of time."
 The evolution of the SLO for models ] ard 2 is shown in Fig. 9..," The evolution of the SLO for models 1 and 2 is shown in Fig. \ref{fig:slo},"
 where the successive profiles are separated by 5 years., where the successive profiles are separated by 5 years.
 The larger amplitude in model | compared to model 2 can be understood in terms of thermal diffusivity: since Ay increases with radius. as the top of the radiative zone migrates outwards with evolution. waves are dissipated more efficiently in model 2 than in model 1.," The larger amplitude in model 1 compared to model 2 can be understood in terms of thermal diffusivity: since $K_T$ increases with radius, as the top of the radiative zone migrates outwards with evolution, waves are dissipated more efficiently in model 2 than in model 1."
 In model 3. damping ts so large that the formation of a SLO is no longer possible. and all low-frequency waves are dissipated in model 4.," In model 3, damping is so large that the formation of a SLO is no longer possible, and all low-frequency waves are dissipated in model 4."
 The turbulent diffusion coefficient associated with this SLO. D (see Eq.," The turbulent diffusion coefficient associated with this SLO, $D$ (see Eq."
 10. and the discussion around it). is shown in Fig.," \ref{maeder95}
 and the discussion around it), is shown in Fig."
 10. às a function of reduced radius., \ref{fig:difPMS} as a function of reduced radius.
 While the shear is much stronger in model 1 than in model 2. the magnitude of D and width of the turbulent regions are quite similar.," While the shear is much stronger in model 1 than in model 2, the magnitude of $D$ and width of the turbulent regions are quite similar."
 This is because the reduced shear is compensated for by the larger thermal diffusivity in model 2 .EEq. 10))., This is because the reduced shear is compensated for by the larger thermal diffusivity in model 2 Eq. \ref{maeder95}) ).
 In both models. D ts relatively high just below the convective envelope. and it drops rapidly with depth.," In both models, $D$ is relatively high just below the convective envelope, and it drops rapidly with depth."
 It is negligible in the region with a temperature Tom2.5x 10°K) high enough for efticient lithium burning.," It is negligible in the region with a temperature $T \apprge 2.5 \times 10^6\,{\rm K}$ ) high enough for efficient lithium burning."
 This temperature is attained at reduced radii of r/R~0.43 and 0.41 in models | and 2 respectively., This temperature is attained at reduced radii of $r/R \sim 0.43$ and $0.41$ in models 1 and 2 respectively.
 This implies that. for this mass. the transport of elements through wave-inducec turbulence has no impact on the pre-main sequence evolutio of the surface abundance of the light elements such as lithium. beryllium. or boron (the latest two elements burning at an eve higher).," This implies that, for this mass, the transport of elements through wave-induced turbulence has no impact on the pre-main sequence evolution of the surface abundance of the light elements such as lithium, beryllium, or boron (the latest two elements burning at an even higher."
 The role of IGWs is more uncertain for the rotatio profile., The role of IGWs is more uncertain for the rotation profile.
 In the case of rapid rotation (v.=100kms7'). the initial rotation profile is deleted by meridional circulatio in a fraction of the main-sequence lifetime.," In the case of rapid rotation $v \apprge 100\,{\rm km\,s^{-1}}$ ), the initial rotation profile is deleted by meridional circulation in a fraction of the main-sequence lifetime."
 In à. slow rotator. however. this is not the case. and the memory of angular momentum transport by IGWs will influencemixing’.," In a slow rotator, however, this is not the case, and the memory of angular momentum transport by IGWs will influence."
.. Detailed evolutionary calculations of angular momentum in the presence of internal waves should thus take this evolutionary phase into account especially in the case of slow rotators (Charbonnel Talon. in preparation).," Detailed evolutionary calculations of angular momentum in the presence of internal waves should thus take this evolutionary phase into account especially in the case of slow rotators (Charbonnel Talon, in preparation)."
Taking the fourth differences amplifies the signals enough to be able to isolate them. keeping the errors al a manageable level (see Basu. Antia Narasimha 1994 for a cliscussion).,"Taking the fourth differences amplifies the signals enough to be able to isolate them, keeping the errors at a manageable level (see Basu, Antia Narasimha 1994 for a discussion)."
" The fourth differences can be fitted to the finetional form given by Basu (1997): where The coefficients a, αι define an overall smooth term. as ay. 7, and ο define the oscillatory contribution due to the IIe II ionization zone ancl the remaining terms define the oscillatory contribution Ilrom the CZ base."," The fourth differences can be fitted to the functional form given by Basu (1997): where The coefficients $a_1$ $a_4$ define an overall smooth term, $a_5$ $a_8$, $\tau_a$ and $\psi_a$ define the oscillatory contribution due to the He II ionization zone and the remaining terms define the oscillatory contribution from the CZ base."
" The frequencies of the two components. 7,ha and τι are approximately (he acoustic depths of the He IE ionization zone and the CZ base respectively. but thev also include a contribution from the frequency dependent part of the phases ο and ey which is not taken into account explicitly."," The frequencies of the two components, $\tau_a$ and $\tau_b$ are approximately the acoustic depths of the He II ionization zone and the CZ base respectively, but they also include a contribution from the frequency dependent part of the phases $\psi_a$ and $\psi_b$ which is not taken into account explicitly."
 The coefficients in (his expression are determined by a least squares fit to the fourth dillerences., The coefficients in this expression are determined by a least squares fit to the fourth differences.
 We use oscillation Ireeuencies determined bx the Michelson Doppler Imager (MDI) on board the SOLIIO spacecraft. as well as data from the Global Oscillations Network Group (GONG).," We use oscillation frequencies determined by the Michelson Doppler Imager (MDI) on board the SOHO spacecraft, as well as data from the Global Oscillations Network Group (GONG)."
 We have used 38 data sets [roin MDI (Schou 1999). each covering a period of 72 davs starting from 1996 May. 1 and ending on 2004 March 19.," We have used 38 data sets from MDI (Schou 1999), each covering a period of 72 days starting from 1996 May 1 and ending on 2004 March 19."
 We use 28 data sets [rom GONG 1 et al., We use 28 data sets from GONG (Hill et al.
 1996). each covering a period of 108 davs. starting Irom 1995 May. 7 and ending on 2003 August 16.," 1996), each covering a period of 108 days, starting from 1995 May 7 and ending on 2003 August 16."
 As a measure of the solar activity for each data set. we use the mean radio [hix at 10.7 cm curing the time interval covered by each data set as obtained from the US National Geophysical Data Center (wvww.nedc.noaa.gov/stp/stp.html).," As a measure of the solar activity for each data set, we use the mean radio flux at 10.7 cm during the time interval covered by each data set as obtained from the US National Geophysical Data Center (www.ngdc.noaa.gov/stp/stp.html)."
" For each data set we use modes with (wo sets of degrees: one set consists of modes with degree 5«€<25 and we refer to this as the low-degree set. the second set consists of all modes with 6x60 that have their lower turning point. r,>0.715 R. and we refer to this as the intermediate-degree set."," For each data set we use modes with two sets of degrees; one set consists of modes with degree $5\le\ell\le25$ and we refer to this as the low-degree set, the second set consists of all modes with $\ell\le60$ that have their lower turning point, $r_t > 0.715$ $_\odot$ and we refer to this as the intermediate-degree set."
 The lowest degree modes (hat satisfies this criterion for our frequency range of 2-3.5mIlz is (=30., The lowest degree modes that satisfies this criterion for our frequency range of 2-3.5mHz is $\ell=30$.
 Very. lew modes below { of 42 satisfy. Ulis criterion., Very few modes below $\ell$ of 42 satisfy this criterion.
 The first set of modes sample both the CZ base and the He II ionization zone and hence we fit the entire form given by Eq. (2)).," The first set of modes sample both the CZ base and the He II ionization zone and hence we fit the entire form given by Eq. \ref{eq:func}) ),"
 while modes in the second set saniple only the He I] ionization zone and hence we fit the mode only to the first (wo terms of Eq. (2))., while modes in the second set sample only the He II ionization zone and hence we fit the mode only to the first two terms of Eq. \ref{eq:func}) ).
 We do not use modes with £«5 because (μον have larger errors., We do not use modes with $\ell < 5$ because they have larger errors.
 We do not use modes with degree, We do not use modes with degree
lus them data extends to Pos173.,thus their data extends to $\Gamma > 173$.
 Error bars are displaved when there was sufficient information in the iterature to determine them., Error bars are displayed when there was sufficient information in the literature to determine them.
 Between these results. there is disagreemoeut on the level.," Between these results, there is disagreement on the level."
 The solid line displays he analytic result of the OCT (no-screeuiug) lianit of he Yukawa system as calculated by Muillo(2000)., The solid line displays the analytic result of the OCP (no-screening) limit of the Yukawa system as calculated by \citet{mur00}.
". The dotted line shows our fit to the data that reproduces the reduced viscosity προς, from D, of equation (1)) and he Stokes-Eiusteim relation in the regine P.«173 while approximately fitting the data in the of Tanaka Ichimazu or Doss173.", The dotted line shows our fit to the data that reproduces the reduced viscosity implied from $D_s$ of equation \ref{eq:diff}) ) and the Stokes-Einstein relation in the regime $\Gamma < 173$ while approximately fitting the data in the of Tanaka Ichimaru for $\Gamma \gg 173$.
 This fit is given by The first term im equation (5)) is the relation between g aud D interred from equation (1))., This fit is given by The first term in equation \ref{eq:etaglfit}) ) is the relation between $\eta^*$ and $\Gamma$ inferred from equation \ref{eq:diff}) ).
 We utilize this fit to explore the consequences of WD interiors not crystallizing but imstead undergoing a transition to a elassy state., We utilize this fit to explore the consequences of WD interiors not crystallizing but instead undergoing a transition to a glassy state.
 With all of these uncertainties in mind. we follow DIIOL and take D=Dy. as our fiducial best euess for the diffusion coefficient.," With all of these uncertainties in mind, we follow BH01 and take $D= D_s$ as our fiducial best guess for the diffusion coefficient."
" But. as we consider the actual value of D to be uncertain eveu in the case of diffuing in a single componcut plasma. we carry out caleulatious with D=D,. 5D.. and LOD, to determine to what exteut a reasonable variation iu the diffusiou rate may alter the iupact of sedimentation."," But, as we consider the actual value of $D$ to be uncertain even in the case of diffusing in a single component plasma, we carry out calculations with $D= D_s$, $5 D_s$, and $10 D_s$ to determine to what extent a reasonable variation in the diffusion rate may alter the impact of sedimentation."
 Although the actual D may © less than Dy. we oulv consider D>Dy. here since for D«Dy. the pact of ou WD evolution is irrelevant.," Although the actual $D$ may be less than $D_s$, we only consider $D>D_s$ here since for $D \ll D_s$, the impact of on WD evolution is irrelevant."
" The variation of up to 10 times D, is probably large for he more typical low mass WDs. but for the more massive WDs. where the interior plasma cau be a quanti liquid over a substantial portion of the WD before it crystallizes. he use of the classical OCP formula may be seriously iu CLYOY."," The variation of up to 10 times $D_s$ is probably large for the more typical low mass WDs, but for the more massive WDs, where the interior plasma can be a quantum liquid over a substantial portion of the WD before it crystallizes, the use of the classical OCP formula may be seriously in error."
" Ouce D is specified. the dowuward drift velocity. 6. of he ious uuder a macroscopic force 21,9 Is where the latter equality holds for D=D.."," Once $D$ is specified, the downward drift velocity, $v$, of the ions under a macroscopic force $- 2 m_p g$ is where the latter equality holds for $D=D_s$."
 This is the approxination we use to cetermune e m our calculations., This is the approximation we use to determine $v$ in our calculations.
 As KTU(r)=Ημ] over most of the stellar radius. the motion of is dominated by falling at the local drift velocity set by equation (6)) (BITOL).," As $k T \ll U(r) \equiv \int_0^r 2 m_p g dr$ over most of the stellar radius, the motion of is dominated by falling at the local drift velocity set by equation \ref{eq:v}) ) (BH01)."
" Equation (6)) also allows a rough estimate of the sedimentation timescale,", Equation \ref{eq:v}) ) also allows a rough estimate of the sedimentation timescale.
 The time it takes a iow to sink from the surface to the center of a constant density star whe it falls at the rate given by equation (6)) is (BITOL)., The time it takes a ion to sink from the surface to the center of a constant density star when it falls at the rate given by equation \ref{eq:v}) ) is (BH01).
 Using the central ασιαος as an estimate. a pure oxvecn(carbon) O.6AL. WD at T.—10 K has f.cm[2(27) Cor. so that complete sedimentation is uulikelv for the most common 0.637. WDs.," Using the central densities as an estimate, a pure oxygen(carbon) $0.6
M_\odot$ WD at $T_c=10^7 \ {\rm K}$ has $t_s\approx 42(27) \ {\rm
Gyr}$ , so that complete sedimentation is unlikely for the most common $0.6M_\odot$ WDs."
 Towever. the strong deusitv dependence moeans that this time decreases for massive WDs.," However, the strong density dependence means that this time decreases for massive WDs."
" For example. à pure oxvecn 1.227. WD at 10K has tyz(αντ, "," For example, a pure oxygen $1.2M_\odot$ WD at $10^7 {\rm K}$ has $t_s\approx 9 \,{\rm Gyr}$."
Oue caveat of our specification of the relation between e and D should be mentioned., One caveat of our specification of the relation between $v$ and $D$ should be mentioned.
 Our WD model is constructed using a fully degenerate EOS evervwhere Qvith no inclusion of jou contributions iu the structure calculation) and we take initially the mass fraction to be =00.02 everwhlere., Our WD model is constructed using a fully degenerate EOS everywhere (with no inclusion of ion contributions in the structure calculation) and we take initially the mass fraction to be 0.02 everywhere.
 This choice of iiodel aud profile confines us to the WD interior where the electron Feri euergev op>ET., This choice of model and profile confines us to the WD interior where the electron Fermi energy $\varepsilon_F > k T$.
 Modeling the outer non-degenerate parts of the star requires a lore accurate background model aud electric field calculation., Modeling the outer non-degenerate parts of the star requires a more accurate background model and electric field calculation.
 Since we are most interested here iu the overall cucreetics. we can safely ucelect the role of the thin outer lavers.," Since we are most interested here in the overall energetics, we can safely neglect the role of the thin outer layers."
 Our WD unodel consists of a plasma of a sinele xedonmiuaut ion species phis a trace amouut of at a local umber deusity of 2»=poo2211., Our WD model consists of a plasma of a single predominant ion species plus a trace amount of at a local number density of $n_{22} = \rho_{22}/22 m_p$.
 The energeties included are the heat provided by sedimentation. ateut leat released upon crystallization of the major species. aud the WD's radiative losses.," The energetics included are the heat provided by sedimentation, latent heat released upon crystallization of the major species, and the WD's radiative losses."
 Since we are nodeling the late (100. My). evolution. we ucelect the slight eravitational cucrev release as the WD coutracts o its final radius.," Since we are modeling the late $ 100$ Myr) evolution, we neglect the slight gravitational energy release as the WD contracts to its final radius."
 Also. other poteutial heat sources such as the potential IT burning in the iuner envelope and fractionation effects are not included for the sake of sinplicitv as our goal here is to demonstrate the poteutial impact of sedimenutation separated from other sources of uncertainty iu WD cooling.," Also, other potential heat sources such as the potential H burning in the inner envelope and fractionation effects are not included for the sake of simplicity as our goal here is to demonstrate the potential impact of sedimentation separated from other sources of uncertainty in WD cooling."
 The power soncrated per jou bw the sedinieutation process is Fevpnip. Where Όρων is the bulk velocity of the flow.," The power generated per ion by the sedimentation process is $\vec{F}\cdot\vec{v}_{bulk}$, where $\vec{v}_{bulk}$ is the bulk velocity of the flow."
 This quantity is different. frou the dift velocity since we must take into account both diffusive aud drifting currents., This quantity is different from the drift velocity since we must take into account both diffusive and drifting currents.
" Thus v4.7,=J»o»/po» aud the net eravitational energv release is given by where go»=poso/22m,.", Thus $\vec{v}_{bulk} = \vec{J}_{22}/\rho_{22}$ and the net gravitational energy release is given by where $n_{22} = \rho_{22}/22 m_p$.
 The lateut leat released upou crystallization is /=W77TALE per nuclei (Saluis 20003., The latent heat released upon crystallization is $l=0.77 k T$ per nuclei \citep{sal00}.
. For the energy lost to radiation. we use the LTe: relations of Althaus&Beuvenuto (1998)..," For the energy lost to radiation, we use the $L-T_C$ relations of \citet{ab98}. ."
 For the heat capacity of the liquid state. Cy-. we use equation (17) of Potekhiu&Chabrier(2000) which takes into account the ideal jou and jon-ion interactions.," For the heat capacity of the liquid state, $C_V$, we use equation (17) of \citet{pot00} which takes into account the ideal ion and ion-ion interactions."
 We neglect the sinall electron contributions to the heat capacity., We neglect the small electron contributions to the heat capacity.
 Quantum liquid. effects are not taken iuto account., Quantum liquid effects are not taken into account.
 In the solid state we utilize equation (5) of Chabiier (1993).., In the solid state we utilize equation (5) of \citet{cha93}. .
 Since the thermal conduction time across the WD core is mich less than the evolution time. we treat all cherev sources and sinks as elobal quantities.," Since the thermal conduction time across the WD core is much less than the evolution time, we treat all energy sources and sinks as global quantities."
" The thermal evolution of the instantaneously isothermal WD at temperatureTe: is then where M, Is the mass of the WD that has crystallized.", The thermal evolution of the instantaneously isothermal WD at temperature$T_C$ is then where $M_{crys}$ is the mass of the WD that has crystallized.
 For a spherically sviumetric flow. equation (2))reads We discretize equation (10)) according το. the," For a spherically symmetric flow, equation \ref{eq:cont}) )reads We discretize equation \ref{eq:SPcont})) according to the"
polarity effect.,polarity effect.
" Other than this, the three ZCs profiles within the seen in 630.25 nm V image most likely umbradue to the Stokeseffect (see Fig."," Other than this, the three ZCs profiles within the umbra seen in 630.25 nm Stokes $V$ image are most likely due to the magneto-optical effect (see Fig."
"are 3c)causedbystrongZeemanef magneto-opticalfect(e.g.,West&Hagyard1983)"," \ref{FIG:ABVSample}$ $c$ ) caused by strong Zeeman effect \citep[e.g.,][]{West+Hagyard1983SoPh...88...51W}."
" In the upper panels of Figure 4,, we show Dopplergrams (Vj, Vze and for each spectral line."," In the upper panels of Figure \ref{FIG:VELO}, , we show Dopplergrams $\nu_i$, $\nu_{zc}$ and $\nu_{lp}$ ) for each spectral line."
" are constructed using the LOS vj»)velocities derived from the Theyshifts of Stokes /, V and LP profilesat each spatial point, respectively."," They are constructed using the LOS velocities derived from the shifts of Stokes $I$, $V$ and $LP$ profilesat each spatial point, respectively."
 We note, We note
will primarily affect sources with low flux.,will primarily affect sources with low flux.
 These sources will have poorly measured. flux ratios. and will suffer from optical selection effects due to their more poorly determined positions.," These sources will have poorly measured flux ratios, and will suffer from optical selection effects due to their more poorly determined positions."
 Other biases include (he locations of optical telescopes. source clustering. ancl other effects.," Other biases include the locations of optical telescopes, source clustering, and other effects."
 It will be important to catalog these ellects explicitlv: in particular. insuring adequate optical coverage may require active preparation and participation.," It will be important to catalog these effects explicitly; in particular, insuring adequate optical coverage may require active preparation and participation."
 GLAST will be able to measure the differences in blazar attenuation in the cosmologically interesting range in redshift που z=1 up io :=5., GLAST will be able to measure the differences in blazar attenuation in the cosmologically interesting range in redshift from $z=1$ up to $z=5$.
 This is in contrast to observations of TeV attenuation by IR. radiation. whieh will only be able to measure differences well below :=1. where the IR becomes opaque.," This is in contrast to ground-based observations of TeV attenuation by IR radiation, which will only be able to measure differences well below $z=1$, where the IR becomes opaque."
 As the enerev threshold of the ground based. experiments drops over “ime. their redshift range will increase. bul will remain limited to low redshifts except for exceptional. statistically insignificant special cases. especially given their generally sinall fields of view.," As the energy threshold of the ground based experiments drops over time, their redshift range will increase, but will remain limited to low redshifts except for exceptional, statistically insignificant special cases, especially given their generally small fields of view."
 More (han establishing that EBL altenuation occurs. GLAST will be able to distinguish between different EBL models.," More than establishing that EBL attenuation occurs, GLAST will be able to distinguish between different EBL models."
 This would validate EBL attenuation as a direct cosmological probe., This would validate EBL attenuation as a direct cosmological probe.
 We emphasize (hat (his analvsis will require redshift determinations of a large Iraction ol GLAST blazars., We emphasize that this analysis will require redshift determinations of a large fraction of GLAST blazars.
 This is another example of the importance of eross-wavelength studies: by using optical measurements of blazar redshifts. eamma-rayv measurements can uniquelv probe the optical-UV EBL.," This is another example of the importance of cross-wavelength studies: by using optical measurements of blazar redshifts, gamma-ray measurements can uniquely probe the optical-UV EBL."
 A redshift measurement for thousands of high-redshift sources is nol a trivial undertaking. but the effort. will be well rewaircec.," A redshift measurement for thousands of high-redshift sources is not a trivial undertaking, but the effort will be well rewarded."
 Even alter observation of a redshift-dependent effect. (he possibility would remain that the spectral evolution of ganuna-ray blazars might coincidentally mimic redshift«dependent EBL absorption.," Even after observation of a redshift-dependent effect, the possibility would remain that the spectral evolution of gamma-ray blazars might coincidentally mimic redshift-dependent EBL absorption."
 For example. if blazars that. formed in the early universe sullered. more internal attenuation than blazars that formed later. the same effect could be produced.," For example, if blazars that formed in the early universe suffered more internal attenuation than blazars that formed later, the same effect could be produced."
 Note that blazars are variable. aud there are some imdications (hat their spectra can become harcler when {μον flare (Sreekumar et al.," Note that blazars are variable, and there are some indications that their spectra can become harder when they flare (Sreekumar et al."
 1996)., 1996).
 Evolution in flaring probability could produce the sanie elfect as actual spectral evolution [rom a statistical standpoint (for example. a higher percentage of high-redshift blazars might be observed in a quiescent phase). although one would expect the GLAST [flux limit to produce a selection effect in the opposite direction.," Evolution in flaring probability could produce the same effect as actual spectral evolution from a statistical standpoint (for example, a higher percentage of high-redshift blazars might be observed in a quiescent phase), although one would expect the GLAST flux limit to produce a selection effect in the opposite direction."
 In anv case. observation of a redshilt-dependent spectral softening will provide au important constraint.," In any case, observation of a redshift-dependent spectral softening will provide an important constraint."
 Theorists will have to decide the likelihood of an evolutionary conspiracy., Theorists will have to decide the likelihood of an evolutionary conspiracy.
 We acknowledge useful conversations with Bill Atwood. who first suggested using the large statisties of GLAST AGNs to look for svstematie effects of extragalactic background light attenuation with redshift.," We acknowledge useful conversations with Bill Atwood, who first suggested using the large statistics of GLAST AGNs to look for systematic effects of extragalactic background light attenuation with redshift."
 We would also like to thank David Thompson. Seth Diegel. Flovd Stecker. and. Mike Salamon lor their useful comments. and James Bullock ancl Joel Primack for providing calculations from their models.," We would also like to thank David Thompson, Seth Digel, Floyd Stecker, and Mike Salamon for their useful comments, and James Bullock and Joel Primack for providing calculations from their models."
 We would like to thank (he referees, We would like to thank the referees
"visual extinction through the CSE is ~100 mag, this scenario requires the existence of channels of greatly reduced extinction through which the ultraviolet radiation can penetrate.","visual extinction through the CSE is $\sim 100$ mag, this scenario requires the existence of channels of greatly reduced extinction through which the ultraviolet radiation can penetrate."
" The consequences of this model for the abundances and spatial distribution of the many other molecules detected in IRC+10216, many af which have been observed interferometrically, has yet tofbe fully investigated."," The consequences of this model for the abundances and spatial distribution of the many other molecules detected in IRC+10216, many of which have been observed interferometrically, has yet to be fully investigated."
" One possible test of this model might involve a search for H3’O or H38O. Because the !?CO photodissociation rate is sharply reduced by self-shielding, D10 have emphasized the importance of 1300 as a source of atomic O that tan reads to fom H2ONThe ο... todissociation c fo C!7OLand C&O wauld be at st afΏιτρο 4s that (@ 139002"," One possible test of this model might involve a search for $_2^{17}$ O or $_2^{18}$ O. Because the $\rm ^{12}CO$ photodissociation rate is sharply reduced by self-shielding, D10 have emphasized the importance of $\rm ^{13}CO$ as a source of atomic O that can react to form $_2$ O. The photodissociation rates for $^{17}$ O and $^{18}$ O would presumably be at least as large as that for $\rm ^{13}CO$."
" Thu$dfofestomic CO, C''O d C80 were the only suppliers oxygen, due HyH38Oμιά and H250WHO/H3""Q'/ ratiog eyould approach (be ne C3 ratios respectively."," Thus, if $\rm ^{13}CO$, $^{17}$ O and $^{18}$ O were the only suppliers of atomic oxygen, the $\rm H_2^{16}O/H_2^{18}O$ and $\rm H_2^{16}O/H_2^{17}O$ ratios would approach the $\rm ^{13}C$ $^{18}$ O and $\rm ^{13}C$ $^{17}$ O ratios respectively."
" (Given theη PC/PC etand 149/1950 -sotopic ratios determined by. Kahane 4(1992) pd by Gemicharo Gai| Kahane 00 for IRC+10216, “πω 3 C/C)ios are respectively =BO smaller than ratios in thejCSE."," Given the $\rm ^{13}C/^{12}C$, $\rm ^{18}O/^{16}O$ and $\rm ^{17}O/^{16}O$ isotopic ratios determined by Kahane et (1992) and by Cernicharo, Guéllin Kahane (2000) for IRC+10216, the $\rm ^{13}CO$ $^{17}$ O and $\rm ^{13}CO$ $^{18}$ O ratios are respectively $\sim$ 18 and 28, each a factor of $45$ (= $\rm ^{12}C/^{13}C$ ) smaller than the elemental $\rm ^{16}O$ $^{17}$ O and $\rm ^{16}O$ $^{18}$ O ratios in the CSE."
 If en a H4°O/H}'O ⊳⊋⊙∕∐≟∂⊙ ratio larger than To e," If SiO or $\rm ^{12}CO$ are significant additional sources of atomic oxygen in the D10/ACG picture, then a $\rm H_2^{16}O$ $_2^{17}$ O ratio larger than 18 – or a $\rm H_2^{16}O$ $_2^{18}$ O ratio larger than 28 – might still be consistent with the model."
t 22010) plade upper 0 (552.05 Hz) and H3°O i, Data obtained in a full HIFI spectral survey carried out toward IRC+10216 (Cernicharo et 2010) place upper limits on the flux of the $_2^{17}$ O (552.021 GHz) and $_2^{18}$ O (547.676 GHz) $1_{10}-1_{01}$ transitions.
" g these 10—101 transition, we determine at the spectral survey places 3 c lower limits on both the 556.936 GHz/552.021 GHz and 556.936 GHz/547.676 GHz line ratios of ~100."," Comparing these with the observed flux in the 556.936 GHz $_2^{16}$ O $1_{10}-1_{01}$ transition, we determined that the spectral survey places 3 $\sigma$ lower limits on both the 556.936 GHz/552.021 GHz and 556.936 GHz/547.676 GHz line ratios of $\sim 100$."
" Taking account of optical depth effects, we find that these correspond to a lower limit of ~200 on both the H160/H1'O and Hi50/H15O0 abundance ratios."," Taking account of optical depth effects, we find that these correspond to a lower limit of $\sim 200$ on both the $\rm H_2^{16}O/H_2^{17}O$ and $\rm H_2^{16}O/H_2^{18}O$ abundance ratios."
" Given the elemental isotopic ratios ~1260 and 160/170~840 (Kahane et 190/15011992), these limits imply that the abundance of the minor isotopologues could only be enhanced (by means of isotope-selective photodissociation of CO, for example) by at most a factor 4 (H1'O) or 6 (H15O)."," Given the elemental isotopic ratios $\rm ^{16}O/^{18}O \sim 1260$ and $\rm ^{16}O/^{17}O \sim 840$ (Kahane et 1992), these limits imply that the abundance of the minor isotopologues could only be enhanced (by means of isotope-selective photodissociation of CO, for example) by at most a factor 4 $_2^{17}$ O) or 6 $_2^{18}$ O)."
 Detailed modeling will be required to determine whether our non-detections of HÀ*O and H3°0 are consistent with the water production mechanism proposed by D10/ACG., Detailed modeling will be required to determine whether our non-detections of $_2^{17}$ O and $_2^{18}$ O are consistent with the water production mechanism proposed by D10/ACG.
" HIFI has been designed and built by a consortium of institutes and university departments from across Europe, Canada and the United States under the leadership of SRON Netherlands Institute for Space Research, Groningen, The Netherlands and with major contributions from Germany, France and the US."," HIFI has been designed and built by a consortium of institutes and university departments from across Europe, Canada and the United States under the leadership of SRON Netherlands Institute for Space Research, Groningen, The Netherlands and with major contributions from Germany, France and the US."
" Consortium members are: Canada: CSA, U. Waterloo; France: CESR, LAB, LERMA, IRAM; Germany: KOSMA, MPIfR, MPS; Ireland, NUI Maynooth; Italy: ASI, IFSI-INAF, Osservatorio Astrofisico di Arcetri- INAF; Netherlands: SRON, TUD; Poland: CAMK, CBK; Spain: Observatorio Astronómmico Nacional (IGN), Centro de Astrobiologá (CSIC-INTA)."," Consortium members are: Canada: CSA, U. Waterloo; France: CESR, LAB, LERMA, IRAM; Germany: KOSMA, MPIfR, MPS; Ireland, NUI Maynooth; Italy: ASI, IFSI-INAF, Osservatorio Astrofisico di Arcetri- INAF; Netherlands: SRON, TUD; Poland: CAMK, CBK; Spain: Observatorio Astronómmico Nacional (IGN), Centro de Astrobiologá (CSIC-INTA)."
" Sweden: Chalmers University of Technology - MC2, RSS GARD; Onsala Space Observatory; Swedish National Space Board, Stockholm University - Stockholm Observatory; Switzerland: ETH Zurich, FHNW; USA: Caltech, JPL, NHSC."," Sweden: Chalmers University of Technology - MC2, RSS GARD; Onsala Space Observatory; Swedish National Space Board, Stockholm University - Stockholm Observatory; Switzerland: ETH Zurich, FHNW; USA: Caltech, JPL, NHSC."
"This work has been supported by the Australian Research Council. the University of Svednev. the Hungarian. οΓΙΑ Grants. WGSLG anc ALBOC 51019. and the ""Lendüllet Young Researchers’ Program of the Hungarian. Acaclemy of Sciences.","This work has been supported by the Australian Research Council, the University of Sydney, the Hungarian OTKA Grants K76816 and MB0C 81013, and the “Lendüllet” Young Researchers' Program of the Hungarian Academy of Sciences."
 [XMs. visits. το Syeney University were supported by the University of Svdney International Visiting Fellowship., KW's visits to Sydney University were supported by the University of Sydney International Visiting Fellowship.
The presence of (dynamically decoupled) sub-compponnents is one of the possible explanations of such complex kinematics.,The presence of (dynamically decoupled) nents is one of the possible explanations of such complex kinematics.
 Complicated velocity fields may arise as well from triaxial systems due to competing contributions of different families of orbits (Statler 1991a.b: Statler et al.," Complicated velocity fields may arise as well from triaxial systems due to competing contributions of different families of orbits (Statler 1991a,b; Statler et al."
 1997. and references therein). like. for instance. the possible presence of a substantial fraction of retrograde orbits (Wozniak Pfenninger 1997: compare their Fig.," 1997, and references therein), like, for instance, the possible presence of a substantial fraction of retrograde orbits (Wozniak Pfenninger 1997; compare their Fig."
 5(c) to some of our RC3)., 5(c) to some of our RCs).
 A concluding remark is in order: all the data presented here show very complicated kinematics. which is only in part accounted for by misclassification of SOs into Es.," A concluding remark is in order: all the data presented here show very complicated kinematics, which is only in part accounted for by misclassification of S0s into Es."
 It seems. instead. conceivable that complex kinematics is à common. and dynamically stable. characteristic of elliptical galaxies.," It seems, instead, conceivable that complex kinematics is a common, and dynamically stable, characteristic of elliptical galaxies."
test iudicates that lxο ions are different at a significance level ofe,test indicates that both distributions are different at a significance level of.
 However. the differe1ος occurs iu the opposite seuse of what we were expecting as group thy shows luc1 niore nietalaxich objects than the eroup ubvk. Aso there seenis to be nore moetd-poor stars amouest the ubvk «warts.," However, the difference occurs in the opposite sense of what we were expecting as group tK shows much more metal-rich objects than the group ubvK. Also there seems to be more metal-poor stars amongst the ubvK dwarfs."
 Therefoxc these groups do iot show any bias derived ron CDVRIparallaxes. althoteh ποιο CXCCSS of ucta-yk ‘stars isapwent in the eroup of Ik dwarfs wih trigoroluctric para]axes.," Therefore, these groups do not show any bias derived from parallaxes, although some excess of metal-rich stars is apparent in the group of K dwarfs with trigonometric parallaxes."
 Ilowever. this restIt is nof 'ouclusive. since the metallicitv disributioi of group I<. is more strongly dependent oi the calihtation by Osen (1981) han eroup ubvkk (the yaction of stars in tlose ΤΟ)» that have (by)5! Or 0.55 axd 0.37. res]ectivelv).," However, this result is not conclusive, since the metallicity distribution of group tK is more strongly dependent on the calibration by Olsen \cite{olsen84}) ) than group ubvK (the fraction of stars in these groups that have $(b-y) > 0.550$ is 0.55 and 0.37, respectively)."
 T is worth to note that tli unetallicity distribution of eroup uIk aerees better with he eroups of C dwar*., It is worth to note that the metallicity distribution of group ubvK agrees better with the groups of G dwarfs.
 The hypothesis that the metaicity αςαπλαtion of eroip ubviy is 10 Sade as those of groups ubvCG and t€ can only be reqected at a sijeuificauce evel of 0.2216 and 0.2339. respecively.," The hypothesis that the metallicity distribution of group ubvK is the same as those of groups ubvG and tG can only be rejected at a significance level of 0.2246 and 0.2339, respectively."
 Itistlen particuarly iuportaut to seo wlrether there are differences amongst t1e eroups of deviating stars. that is. those ob.ects to he right of the dot-dasred lines in Fig. 5..," It is then particularly important to see whether there are differences amongst the groups of deviating stars, that is, those objects to the right of the dot-dashed lines in Fig. \ref{cnship},"
 and he roiuaiuing Ooeroups., and the remaining groups.
 The metallicites of the deviating G stars raiege Yolu 0.5 to BUSex. with an average around 0.10 dex.," The metallicites of the deviating G stars range from $-0.5$ to $+0.1$ dex, with an average around $-0.10$ dex."
 There is no indication that this eroup has more metal-ricT stars compared to the others., There is no indication that this group has more metal-rich stars compared to the others.
 The absence of stars wih metallicites lower than 0.5 dex can well be ascribed to the size of the sample., The absence of stars with metallicites lower than $-0.5$ dex can well be ascribed to the size of the sample.
" As au illustration. the KS est eives significance levels of 0.517 aud 0,511 for this distribution not to be taken from the sale population of eroups ubvG aud tC. respectively,"," As an illustration, the KS test gives significance levels of 0.517 and 0.314 for this distribution not to be taken from the same population of groups ubvG and tG, respectively."
 The situation is cdiffereut or Ix dwarts., The situation is different for K dwarfs.
 The group of deviating I dwarfs has ietalcities ranging fro l.6 to 0.05 dex. with alb average around 0.65 dex.," The group of deviating K dwarfs has metallicities ranging from $-1.6$ to $-0.05$ dex, with an average around $-0.65$ dex."
 This result is very peculiar since 1 sugecsts that the stars which have Svsteniaically uncderestinated disances nu he CNSS3 are metal-poor. while we woud expec that mecal-poor stars would have overestimated distances according to Figure T.," This result is very peculiar since it suggests that the stars which have systematically underestimated distances in the CNS3 are metal-poor, while we would expect that metal-poor stars would have overestimated distances according to Figure \ref{25pclim}."
 However. this question cautot be properly answered because the metallicity of he group of deviaine I& darts also strougly depends ou the imeallicity calibration for stars cooler than (bi)0.550.," However, this question cannot be properly answered because the metallicity of the group of deviating K dwarfs also strongly depends on the metallicity calibration for stars cooler than $(b-y)=0.550$."
 Tn spite of that. if such jas is likely to be esent 1ithe catalogue. it should occur or both €t aud I& dwarts. |clue in fact stronger for the loter stars.," In spite of that, if such bias is likely to be present in the catalogue, it should occur for both G and K dwarfs, being in fact stronger for the hotter stars."
 The nou-existeuce of such bias amongst the € dwarts. which inve even Wore accurate photometric metalicites. indicates that it docs rot affect the conteut of he CNS3.," The non-existence of such bias amongst the G dwarfs, which have even more accurate photometric metallicites, indicates that it does not affect the content of the CNS3."
 This cau happen jecauxe the limit for inclusion o: objects 1i the CNS3 due o spectroscopic andVRZ parallaxes is more flexible han the limit for trizonome11 paralaxes., This can happen because the limit for inclusion of objects in the CNS3 due to spectroscopic and parallaxes is more flexible than the limit for trigonometric parallaxes.
 This is evideut roni Fieure &.., This is evident from Figure \ref{cnship}.
 Iu this plot we see that here are lualiv stars dm the CNS3 whose distauces in this catalogue are ercater than 25 pc. alongs those 1icluded withΕΡΤ xrallaxes.," In this plot we see that there are many stars in the CNS3 whose distances in this catalogue are greater than 25 pc, amongst those included with parallaxes."
 Thus. in the CNS3 there is not a fixed limit at 25 pe for the inclusion of stars with parallaxes. and there seeuis to be uo correspolnine moetallicitv-bias.," Thus, in the CNS3 there is not a fixed limit at 25 pc for the inclusion of stars with parallaxes, and there seems to be no corresponding metallicity-bias."
 The siuplest hvpothesis to accoui for the results found by Favata ct al. (1997)), The simplest hypothesis to account for the results found by Favata et al. \cite{favata1}) )
 is that their sample is not representaive of the ooOe:dactic populatic1i of I& dwarfs. due to its sia] size.," is that their sample is not representative of the galactic population of K dwarfs, due to its small size."
" The originalOo sample randomly selected from the €'""NS2. and consisting of arouxd 100 Ce and 100 I dwarfs (Favata ct al. 1996))."," The original sample randomly selected from the CNS2, and consisting of around 100 G and 100 K dwarfs (Favata et al. \cite{favata2}) ),"
 can ο expected. to be representaive., can be expected to be representative.
 However. the muaver of stars that were effectively. observed is 63 Co dwarfs ik 26 I& dwiurfs.," However, the number of stars that were effectively observed is 63 G dwarfs and 26 K dwarfs."
 As Favata et al. (1997)), As Favata et al. \cite{favata1}) )
 themselves state. relatively fewer cooler stars were observed due to hei faint maguitudes.," themselves state, relatively fewer cooler stars were observed due to their faint magnitudes."
 This observational selection is rot likely to remove the roepresentativeness of a large data sample., This observational selection is not likely to remove the representativeness of a large data sample.
 Towever. simall samples are mich casily affected by statistical fucuations due to the elimination of some sstars.," However, small samples are much easily affected by statistical fluctuations due to the elimination of some stars."
 This can explain why the distributions by Favata ot al. (1997)), This can explain why the distributions by Favata et al. \cite{favata1}) )
 show large fluctuations and not a single prouineut peak as f1ο other metallicity distributions iu the licrature., show large fluctuations and not a single prominent peak as the other metallicity distributions in the literature.
lu inost cases the solution for E is executed Mice: first. a coarse fit at order No=2 is attenpted.,"In most cases the solution for ${\bf E}$ is executed twice: first, a coarse fit at order $N=2$ is attempted."
 If this fails. the flag is set.," If this fails, the flag is set."
 [fit succeeds. the coarse solution is used as 16 starting point for a fit to the full requested IN.," If it succeeds, the coarse solution is used as the starting point for a fit to the full requested $N$."
 Another subtlety is that the pixel mask is not Cladiced between iterationstlie mask maintaius 1e shape aud size specified by the initial guess (or 1C result of the coarse fit)., Another subtlety is that the pixel mask is not changed between iterations—the mask maintains the shape and size specified by the initial guess (or the result of the coarse fit).
 This should not bias 1C fit toward the initial guess. however. because We 0we fitting to the data within the mask rather lal1 executing a stun of monients over pixels within the mask.," This should not bias the fit toward the initial guess, however, because we are fitting to the data within the mask rather than executing a sum of moments over pixels within the mask."
 If the size jp ofthe object changes οσαnatically diving the coarse fit due to a very poor iuitial estimate from SExtractor). hen we (lo resize the masks aud start over.," If the size $\mu$ of the object changes dramatically during the coarse fit due to a very poor initial estimate from ), then we do resize the masks and start over."
 If at any tic the mask ou an image grows to cucompass too lareeOo a iuuberof pixels. we set the flagOo aud quit.," If at any time the mask on an image grows to encompass too large a number of pixels, we set the flag and quit."
 This preveuts the program from ecttiug hung up calculating pixel sus for nouseusicallv laree objects., This prevents the program from getting hung up calculating pixel sums for nonsensically large objects.
1n addition to spectroscopic observations already reported in he literature and. described in the Appendix. new spectra or some of the starburst radio galaxies were obtained in various runs on the ESO: Very. Large ‘Telescope (VET). 3.6m and New Technology Telescope (NEL) telescopes.,"In addition to spectroscopic observations already reported in the literature and described in the Appendix, new spectra for some of the starburst radio galaxies were obtained in various runs on the ESO Very Large Telescope (VLT), 3.6m and New Technology Telescope (NTT) telescopes."
 Details of hese new observations are given in Table 3., Details of these new observations are given in Table 3.
 The reduction of these data followed the standard steps of bias subtraction. lat fielding. wavelength calibration. atmospheric extinction correction. Hux calibration. and registration of blue and red spectral images.," The reduction of these data followed the standard steps of bias subtraction, flat fielding, wavelength calibration, atmospheric extinction correction, flux calibration, and registration of blue and red spectral images."
" Dased on measurements of night sky lines. he uncertaintv in the wavelength. calibration is estimated o be smaller than for all datasets: the relative Dux. calibration uncertainty. sed on the comparison of observations of several spectroscopic standard: stars. is better than Επι,"," Based on measurements of night sky lines, the uncertainty in the wavelength calibration is estimated to be smaller than for all datasets; the relative flux calibration uncertainty, based on the comparison of observations of several spectroscopic standard stars, is better than $\pm$."
 Prior o the analvsis of the data. the spectra were corrected. for Galactic extinction using values of E(GD-V) derived. from Schlegel.Finkbeiner&Davis(1998). along with the Seaton(1979) extinction. law.," Prior to the analysis of the data, the spectra were corrected for Galactic extinction using values of E(B-V) derived from \citet{schlegel98} along with the \cite{seaton79} extinction law."
 Note that. in the Cases of the VLE/EORRSI observations of. PINSO023-26 and. PINS1549-79. the data were taken in spectropolarimetric mode with he light.5 passing5 through5 a half-wave plate as well as an analvsing prism.," Note that, in the cases of the VLT/FORS1 observations of PKS0023-26 and PKS1549-79, the data were taken in spectropolarimetric mode with the light passing through a half-wave plate as well as an analysing prism."
 In the latter two cases. the ο and ο rav images of the long-slit spectra were spatially registered and. co-adcded. and the Dux calibration was performed. using spectropolarimetric standard: stars observed. through the same instrument configuration as the target galaxies.," In the latter two cases, the `o' and `e' ray images of the long-slit spectra were spatially registered and co-added, and the flux calibration was performed using spectropolarimetric standard stars observed through the same instrument configuration as the target galaxies."
 For some of the objects discussed in this paper (Fornax A. PINSO652-20. PINS1549-79) we have carried. out. new spectral svnthesis modelling in an attempt to improve our knowledge of the properties of the voung stellar populations (YSD).," For some of the objects discussed in this paper (Fornax A, PKS0652-20, PKS1549-79) we have carried out new spectral synthesis modelling in an attempt to improve our knowledge of the properties of the young stellar populations (YSP)."
" In these cases we followed t1ο procedures outlined in Tadhunteretal.(2005).. Holtetal. (2007).. and. Willsetal. (2008): modelling the continuum spectrum using a combination of a reddened. YSP (0.005<£44<5 Gyr. 0<E(B.V)« 2.0) and an unreddened OSP (1,4,=12.5 Gyr). taking full account of AGN continuum components (where appropriate)."," In these cases we followed the procedures outlined in \citet{tadhunter05}, \citet{holt07}, and \citet{wills08}: modelling the continuum spectrum using a combination of a reddened YSP $0.005 < t_{ysp} < 5$ Gyr, $0 < E(B-V) < 2.0$ ) and an unreddened OSP $t_{ysp}=12.5$ Gyr), taking full account of AGN continuum components (where appropriate)."
 In all cases we used instantaneous burst. solar metallicity spectral synthesis templates from Muzual&Charlot(2005).," In all cases we used instantaneous burst, solar metallicity spectral synthesis templates from \citet{bruzual03}."
.. Detailed CLISCUSSLODS of the individual starburst racio ealaxies. inclucling the new spectroscopic results./ are presented in the Appendix.," Detailed discussions of the individual starburst radio galaxies, including the new spectroscopic results, are presented in the Appendix."
 In this section we consider the eeneral results and. patterns obtained for the population of starburst radio galaxies as a whole., In this section we consider the general results and patterns obtained for the population of starburst radio galaxies as a whole.
Il line in several CASLEO spectra ancl the fact that this line can be contaminated by the Ile line. the S index is not suitable to study (he chromospheric activity on the faint stars in (his study.,"H line in several CASLEO spectra and the fact that this line can be contaminated by the $\epsilon$ line, the $S$ index is not suitable to study the chromospheric activity on the faint stars in this study."
 In this paper we use as a proxy of stellar activity the Cau Ix line-core flix. integrated with a triangular prolile of 1.09 EWIIM to mimic the response of the Mount. Wilson instrument.," In this paper we use as a proxy of stellar activity the Ca K line-core flux, integrated with a triangular profile of 1.09 FWHM to mimic the response of the Mount Wilson instrument."
 We complement our data with photometry from the All Sky Automated Survey (ASAS)., We complement our data with photometry from the All Sky Automated Survey (ASAS).
 This project is dedicated to constant photometric monitoring of the whole available sky. including approximately 10* stars brighter than 14 magnitude 2002)..," This project is dedicated to constant photometric monitoring of the whole available sky, including approximately $^7$ stars brighter than 14 magnitude \citep{2002AcA....52..397P}."
 Presently. ASAS consists of two observing stations. one in Las Campanas Observatory. Chile (since 1997) and the other on Haleakala. Maii (since 2006).," Presently, ASAS consists of two observing stations, one in Las Campanas Observatory, Chile (since 1997) and the other on Haleakala, Maui (since 2006)."
 Each site is equipped with (svo wide-field 200/2.8 instruments. observing simultaneously in the V and Ibands?.. In this work we used the mean V magnitude of each observing season {ο analvze the long-term activity of the stars.," Each site is equipped with two wide-field 200/2.8 instruments, observing simultaneously in the V and I. In this work we used the mean V magnitude of each observing season to analyze the long-term activity of the stars."
 C] 229 A is a moderately active dMI fIare star (Byrneetal. 1985)... with a brown dwarl companion (Nakajimaetal.1995)..," Gl 229 A is a moderately active dM1 flare star \citep{1985MNRAS.214..119B}, , with a brown dwarf companion \citep{1995Natur.378..463N}."
 We observed GI 229 A for 20 nights between the vears 2002 and 2010., We observed Gl 229 A for 20 nights between the years 2002 and 2010.
 In Fig., In Fig.
 1 we plot each spectrum. with a label which indicates ihe month ancl vear of the observation (e.g. 0302 means that the observation was obtained in March. 2002).," \ref{fig.sp_hd42581} we plot each spectrum, with a label which indicates the month and year of the observation (e.g. 0302 means that the observation was obtained in March, 2002)."
 We excluded (he spectrum obtained in November 2004 due to ils low signal-to-noise near the Cau lines and the one observed in March 2010 because it was obtained in acloudy. night., We excluded the spectrum obtained in November 2004 due to its low signal-to-noise near the Ca lines and the one observed in March 2010 because it was obtained in acloudy night.
The data from low-frequeney Galactic surveys are uxuallv obtained by scanning the sky with transit instruments.,The data from low-frequency Galactic surveys are usually obtained by scanning the sky with transit instruments.
 By “scanning effects”. we refer to the stripes xoduced. in the temperature maps (aud particularly roticeable im the spectral index map) because of the different observing conditions. e.g. differences iu eal or krospheric absorption. between adjacent scans lade wo the instrmucut.," By “scanning effects”, we refer to the stripes produced in the temperature maps (and particularly noticeable in the spectral index map) because of the different observing conditions, e.g. differences in gain or ionospheric absorption, between adjacent scans made by the instrument."
 Another spurious effect is produced when surveys made at differeut epochs or zoues or with different instruments are combined in a process described. or example. iu ? where the los MIIZ all-sky: πα was obtained.," Another spurious effect is produced when surveys made at different epochs or zones or with different instruments are combined in a process described, for example, in \citet{Haslam1981} where the 408 MHz all-sky map was obtained."
 This effect is caused by the matching of partial uaps that together form the final map. and we refer o this as the “matching effect”.," This effect is caused by the matching of partial maps that together form the final map, and we refer to this as the “matching effect”."
 All these spurious catures and residues are usually produced along coustaut declination and/or right The [5 ΑΠΕ map has a matching effect. produced. by combining the northern aud southern surveys at àzz107.paper., All these spurious features and residues are usually produced along constant declination and/or right The 45 MHz map has a matching effect produced by combining the northern and southern surveys at $\delta\approx10^\circ$.
 Scanning effects iu this survey are preseut especially at thepositious RA: 15hzRAz21h aud 30°XDECz 307., Scanning effects in this survey are present especially at thepositions RA: $15h\lesssim RA \lesssim 24h$ and $-30^\circ\lesssim DEC\lesssim 30^\circ$ .
 They appear, They appear
present the model used to derive the coluun deusities and temperatures for the various species m section 3 and describe the results for each of the molecules iu section [..,present the model used to derive the column densities and temperatures for the various species in section \ref{sec:fit} and describe the results for each of the molecules in section \ref{sec:results}.
 Subsequenuthe im section 5.. we discuss the possible scenario iu which we are probing chemistry in a circtuustellay disk.," Subsequently, in section \ref{sec:mod}, we discuss the possible scenario in which we are probing chemistry in a circumstellar disk."
 In section 6.. we provide some conchiding remarks.," In section \ref{sec:conc}, we provide some concluding remarks."
 The observations were made with TENES. the Texas Echelon Cross Echelle Spectrograph (Lacy.ctal.2002).. on the NASA DIufrared Telescope Facility (RTF) in 2001 June and November. 2002 September auc December. and 2005 December.," The observations were made with TEXES, the Texas Echelon Cross Echelle Spectrograph \citep{lacy02}, on the NASA Infrared Telescope Facility (IRTF) in 2001 June and November, 2002 September and December, and 2005 December."
 TENES is a high-resolution dispersed spectrograph operating at wwavelengths between 5 and 25jan., TEXES is a high-resolution cross-dispersed spectrograph operating at wavelengths between 5 and 25.
 It achieves a spectral resolution of R = 75.000-100.000 |km ij) shortward of 11 wwith a slit width of ~1.5”« (," It achieves a spectral resolution of $R$ = 75,000-100,000 $\Delta v =
3-4$ ) shortward of 14 with a slit width of $\sim$. ("
Of the observations presented here. only the spectra near llpiu. which used a sslit. achieved the ligher resolution.),"Of the observations presented here, only the spectra near 11, which used a slit, achieved the higher resolution.)"
" The spatial resolution aloug the ὃςxth-South orieuted slit is —1"". slightly larger than the «iffvaction lint of the IRTF."," The spatial resolution along the North-South oriented slit is $\sim$, slightly larger than the diffraction limit of the IRTF."
" The spectral aud spatial sauliue by the 256« pixel Si:As detector array are L0 ad.0.35"". respectively,"," The spectral and spatial sampling by the $256 \times 256$ pixel Si:As detector array are 1.0 and, respectively."
" At cach spectral setting. 5-10 orders of the high resolution echelou erating are recorded. eiving a spectral coverage of AA4,4,/A.~0.5..."," At each spectral setting, 5-10 orders of the high resolution echelon grating are recorded, giving a spectral coverage of $\Delta\lambda_{tot}/\lambda~\sim~$."
 At wavelengths shortward of 11 ffull order widths are observed. eivinge continuous spectral coverage. but longward of 11 tthe order width exceeds the array width. leaving gaps in the observed spectra.," At wavelengths shortward of 11 full order widths are observed, giving continuous spectral coverage, but longward of 11 the order width exceeds the array width, leaving gaps in the observed spectra."
 The location of the gaps cau be shifted in order to optimize line coverage., The location of the gaps can be shifted in order to optimize line coverage.
 However. since the gaps are smal aud. at 13pan.. there are many telluric lines. it would be impractical to repeat observations with sinall shifts in order to get complete spectral coverage for cach setting.," However, since the gaps are small and, at 13, there are many telluric lines, it would be impractical to repeat observations with small shifts in order to get complete spectral coverage for each setting."
 Mich of the data were taken in high water vapor conditions. but the observing couditious did not seriously affect the spectra except for increased noise. especially on atimospheric lues;," Much of the data were taken in high water vapor conditions, but the observing conditions did not seriously affect the spectra except for increased noise, especially on atmospheric lines."
 Observations of wwere interspersed with observations of a bright asteroid (usually Ceres (2500 Jv) but also IEvecia (150 Jw). which served as a comparison. source for removal of tellure absorption features.," Observations of were interspersed with observations of a bright asteroid (usually Ceres $> 500$ Jy), but also Hygeia $\sim 150$ Jy), which served as a comparison source for removal of telluric absorption features."
 The telescope was nodded by iuoviug the source along the spectrograph cutrance slit to allow subtraction of background cimission.," The telescope was nodded by, moving the source along the spectrograph entrance slit to allow subtraction of background emission."
 At the beeinnine of each set of nodded observations an ambient temperature blackbody was observed for flat-fieldine., At the beginning of each set of nodded observations an ambient temperature blackbody was observed for flat-fielding.
 The first nods were used to peak up on the source by maximizing the throughput of the shit., The first nods were used to peak up on the source by maximizing the throughput of the slit.
 The nod pairs taken during peak up were not iucluded in the final sui., The nod pairs taken during peak up were not included in the final sum.
 Data veduction followed the standard TENES procedure (Lacyetal.2002) of first subtracting the two nod position spectra. flat-fielding with the blackbody-skv difference echelloeram (Lacyctal. 1989).. and interpolating over spikes and bad pixels.," Data reduction followed the standard TEXES procedure \citep{lacy02}
 of first subtracting the two nod position spectra, flat-fielding with the blackbody-sky difference echellogram \citep{lacy89irsh}, , and interpolating over spikes and bad pixels."
 Optimally-weighted poiut-source spectra were then extracted from the echellograms., Optimally-weighted point-source spectra were then extracted from the echellograms.
 The spectra were linearized in Πατ) based on the known optical distortions. and absolute wavenuuniber calibration was obtained from telluric Cluission features and corrected for the Earth's motion relative to LSR.," The spectra were linearized in $(\nu)$ based on the known optical distortions, and absolute wavenumber calibration was obtained from telluric emission features and corrected for the Earth's motion relative to LSR."
 The resulting waveuuniber scale is correct to within +., The resulting wavenumber scale is correct to within 1.
" The same procedure was used or the asteroid spectra. aud then the sspectra were divided by the asteroid spectra with correction for differences in airmass,"," The same procedure was used for the asteroid spectra, and then the spectra were divided by the asteroid spectra with correction for differences in airmass."
 With this procedure. elluric absorption lines as deep as can be removed. although of course the noise increases on lines.," With this procedure, telluric absorption lines as deep as can be removed, although of course the noise increases on lines."
 However. woad response variations often remained. which were removed by fitting the coutimmim in cach echelou order 6 a quartic polynomial aud dividing.," However, broad response variations often remained, which were removed by fitting the continuum in each echelon order to a quartic polynomial and dividing."
 This procedure also sets the continui in each order to one., This procedure also sets the continuum in each order to one.
 Ten spectral settings were observed in the S20 ((13.712.2 pan) region. two in the 860930 ((11.510.7 ) region (soe Figure 1)). aud three in the 12101312 (8.07.6 fan) region (sec Figure 2)).," Ten spectral settings were observed in the 728--820 (13.7–12.2 ) region, two in the 860–930 (11.5–10.7 ) region (see Figure \ref{fig:twelveum}) ), and three in the 1240--1312 (8.0–7.6 ) region (see Figure \ref{fig:eightum}) )."
 A total of 110 echelou orders. or 70 +. were observed.," A total of 110 echelon orders, or $\sim$ 70, were observed."
 The original observations were meant to study the line profiles of aan TICN. which hac previously Όσοι. detected.," The original observations were meant to study the line profiles of and HCN, which had previously been detected."
 Iu addition to the promiqnt fundamental baud lines of aud IICN. many thot band lines were deected in our spectra iucbludiug ines from the aud the bands.," In addition to the prominent fundamental band lines of and HCN, many hot band lines were detected in our spectra including lines from the – and the – bands."
 After the hot baids were identified. uudeutibed ines still remained. some of which were double dipped.," After the hot bands were identified, unidentified lines still remained, some of which were double dipped."
 The double lines turned out to be lines which are split chο to spin-rotatiou intcractious (see Appendix)., The double lines turned out to be lines which are split due to spin-rotation interactions (see Appendix).
 This is the first detection of toward deuse gas and tιο first eround-based detection ofCIT;., This is the first detection of toward dense gas and the first ground-based detection of.
. Previously. ibad ouly been observed with toward the Galactic ceuter (Feuchteruberetal. 2000).," Previously, had only been observed with toward the Galactic center \citep{feucht00}."
. Iu efforts to detect ethane.Collg.. at SIO-8205. uanv lines were detected. though noue corresponded toCog.," In efforts to detect ethane, at 810-820, many lines were detected, though none corresponded to."
.. Some were due toNIT;.. but may lines remained nuidentified (see Figures 3. and 1))," Some were due to, but many lines remained unidentified (see Figures \ref{fig:hncoall} and \ref{fig:hncoblowup}) )."
 Based on their separation we wore able to ideutifv the lues as due to IINCO., Based on their separation we were able to identify the lines as due to HNCO.
 Iuterstellar TINCO is a well known hot core molecule at radio waveleneths (Zincheukoetal.2000).. but it had never been scen at infrared wavelengths.," Interstellar HNCO is a well known hot core molecule at radio wavelengths \citep{zinchenko00}, but it had never been seen at infrared wavelengths."
 Iu total. five molecules were observed in the 728820 region: (Gucludiug Πο). ICN. IINCO.CIL;. aud ΣΠ.," In total, five molecules were observed in the 728–820 region: (including ), HCN, HNCO, and ."
".. At higher frequencies.CIT,... aaud CS were observed in the 12101320 ireeion."," At higher frequencies, and CS were observed in the 1240–1320 region."
 oobservatious frou the grouud are possible for this source because of the large Doppler shift with respect to the telluric lines., observations from the ground are possible for this source because of the large Doppler shift with respect to the telluric lines.
 Table1 gives a sununary of the detected lines, Table\ref{sumlines} gives a summary of the detected lines.
Atleast 5 lines were observed for cach molecule. with some molecules like (Guclicding isotopic lines aud hot bands). aud,"Atleast 5 lines were observed for each molecule, with some molecules like (including isotopic lines and hot bands), and"
al. 1992:,al. \cite{kos92};
 Antia Basu 1991b:: Basu Autia 1995)) aud the temperature can then be determined usine the inverted sound speed., Antia Basu \cite{ab94b}; Basu Antia \cite{ba95}) ) and the temperature can then be determined using the inverted sound speed.
" Πωπονα, iu this work we have restricted ourselves to the radiative interior oulv."," However, in this work we have restricted ourselves to the radiative interior only."
 The estimated value of X at the base of the convection gone can be compared with the iudepeudoenutlv estimated values iuside the couvection zoue., The estimated value of $X$ at the base of the convection zone can be compared with the independently estimated values inside the convection zone.
 Once the Z7 aud X profiles inside the Sun are known it is straightforward to estimate the neutrino fluxes for the various solar neutriuo experiucuts., Once the $T$ and $X$ profiles inside the Sun are known it is straightforward to estimate the neutrino fluxes for the various solar neutrino experiments.
 The represcutation (6) takes care of both the nodifications in the opacity aud unclear euergey generation rate., The representation (6) takes care of both the modifications in the opacity and nuclear energy generation rate.
 Thus. in the linut of a=0 we essentially recover he approach of ST96. for which the required opacity nodification is negligible.," Thus, in the limit of $\alpha=0$ we essentially recover the approach of ST96, for which the required opacity modification is negligible."
 When we try the least squares solution with a=0. the resulting opacity modifications urn out to be ecuerally very small aud merely represcut numerical errors.," When we try the least squares solution with $\alpha=0$, the resulting opacity modifications turn out to be generally very small and merely represent numerical errors."
 We have. in fact. verified that the two solutions are almost identical.," We have, in fact, verified that the two solutions are almost identical."
 Most of the results in this oper have been obtained using this prescription., Most of the results in this paper have been obtained using this prescription.
 In our ornulation bv increasing the value of à we can also fud solutions which satisfv the huuinositv constraint at he expense of allowing for some opacity variatious., In our formulation by increasing the value of $\alpha$ we can also find solutions which satisfy the luminosity constraint at the expense of allowing for some opacity variations.
 This approach would be similar to that of IKosovichev (1996)) who has estimated Z variations. which is esscutially equivalent to fiudiug the opacity modifications.," This approach would be similar to that of Kosovichev \cite{kos96}) ) who has estimated $Z$ variations, which is essentially equivalent to finding the opacity modifications."
 We thus have the choice of cither moditving the nuclear energy eeneration rate to match the solar huninositv (a= 0) or to uodifv the opacitics keeping the nuclear energy eeneration rate fixed (a>> 1) to achieve the sue purpose., We thus have the choice of either modifying the nuclear energy generation rate to match the solar luminosity $\alpha=0$ ) or to modify the opacities keeping the nuclear energy generation rate fixed $\alpha>>1$ ) to achieve the same purpose.
" Naturally, by using intermediate values of à we can obtain solutions which require modifications of both opacitics and nuclear energev generation rates by varviug amount."," Naturally, by using intermediate values of $\alpha$ we can obtain solutions which require modifications of both opacities and nuclear energy generation rates by varying amount."
 Clearly the resulting solution is not unique. but all possible solutions may not be acceptable since some of them may require unacceptably large iiodificatious in opacity or nuclear euergy eeueration rates;," Clearly the resulting solution is not unique, but all possible solutions may not be acceptable since some of them may require unacceptably large modifications in opacity or nuclear energy generation rates."
 Of course. if we can ect the observed solar huninositv for a solution with a—0. then effectively no modification would be required in mput nmicroplivsics.," Of course, if we can get the observed solar luminosity for a solution with $\alpha=0$, then effectively no modification would be required in input microphysics."
 In fact. for some clioices of nuclear reaction rates and opacities we do fud such solutious where no significant imodificatious in inucroplhivesies are required.," In fact, for some choices of nuclear reaction rates and opacities we do find such solutions where no significant modifications in microphysics are required."
 It should be stressed that our technique for inferring the T and X xofiles in the solu interior is absolute. in the seuse that no reference model is required aud the actual profiles of T and [X are determined directly from the sound speed and density profiles.," It should be stressed that our technique for inferring the $T$ and $X$ profiles in the solar interior is absolute, in the sense that no reference model is required and the actual profiles of $T$ and $X$ are determined directly from the sound speed and density profiles."
 Of course. tlhe density and sound speed are determuned bv using a differeutial technique where the differences are linearized about a reference model.," Of course, the density and sound speed are determined by using a differential technique where the differences are linearized about a reference model."
 We use data sets of panode frequencies frou the GONG data (Thll et al. 1996)), We use data sets of $p$ -mode frequencies from the GONG data (Hill et al. \cite{hil96}) )
 aud from the Bie Bear Solar Observatory (BBSO) data (Libbrecht et al. 1990)), and from the Big Bear Solar Observatory (BBSO) data (Libbrecht et al. \cite{lib90}) )
 along with the low degree modes from DiSON data (Elsworth ot al. 1991)), along with the low degree modes from BiSON data (Elsworth et al. \cite{els94}) )
 to iufer the sound speed and density profiles using the Reeularized Least Squares technique (Antia 1996))., to infer the sound speed and density profiles using the Regularized Least Squares technique (Antia \cite{a96}) ).
 The results obtained using different input frequencies are shown in Fie., The results obtained using different input frequencies are shown in Fig.
 1 which displays the relative difference betsyeen inverted profiles aud those iu the model S of CC al. (1996))., \ref{obsinv} which displays the relative difference between inverted profiles and those in the model S of et al. \cite{jcd96}) ).
 It cau be seen that the results obtained using different sets of frequencies agree with cach other to within the estimated errors., It can be seen that the results obtained using different sets of frequencies agree with each other to within the estimated errors.
 The most siguificaut difference in the sound speed between the model aud the Sui. occurs just below the base of the convection zoue aud is very likely ou account of the NX profile in the Sun beiug smoother than that iu the model (Basu Antia 1991: Cough et al. 1996::," The most significant difference in the sound speed between the model and the Sun, occurs just below the base of the convection zone and is very likely on account of the $X$ profile in the Sun being smoother than that in the model (Basu Antia \cite{ba94}; Gough et al. \cite{dog96};"
 Basu 1997))., Basu \cite{b97}) ).
 Apart from this. another noteworthy smaller hump occurs around 7=0.2/2. which is opposite in sigu to that below the convection zone.," Apart from this, another noteworthy smaller hump occurs around $r=0.2R_\odot$ which is opposite in sign to that below the convection zone."
 It is likely that iu lis region the .X profile in the Sun is steeper than that in the model., It is likely that in this region the $X$ profile in the Sun is steeper than that in the model.
 The third minor bump around r=(0.05R. is not very significant and occurs iu a region where the primary inversions are rot likely to be in aux case very reliable., The third minor hump around $r=0.05R_\odot$ is not very significant and occurs in a region where the primary inversions are not likely to be in any case very reliable.
 Simularly. the most siguificaut difference in the density profile occurs inside the convection zouc and is probably due to sinall errors iu opacity. equation of state. surface abundances and/or the depth of the convection zone (Basu Autia 1997)).," Similarly, the most significant difference in the density profile occurs inside the convection zone and is probably due to small errors in opacity, equation of state, surface abundances and/or the depth of the convection zone (Basu Antia \cite{ba97}) )."
 Receutly. it iis been sugeestedee hat the standard value of solar radius (Allen 19733) needs to be reduced (Antia 1997: Schou et al. 1997))," Recently, it has been suggested that the standard value of solar radius (Allen \cite{all73}) ) needs to be reduced (Antia \cite{a97}; Schou et al. \cite{sch97}) )"
 aud 1 would be interesting to estimate the effect of error in the solar radius ou helioseisnuüc inversions., and it would be interesting to estimate the effect of error in the solar radius on helioseismic inversions.
 We have shown in Fie., We have shown in Fig.
 1 the results obtained using the «αποτι value of 695.99 Nha for the solar radius., \ref{obsinv} the results obtained using the standard value of 695.99 Mm for the solar radius.
 We have also performed mversious with a reduced radius of 695.78 Mii aud Fig., We have also performed inversions with a reduced radius of 695.78 Mm and Fig.
 2 compares the results obtained with two different values for solar radius using the same set of observed. frequencies from CONG mouths [10 cata., \ref{invrad} compares the results obtained with two different values for solar radius using the same set of observed frequencies from GONG months 4–10 data.
 It is clear that the small error in solar radius affects the iuversion results to an extent which is much larger than the estimated errors due to those in frequeucies., It is clear that the small error in solar radius affects the inversion results to an extent which is much larger than the estimated errors due to those in frequencies.
 Applviug the procedure outlined iu Section 2 to the inverted profiles for sound speed and density shown iu Fie. 1..," Applying the procedure outlined in Section 2 to the inverted profiles for sound speed and density shown in Fig. \ref{obsinv},"
 we obtain the Z aud NX profiles aud the results are shown in Fie. 3.., we obtain the $T$ and $X$ profiles and the results are shown in Fig. \ref{invtx}.
 All these results have heen obtaiuca with a=0 in equation (6)) and using the pp reactio- rate from Dalicall (1989))., All these results have been obtained with $\alpha=0$ in equation \ref{chisq}) ) and using the pp reaction rate from Bahcall \cite{bah89}) ).
 A Z profile including diffusion (RICII) of heavy elements with surface value of Z=0.015 was assumed for these inversions., A $Z$ profile including diffusion (RICH) of heavy elements with surface value of $Z=0.018$ was assumed for these inversions.
 Ouce again it is clear that he results obtained usine differeut input frequencies are close to one another., Once again it is clear that the results obtained using different input frequencies are close to one another.
" The computed hunuinositv iu these seiuuc models turus out to be between 0.993 ,.. which is roughly cousisteut with the actual solar Iumuinositv."," The computed luminosity in these seismic models turns out to be between $L_\odot$ , which is roughly consistent with the actual solar luminosity."
 The errors iu secondary iuversions arisiug from estimated uncertaiuties iun the input frequencies can be caleulated with a Monte-Carlo simulation., The errors in secondary inversions arising from estimated uncertainties in the input frequencies can be calculated with a Monte-Carlo simulation.
 For, For
the chemical evolution of the interstellar medium.,the chemical evolution of the interstellar medium.
" As time evolves SNe contribute to the enrichment of the medium with iron, hence stellar populations born at low redshifts exhibit higher iron abundances."," As time evolves SNe contribute to the enrichment of the medium with iron, hence stellar populations born at low redshifts exhibit higher iron abundances."
" This enrichment is stronger in scenarios Scl and Sc2, where the metallicity cut-offs are not so restrictive."," This enrichment is stronger in scenarios Sc1 and Sc2, where the metallicity cut-offs are not so restrictive."
 In Fig., In Fig.
 4 we present the distribution of [Si/Fe] for LGRB progenitors., \ref{cum_sife} we present the distribution of [Si/Fe] for LGRB progenitors.
" We observe that, for all our scenarios, they exibit a higher [Si/Fe] as redshift increases, indicating an enhacement in a elements at high redshifts."," We observe that, for all our scenarios, they exibit a higher [Si/Fe] as redshift increases, indicating an enhacement in $\alpha$ elements at high redshifts."
 This is consistent with the fact that SNIa and SNII enrich the interstellar medium., This is consistent with the fact that SNIa and SNII enrich the interstellar medium.
" Due to the fact that SNIa progenitors have lifetimes 1Gyr while those of SNII live only10Myr, at high redshift only the contribution of the latter to the interstelar medium enrichment is significant, rendering stellar populations rich in a elements."," Due to the fact that SNIa progenitors have lifetimes $\sim1{\rm Gyr}$ while those of SNII live only, at high redshift only the contribution of the latter to the interstelar medium enrichment is significant, rendering stellar populations rich in $\alpha$ elements."
" As redshift decreases, the contribution of SNIa becomes important, decreasing the abundance of a elements relative to iron."," As redshift decreases, the contribution of SNIa becomes important, decreasing the abundance of $\alpha$ elements relative to iron."
" The trend of [Si/Fe] to decrease towards lower redshift is also observed in Fig.5 (right panel), where we plot the median value of the ratio [Si/Fe] as a function of redshift."," The trend of [Si/Fe] to decrease towards lower redshift is also observed in \ref{medians_sife_feh} (right panel), where we plot the median value of the ratio [Si/Fe] as a function of redshift."
 In this figure we also observe that the values of [Si/Fe] of the LGRB progenitors is lower and evolve more strongly with redshift in the scenarios where the metallicity cut-off is more restrictive., In this figure we also observe that the values of [Si/Fe] of the LGRB progenitors is lower and evolve more strongly with redshift in the scenarios where the metallicity cut-off is more restrictive.
 In the left panel of Fig., In the left panel of Fig.
 5 we show the median value of the ratio [Fe/H] as a function of redshift., \ref{medians_sife_feh} we show the median value of the ratio [Fe/H] as a function of redshift.
 We find that the median value of [Fe/H] decreases with Ze as expected for cut-offs progressively more restrictive in metallicity., We find that the median value of [Fe/H] decreases with $Z_{\rm c}$ as expected for cut-offs progressively more restrictive in metallicity.
" These results indicate that the metallicity cut-off tends to eliminate old stellar populations highly enriched by SNII, located mainly in the central regions of galaxies and originated in first outbreaks of star formation."," These results indicate that the metallicity cut-off tends to eliminate old stellar populations highly enriched by SNII, located mainly in the central regions of galaxies and originated in first outbreaks of star formation."
 This interpretation agrees with the shift of the normalized impact parameters of LGRB progenitors observed in Fig. 1.., This interpretation agrees with the shift of the normalized impact parameters of LGRB progenitors observed in Fig. \ref{bnorm}.
" Aiming at understanding the relation between LGRBs and star formation, we analysed the spatial distribution and chemical abundances of stellar populations producing these phenomena."," Aiming at understanding the relation between LGRBs and star formation, we analysed the spatial distribution and chemical abundances of stellar populations producing these phenomena."
 We investigated four different scenarios for the progenitors of LGRBs based on the collapsar model with different metallicity cut-offs., We investigated four different scenarios for the progenitors of LGRBs based on the collapsar model with different metallicity cut-offs.
 We compared the spatial distribution of LGRBs within their HGs in our scenarios and with the available observations., We compared the spatial distribution of LGRBs within their HGs in our scenarios and with the available observations.
" We found that in all our scenarios LGRB progenitors reside on average in the outer regions of their galaxies at low redshifts, shifting toward the centre"," We found that in all our scenarios LGRB progenitors reside on average in the outer regions of their galaxies at low redshifts, shifting toward the centre"
in combination with the star-formation history obtained by simulation of a hierarchical assembly in a canonical ACDM cosmology (Lietal.2007) for the host galaxy of the second most distant (z= 6.42) quasar ever detected: SDSS JI]14845251.,"in combination with the star-formation history obtained by simulation of a hierarchical assembly in a canonical $\Lambda$ CDM cosmology \citep{Li07}
 for the host galaxy of the second most distant $z = 6.42$ ) quasar ever detected: SDSS J1148+5251."
 This galaxy appears to contain an unexpectedly high dust mass (given its redshift) which has raised doubts about whether stars alone can produce so much dust in such a short time., This galaxy appears to contain an unexpectedly high dust mass (given its redshift) which has raised doubts about whether stars alone can produce so much dust in such a short time.
" Valianteetal.(2009) argued that stellar sources ean account for the dust detected by observations and show also that the contribution of AGB stars to ""ust production in the early Universe may be far from negligible.", \citet{Valiante09} argued that stellar sources can account for the dust detected by observations and show also that the contribution of AGB stars to dust production in the early Universe may be far from negligible.
 Recently. Pipinoetal.2011). reached a conclusion similar to jat of. Valianteetal.(2009)... although their model includes also 22anonstellar sources of dust.," Recently, \citet{Pipino10} reached a conclusion similar to that of \citet{Valiante09}, although their model includes also nonstellar sources of dust."
 However. the model by Valianteetal.Ae2009) as well as that of Pipinoetal.(2011) assumes a gas mass qat is much larger (about an orderof magnitude) than the observec sesmass of molecular gas of ~107 (Walteretal.2004).," However, the model by \citet{Valiante09} as well as that of \citet{Pipino10} assumes a gas mass that is much larger (about an order of magnitude) than the observed mass of molecular gas of $\sim 10^{10} M_\odot$ \citep{Walter04}."
.. A more gas rich galaxy will obviously be A7;able to hold a greater mass of ocust. which led Gall(2010) to suggest that the gas content of many ust rich galaxies at high redshifts may be greatly underestimated —heir preferred models has a few times 10!!M. of gas after | Gyr).," A more gas rich galaxy will obviously be able to hold a greater mass of dust, which led \citet{Gall10thes} to suggest that the gas content of many dust rich galaxies at high redshifts may be greatly underestimated (their preferred models has a few times $10^{11} M_\odot$ of gas after 1 Gyr)."
" An intriguing discovery which has complicated the picture is qat not all quasar host galaxies at high redshifts show significan ""ust components.", An intriguing discovery which has complicated the picture is that not all quasar host galaxies at high redshifts show significant dust components.
 In fact. some of them may be essentially dus Tee (see.e.g.Jiangetal.2006).," In fact, some of them may be essentially dust free \citep[see, e.g.,][]{Jiang06}."
".. Quasars at high redshifts are powered by supermassive black holes with masses >10""M. (see.e.g..Vestergaard2004:Jiangetal. 2006)... which may be part of he explanation to why some quasar hosts are very dust rich while others are not."," Quasars at high redshifts are powered by supermassive black holes with masses $>10^9 M_\odot$ \citep[see, e.g.,][]{Vestergaard04,Jiang06}, which may be part of the explanation to why some quasar hosts are very dust rich while others are not."
 Jiangetal.(2010). have shown that there is a correlation between the estimated size of the central black hole anc he derived dust masses., \citet{Jiang10} have shown that there is a correlation between the estimated size of the central black hole and the derived dust masses.
 This finding may lend some support to a iypothesis presented by Elvisetal.(2002). that dust may form in significant quantities in the outflows of quasars.," This finding may lend some support to a hypothesis presented by \citet{Elvis02}, that dust may form in significant quantities in the outflows of quasars."
 In this paper it will be shown that the stellar dust production required to explain the observed dust mass in SDSS JI14845251 and in other high-z galaxies is higher than what appears reasonable taking both theoretical and observational constraints into account., In this paper it will be shown that the stellar dust production required to explain the observed dust mass in SDSS J1148+5251 and in other $z$ galaxies is higher than what appears reasonable taking both theoretical and observational constraints into account.
 Is some kind of “secondary” dust production. possibly associated with a quasar outflow. required to explain the detections of large dust masses?," Is some kind of 'secondary' dust production, possibly associated with a quasar outflow, required to explain the detections of large dust masses?"
 Or are the high dust-to-gas ratios at high redshifts simply an artefact due to incorrect calibration of conversion factors for gas and dust tracers?, Or are the high dust-to-gas ratios at high redshifts simply an artefact due to incorrect calibration of conversion factors for gas and dust tracers?
" Adopting a canonical ACDM cosmology. a closed form expression or the expansion scale factor S(2) can be obtained (Gron 2002).. where Hy is Hubble's constant. Qy, is the fractional matter energy density today. and OQ is the fractional dark energy density oday."," Adopting a canonical $\Lambda$ CDM cosmology, a closed form expression for the expansion scale factor $S(t)$ can be obtained \citep{Gron02}, , where $H_0$ is Hubble's constant, $\Omega_M$ is the fractional matter energy density today, and $\Omega_\Lambda$ is the fractional dark energy density today."
 Combining thisexpression with z*|=Stu/S() gives the ime-redshift relation Here. a standard scenario with Hy=70 km ! ον= and O4=0.7 is assumed.," Combining thisexpression with $z + 1 = S(t_0)/S(t)$ gives the time-redshift relation Here, a standard scenario with $H_0 = 70$ km $^{-1}$ $^{-1}$, $\Omega_M = 0.3$ and $\Omega_\Lambda = 0.7$ is assumed."
" Galaxy formation is assumed to start at >=14. which corresponds to 4,=290 Myr."," Galaxy formation is assumed to start at $z=14$, which corresponds to $t_{\rm a} \approx 290$ Myr."
 This choice of redshift is somewhat arbitrary. but the first galaxies are believed to have formed during the reionisation epoch. i.e.. between z=20 and σΞ6.," This choice of redshift is somewhat arbitrary, but the first galaxies are believed to have formed during the reionisation epoch, i.e., between $z = 20$ and $z = 6$."
 Galaxy formation at z=[4 may thus be representative for when the vast majority of galaxies formed in the early Universe., Galaxy formation at $z=14$ may thus be representative for when the vast majority of galaxies formed in the early Universe.
 According to the current cosmological paradigm (see Sect. 2.19. ," According to the current cosmological paradigm (see Sect. \ref{cosmology}) ),"
galaxies are formed by mergers with other galaxies and/or infall of baryons at a certain rate A7;., galaxies are formed by mergers with other galaxies and/or infall of baryons at a certain rate $\dot{ M}_{\rm i}$.
" The time scale of infall (mass assembly) may be defined as r,=M,/M,.", The time scale of infall (mass assembly) may be defined as $\tau_{\rm i} = {M_{\rm g} / \dot{ M}_{\rm i}}$.
 It is reasonable to assume the star-formation rate somehow reflects the rate of infall., It is reasonable to assume the star-formation rate somehow reflects the rate of infall.
" In particular. it is especially convenient (mathematically) to assume M,=oM,. which is equivalent to assuming that the gas mass is constant with respect to time (Larson1972)."," In particular, it is especially convenient (mathematically) to assume $\dot{ M}_{\rm i} = \alpha\dot{ M}_\star$, which is equivalent to assuming that the gas mass is constant with respect to time \citep{Larson72}."
 This simplistic model of the rate of infall will be used throughout this paper., This simplistic model of the rate of infall will be used throughout this paper.
 Dust in the ISM is expected to be destroyed by Kinetic-energy injection to the ISM by SNe. but dust grains may also in the ISM.particular. the outer parts of SN remnants.," Dust in the ISM is expected to be destroyed by kinetic-energy injection to the ISM by SNe, but dust grains may also in the ISM, the outer parts of SN remnants."
 Further dust may also take place in molecular clouds., Further dust may also take place in molecular clouds.
" However. dust destruction is probably dominating over nucleationISM). and it is thus fair to assume that introducing an rate of dust destruction in the ISM Aj, is sutficient in order to model the dust cycle."," However, dust destruction is probably dominating over nucleation, and it is thus fair to assume that introducing an rate of dust destruction in the ISM $\dot{M}_{\rm ISM}$ is sufficient in order to model the dust cycle."
" Following (2007) the dust destruction time-scale is where A is the gas mass. z/nsy is the effective gas mass cleared of dust by each SN event. and Asx is the SN rate. which may be approximated as where m, is the mass of the most massive SN-progenitors."," Following \citet{Dwek07} the dust destruction time-scale is where $M_{\rm g}$ is the gas mass, $m_{\rm ISM}$ is the effective gas mass cleared of dust by each SN event, and $R_{\rm SN}$ is the SN rate, which may be approximated as where $m_{\rm u}$ is the mass of the most massive SN-progenitors."
 The integral in Eq. (49) , The integral in Eq. \ref{snr}) )
"is a constant with repspect to time. hence the time scale Ty may be expressed as where ojs\, Will be referred to as the efficiency. of dust destruction in the ISM for each generation of stars."," is a constant with repspect to time, hence the time scale $\tau_{\rm d}$ may be expressed as where $\delta_{\rm ISM}$ will be referred to as the efficiency of dust destruction in the ISM for each generation of stars."
" For a normal IMF and zs=1000M; (thepreferredvaluefortheMilkyWay.seeDweketal.2007.andreferences thereiny.. this parameter is Os,x10.0."," For a normal IMF and $m_{\rm ISM} \approx 1000 M_\odot$ \citep[the preferred value for the Milky Way, see][and references therein]{Dwek07}, this parameter is $\delta_{\rm ISM} \approx 10.0$."
 When considering the contribution from AGB stars there is a time-lag due to the lifetimes of these stars and therfore the destruction of dust depends on how the SN rate evolves., When considering the contribution from AGB stars there is a time-lag due to the lifetimes of these stars and therfore the destruction of dust depends on how the SN rate evolves.
 If the rate of star formation has decreased suthcienly when the release of new dust from AGB stars happens. the dust destruction by SNe will be much less due to the lower SN rate.," If the rate of star formation has decreased sufficienly when the release of new dust from AGB stars happens, the dust destruction by SNe will be much less due to the lower SN rate."
 Effectively the parameter ois may have to be smaller than the number estimated above., Effectively the parameter $\delta_{\rm ISM}$ may have to be smaller than the number estimated above.
" Dwekeal.(2007). considered miss, a free parameter. and found that with Hs=LOOM: and a top-heavy IMF (shallowerthantheSalpeter1955.IMF) the estimated dust mass of SDSS 1114845251 could be explained in terms of dust production in SNe/high-mass stars."," \citet{Dwek07} considered $m_{\rm ISM}$ a free parameter, and found that with $m_{\rm ISM} = 100 M_\odot$ and a top-heavy IMF \citep[shallower than the][IMF]{Salpeter55} the estimated dust mass of SDSS J1148+5251 could be explained in terms of dust production in SNe/high-mass stars."
" Tf Hass,=00: is adopted as the ettective gas mass cleared of dusby a SN-event. then ois= 1.0."," If $m_{\rm ISM} = 100 M_\odot$ is adopted as the effective gas mass cleared of dustby a SN-event, then $\delta_{\rm ISM} \approx 1.0$ ."
 Depletion of dust in the ISM need not be due to any actua destruction of dust., Depletion of dust in the ISM need not be due to any actual destruction of dust.
 If a galactic wind is present it may have essentially the same etfect. although gas and atomic metals wil be lost too.," If a galactic wind is present it may have essentially the same effect, although gas and atomic metals will be lost too."
 Since a galactic wind is beleived to be mainly driven by, Since a galactic wind is beleived to be mainly driven by
tween 12.000 Ix ancl 13.000 Ix. and. when the temperature is higher. bolometric corrections appropriate for οΗ]= 1.5. based on the spectroscopical evidence provided by. Pace et al. (,"between 12,000 K and 13,000 K, and, when the temperature is higher, bolometric corrections appropriate for ${\rm [Fe/H]}=0.5$ , based on the spectroscopical evidence provided by Pace et al. ("
2006. see their Figure 4),"2006, see their Figure 4)."
 This is in principle a erude approximation. but extended (in both mass and initial chemical composition) grids of LB stellar evolution and. atmosphere mocdels that. include consistently the clleet of radiative levitation (see. e.g. Lui-Bon-Boa. LeBlanc Lauschilelt 2000. ancl Michaud. Richer Richarc.," This is in principle a crude approximation, but extended (in both mass and initial chemical composition) grids of HB stellar evolution and atmosphere models that include consistently the effect of radiative levitation (see, e.g. Hui-Bon-Boa, LeBlanc Hauschildt 2000, and Michaud, Richer Richard."
 2008 for [first results on this topic) ave still lacking., 2008 for first results on this topic) are still lacking.
 Also. it appears that additional mixing processes (e.g. meridional circulation) must be coupled. to the effect. of radiative levitation. in order to explain the abrupt. disappearance of the surface abundance anomalies at Zip LLO00-12000 Ix. (sce. e.g. Quievy et al.," Also, it appears that additional mixing processes (e.g. meridional circulation) must be coupled to the effect of radiative levitation, in order to explain the abrupt disappearance of the surface abundance anomalies at $T_{eff} $ 11000-12000 K (see, e.g., Quievy et al."
 2009)., 2009).
 lt is clear that our wav to simulate the elect of radiative levitation is just an approximation. that cannot take into account a possible stratification of the metal abundances in the atmosphere and below the photosphere.," It is clear that our way to simulate the effect of radiative levitation is just an approximation, that cannot take into account a possible stratification of the metal abundances in the atmosphere and below the photosphere."
 In the following we will however show how this treatment seems to. be reasonably appropriate., In the following we will however show how this treatment seems to be reasonably appropriate.
 Asa first step we determined a reference distance moculus and extinction. by fitting our WB models to the RIIB., As a first step we determined a reference distance modulus and extinction by fitting our HB models to the RHB.
 First. we produced synthetic samples of LIB stars. by following the techniques pioneered by Rood (1973: see also Ferraro et al.," First, we produced synthetic samples of HB stars by following the techniques pioneered by Rood (1973; see also Ferraro et al."
 1999) and compared separately the magnitude distributions of svnthetic RIIB stars in the £336. and POSOW filters to the observed ones for the stars populating the box displayed in the lower panel of Figure 5.., 1999) and compared separately the magnitude distributions of synthetic RHB stars in the $F336W$ and $F555W$ filters to the observed ones for the stars populating the box displayed in the lower panel of Figure \ref{teo1}.
 RILB stars ave too faint in the {ΟΠΗ filter to provide any additional constraint., RHB stars are too faint in the $F160BW$ filter to provide any additional constraint.
 We used the set of models with Yoo=0.248 [for the RIB stars., We used the set of models with $Y=0.248$ for the RHB stars.
" Phe synthetic samples have four. free parameters: extinction. distance modulus (02.Mo. mean HIB stellar mass iMgp? and its dispersion m:i44,;.."," The synthetic samples have four free parameters: extinction, distance modulus $(m-M)_{0}$, mean HB stellar mass $\langle M_{\rm HB} \rangle$ and its dispersion $\sigma_{\langle M_{\rm HB} \rangle}$."
 We started by randomly selecting the stellar mass Mgy from a Gaussian distribution centred around a given guess value of Αη). with lo dispersion σιμι.," We started by randomly selecting the stellar mass $M_{\rm HB}$ from a Gaussian distribution centred around a given guess value of $\langle M_{\rm HB} \rangle$, with $\sigma$ dispersion $\sigma_{\langle M_{\rm HB} \rangle}$."
 The WEPC2 magnitudes of the svnthetic star were determined. according to its position along the appropriate LB track with mass Aly (interpolated. when necessary. among the available set. of tracks in the BaSTI database) after an evolutionary time ἐς," The WFPC2 magnitudes of the synthetic star were determined according to its position along the appropriate HB track with mass $M_{\rm HB}$ (interpolated, when necessary, among the available set of tracks in the BaSTI database) after an evolutionary time $t$."
 We determined. / assuming that the stars reach the ZAUB with a constant rate., We determined $t$ assuming that the stars reach the ZAHB with a constant rate.
 We employed a Hat. probability distribution ranging [rom zero to /gi. where fap is the time spent along the LED. (we consider as the end ofthe LLB phase the time when the central He-abundance drops to zero).," We employed a flat probability distribution ranging from zero to $t_{\rm HB}$, where $t_{\rm HB}$ is the time spent along the HB (we consider as the end of the HB phase the time when the central He-abundance drops to zero)."
 The star with the lowest mass has the longest central He-burning lifetime. (Castellani et al., The star with the lowest mass has the longest central He-burning lifetime (Castellani et al.
 L994. Cassisi ct al.," 1994, Cassisi et al."
 2003). and sets fap., 2003) and sets $t_{\rm HB}$.
 This implies that for some masses the randomly selected value of/ will be longer that their HEB lifetime. or in other words that they have already evolved to the following evolutionary stages.," This implies that for some masses the randomly selected value of $t$ will be longer that their HB lifetime, or in other words that they have already evolved to the following evolutionary stages."
" We mocdifed these synthetic magnitudes accounting for distance modulus. (mAl), and extinction.", We modifed these synthetic magnitudes accounting for distance modulus $(m-M)_0$ and extinction.
 We used as broad guidelines the results from. previous investigations (c.g. Bedin ct al., We used as broad guidelines the results from previous investigations (e.g. Bedin et al.
" 2000). exploring a range of extintion values L£(21320.15 0.20. and distance modulus (imAd),= 15.015.5."," 2000), exploring a range of extintion values $E(B-V)=0.15$ –0.20, and distance modulus $(m-M)_0=15.0$ –15.5."
 We finally acicled a Gaussian ranconm error with le dispersion. equal to that obtained [rom the data reduction process.," We finally added a Gaussian random error with $\sigma$ dispersion, equal to that obtained from the data reduction process."
 We then compared the resulting svathetic clistribution o£ mοσο and mpaso values to the observed one. and modified the values of the four free parameters until," We then compared the resulting synthetic distribution of $m_{F555W}$ and $m_{F336W}$ values to the observed one, and modified the values of the four free parameters until"
the precision of the pulsc-tiuiug-based value determined bv Fiugeretal.(1999).,the precision of the pulse-timing-based value determined by \citet{fingeretal99}.
. IIence. we do not present an ASMP-based estimate of the period.," Hence, we do not present an ASM-based estimate of the period."
 was discovered in INTECRAL/IDIS observations of the Calactic Center reeion (Cheruvakovaetal.2003)., was discovered in /IBIS observations of the Galactic Center region \citep{chern03}.
. The orbital period was iudependcutly found around the same time by Levine&Corbet(2006) in the ASAI light curve aud by Seueractal.(2007) in the light curve., The orbital period was independently found around the same time by \citet{lcatel06} in the ASM light curve and by \citet{sguera07} in the light curve.
 The reported values for the period were 18.5540.03 davs and 18.52+0.01 davs. respectively.," The reported values for the period were $18.55 \pm 0.03$ days and $18.52 \pm 0.01$ days, respectively."
 Jainctal.(2000). have receutlv analyzed BAT aud ASAI data aud estimate the period to be Suvft/P=18.5182£0.0088 davs., \citet{jainetal09} have recently analyzed /BAT and ASM data and estimate the period to be $P = 18.5482 \pm 0.0088$ days.
 The period is detected with high significauce in the present study (see Figure 20))., The period is detected with high significance in the present study (see Figure \ref{fig:pds1848}) ).
 We use that yower spectrum to estimate that the orbitalperiod is P=18.515c0.0053[£0.03 days.," We use that power spectrum to estimate that the orbitalperiod is $P = 18.545 \pm 0.003\,[\pm 0.034]$ days."
 is the N-rav1] counterpart to the well-shown object433., is the X-ray counterpart to the well-known object.
" This object most likely consists of a stellar-niass black hole that is accreting at super-Eddington rates from a strong wind ολτος, bv a ligh-uass normal-type star. and produces jets wherein the outflow velocity is approximately 0.266."," This object most likely consists of a stellar-mass black hole that is accreting at super-Eddington rates from a strong wind emitted by a high-mass normal-type star, and produces jets wherein the outflow velocity is approximately $0.26 c$."
 An carly review nay befound in Margon(1981) aud a more recent oue in Fabrika(2001)., An early review may befound in \citet{margon84} and a more recent one in \citet{fabrika04}.
. The precession of the jets with a period of ~162 days is not ouly evident iu the moving optical lines. but also in other phenomena includius the apparent N-rav intensity (sceWenetal.2006:Ces 2002).," The precession of the jets with a period of $\sim
162$ days is not only evident in the moving optical lines, but also in other phenomena including the apparent X-ray intensity \citep[see][]{wen06,giesetal02}."
. À προς of reports discuss the evidence frou optical spectroscopy for Doppler shifts of tle compact and normal components aud their interpretation in ternis of orbital parameters and component masses, A number of reports discuss the evidence from optical spectroscopy for Doppler shifts of the compact and normal components and their interpretation in terms of orbital parameters and component masses.
"2010).. Coranskiietal.(1998) reported a precise nieasuremient of the orbital period of the binary. be. Poa,=3.0821140.00008 days. on the basis of the times of eclipse-like dips in the optical fix."," \citet{goranetal98}
 reported a precise measurement of the orbital period of the binary, i.e., $P_{orb} = 13.08211 \pm 0.00008$ days, on the basis of the times of eclipse-like dips in the optical flux."
 The A-rav intensity also shows eclipse-like dips (e.g...Filip-povaetal.2006:Cherepashchukct2009). but these have not vet been exploited for orbital period estimation.," The X-ray intensity also shows eclipse-like dips \citep[e.g.,][]{filipetal06,cheretal09} but these have not yet been exploited for orbital period estimation."
 (οςetal.(2002). reported the first evidence of the orbital period in the ASM N-vay intensity measurciicuts., \citet{giesetal02} reported the first evidence of the orbital period in the ASM X-ray intensity measurements.
" Very recently. Kubotaetal.(2010) obtained a new time of 1ininmun optical brightuess and used it. together with the earlier observatious. to obtain a revised value of the period. Py,=13.08:EO.00008 d. Iu the present analysis. we find unanbieuous evidence of both the precession and the orbital periods im the power spectrum (see Figure 21))."," Very recently, \citet{kubetal10} obtained a new time of minimum optical brightness and used it, together with the earlier observations, to obtain a revised value of the period, $P_{orb} = 13.08227 \pm 0.00008$ d. In the present analysis, we find unambiguous evidence of both the precession and the orbital periods in the power spectrum (see Figure \ref{fig:pds433}) )."
" Our estimate of the orbital period is P,,4=13.080c0.003[20.017] days."," Our estimate of the orbital period is $P_{orb} = 13.080 \pm 0.003\,[\pm 0.017]$ days."
 This value is consistent with. but somewhat less precise than. those of Corauskietal.(1998) and IWubotaetal.(2010).," This value is consistent with, but somewhat less precise than, those of \citet{goranetal98} and \citet{kubetal10}."
 If we fold the ASM Πο curve with the period and epoch of Kubotaetal.(2010).. we find that the time of N-rav 1iuimuun is coincident with the time of optical iininuin.," If we fold the ASM light curve with the period and epoch of \citet{kubetal10}, we find that the time of X-ray minimum is coincident with the time of optical minimum."
 The shape of the folded ASM light curve is similar to the eclipse-like shapes seen by (e.g...Cherepashichuketal.2009) and sueecsts that the eclipse-like decrease in the flux is the result of periodically occuring partial occultatious of the N-rav cluitting regions. broadly coustrued to inchide regions where the X-ray flux may be scattered (see Fig. 1)).," The shape of the folded ASM light curve is similar to the eclipse-like shapes seen by \citep[e.g.,][]{cheretal09}
 and suggests that the eclipse-like decrease in the flux is the result of periodically occurring partial occultations of the X-ray emitting regions, broadly construed to include regions where the X-ray flux may be scattered (see Fig. \ref{fig:fold2}) )."
 is a ligh-mass X-ray binary which manifests a distinct variation at a period of Pz9.6 days in the ASM light curve (Corbet&Pecle 2001).., is a high-mass X-ray binary which manifests a distinct variation at a period of $P \approx 9.6$ days in the ASM light curve \citep[][]{corbetpeele01}. .
 This source is iucluded iu this report even though the ~9.6- , This source is included in this report even though the $\sim 9.6$ 
Bringing all of these data together. we calculate sin7 for cach star as where P is the measured rotation period and Πο is the stellar radius (the subscript indicates that the measured radius is dependent on the stellar huuinosity. which depends on the assuned cluster distance. D).,"Bringing all of these data together, we calculate $\sin i$ for each star as where $P$ is the measured rotation period and $R_{D}$ is the stellar radius (the subscript indicates that the measured radius is dependent on the stellar luminosity, which depends on the assumed cluster distance, $D$ )."
 Eq., Eq.
| 1 allows us to Oogenerate an observed distribution of sin/ that is dependent on the adopted cluster cistance., \ref{eq:sini} allows us to generate an observed distribution of $\sin i$ that is dependent on the adopted cluster distance.
 We then model the distribution of sin/ assunüug that the rotational axes of stars in the cluster are raucdomly oriented., We then model the distribution of $\sin i$ assuming that the rotational axes of stars in the cluster are randomly oriented.
 By scaling the input cluster distance so that the observed siu; distribution matches the predicted distribution. we obtain an estimate for the true. cluster distance.," By scaling the input cluster distance so that the observed $\sin i$ distribution matches the predicted distribution, we obtain an estimate for the true cluster distance."
 We beein in 62 with a description of the data used to determine the distance to NGC 2261: 583. describes our ineasuremient of projected rotation velocities by applving a cross-correlation routine to high resolution spectra of cluster imenbers: &L describes the distance deteriunuation technique in detail: our results are discussed i ] 51.6..," We begin in $\S$ \ref{data} with a description of the data used to determine the distance to NGC 2264; $\S$ \ref{vsinidetermination}
 describes our measurement of projected rotation velocities by applying a cross-correlation routine to high resolution spectra of cluster members; $\S$ \ref{distancedetermination} describes the distance determination technique in detail; our results are discussed in $\S$ \ref{discussion}."
 As described above. our distance Tüeasurement relics on sin’ values for stars m NGC 2261.," As described above, our distance measurement relies on $\sin i$ values for stars in NGC 2264."
 Four types of data are needed to calculate siny for an individual star using Eq. 1z, Four types of data are needed to calculate $\sin i$ for an individual star using Eq. \ref{eq:sini}:
 the stars period. luuinosity. effective temperature and projected equatorial rotational velocity. esni.," the star's period, luminosity, effective temperature and projected equatorial rotational velocity, $v \sin i$."
 Period (52.1)). effective temperature (52.2)) and huninosity (52.3)) data were obtained from a catalog compiledby Rebulletal.(2006) anc are briefly described below.," Period $\S$ \ref{periodsub}) ), effective temperature $\S$ \ref{teffsub}) ) and luminosity $\S$ \ref{lumsub}) ) data were obtained from a catalog compiled by \citet{Rebull} and are briefly described below."
" Most of the esin/ data. ou the other haud. were calculated from spectra (52.1)) uxing a cross-correlatiou technique further discussed im οι, "," Most of the $v \sin i$ data, on the other hand, were calculated from spectra $\S$ \ref{spectrasub}) ) using a cross-correlation technique further discussed in $\S$ \ref{vsinidetermination}."
"Rotation periods for pre-auaiun sequence stars are deteriuuued by iueasumiug periodic wariatious im the objects"" brightness.", Rotation periods for pre-main sequence stars are determined by measuring periodic variations in the objects' brightness.
 These variations arise from the presence of large (~LOY angular radius) starspots ou the surfaces of these vouug. magnetically active stars (Herbst 1989).," These variations arise from the presence of large $\sim 40^{\circ}$ angular radius) starspots on the surfaces of these young, magnetically active stars \citep{Herbst}."
. The transit of the starspot(s) as tle star rotates diminishes the observed stellar hunuinositv ou the order of a few tenths of a magnitude., The transit of the starspot(s) as the star rotates diminishes the observed stellar luminosity on the order of a few tenths of a magnitude.
 This variation is quite stable and can be used to derive precise periods., This variation is quite stable and can be used to derive precise periods.
 Rebulletxeal.(2006). compiled periods measured for members of 2261 bv Rebulletal.(2002). Malsidouetal.(2001). and Lanunetal. (2005).," \citet{Rebull} compiled periods measured for members of NGC 2264 by \citet{Rebull_periods}, \citet{Makidon} and \citet{Lamm}. ."
. The period lueasurements range from roughly 0.5 to 29 davs and are predicted to be accurate to roughly 6P/P~1, The period measurements range from roughly 0.5 to 29 days and are predicted to be accurate to roughly $\delta P / P \approx 1\%$.
" The catalog assembled by Rebullsolutionetal.(2006). includes Typ, estimates derived from low-re spectral types and from the stars’ optical colors using the V—7 vs. Toy, relation presented. byIülleubraud.(1997).", The catalog assembled by \citet{Rebull} includes $T_{eff}$ estimates derived from low-resolution spectral types and from the stars' optical colors using the $V-I$ vs. $_{eff}$ relation presented by\citet{Hillenbrand}.
 Each stars WoFf color was dereddeued prior to this calculation as described by Rebulletal.(2002): stars with spectra were dereddened so that their observed RF colors matched the intrinsic colors of that spectral type on the zero age main sequence (ZÀMS) defined bv Bessell(1991).. Leeectt(1992) and. Leeecttctal.(1998).," Each star's $V-I$ color was dereddened prior to this calculation as described by \citet{Rebull_periods}: stars with spectra were dereddened so that their observed $R-I$ colors matched the intrinsic colors of that spectral type on the zero age main sequence (ZAMS) defined by \citet{Bessell1991}, \citet{Leggett1992}
 and \citet{Leggett1998}."
. The V.£ colors for stars without spectra were dereddened asstuning the modal reddening of members witli measired spectral types., The $V-I$ colors for stars without spectra were dereddened assuming the modal reddening of members with measured spectral types.
 While the use of an overall reddening for the cluster is less than ideal. out of a otal of OF stars for which we derive siu values. ouly ll lack spectral types.," While the use of an overall reddening for the cluster is less than ideal, out of a total of 97 stars for which we derive $\sin i$ values, only 14 lack spectral types."
 Thus. larger errors associated with photometric based Zpy measurements likely have a icelieible effect on our results.," Thus, larger errors associated with photometric based $T_{eff}$ measurements likely have a negligible effect on our results."
 T Tauri stars almost certainly do not have a single xiotospherie temperature. however: observations of the weak T Tauri star V110 by Herbst(1989). indicated the oeseuce of two large. polar star spots with characteristic eniperatures of 3100 IS. in coustrast to the stars E100 I. ohotosphiere.," T Tauri stars almost certainly do not have a single photospheric temperature, however; observations of the weak T Tauri star V410 by \citet{Herbst} indicated the presence of two large, polar star spots with characteristic temperatures of 3100 K, in constrast to the star's 4400 K. photosphere."
 Observations of larger ensembles of T Tauri stars indicate similar plotosphericspot teniperature differentials. and typical spot covering fractious of ~ (Bouvier&Bertout1989:Johus-IxrullCatford2002).," Observations of larger ensembles of T Tauri stars indicate similar photospheric–spot temperature differentials, and typical spot covering fractions of $\sim$ \citep{Bouvier1989,
Johns-Krull2002}."
. Somewhat counterntuitivebv. however. starspots are unlikely to introduce large errors dn ao stars derived plotospheric temperature. particularly for laree spotphotosphnere temperature cliffercutials.," Somewhat counterintuitively, however, star–spots are unlikely to introduce large errors in a star's derived photospheric temperature, particularly for large spot–photosphere temperature differentials."
 The stars inteerated enüssion is dominated by non-spotted photospheric flux. as the cooler spots cuit sieuificantly less flux per unit area.," The star's integrated emission is dominated by non-spotted photospheric flux, as the cooler spots emit significantly less flux per unit area."
 Observational confirmation of this effect was provided by Frascactal.(2005).. who coustructed detailed starspot models of RS ονι systems to reproduce broadbaud plotometric ight curves and temperature seusitive liue ratios from eniporallv resolved. Ligh resolution spectra of RS CV systems.," Observational confirmation of this effect was provided by \citet{Frasca2005}, who constructed detailed star–spot models of RS CVn systems to reproduce broadband photometric light curves and temperature sensitive line ratios from temporally resolved, high resolution spectra of RS CVn systems."
 Wlile the best fit models identified wo Frascaetal.(2005) possessed. plotosplerespot eniperature differentials of— —1000 Ix. the temperature seusitive spectroscopic lineratios ouly departed froin the shotospheric value by ~150 EK over the course of the observations.," While the best fit models identified by \citet{Frasca2005} possessed photosphere–spot temperature differentials of $\sim$ 1000 K, the temperature sensitive spectroscopic line–ratios only departed from the photospheric value by $\sim$ 150 K over the course of the observations."
 Iu their Appendix A. Frascaetal.(2005) outline a ormalism for calculating the mean temperature observed roni a star with a eiven spot coveriue fraction aud shotospherespot temperature differential.," In their Appendix A, \citet{Frasca2005} outline a formalism for calculating the mean temperature observed from a star with a given spot covering fraction and photosphere–spot temperature differential."
 Using this ornalisn. we calculated the difference between a T Tami stars truce (nou-spotted) photospheric teniperature and the temperature that would be measured from its integrated (spot | plotsphere) spectrum," Using this formalism, we calculated the difference between a T Tauri star's true (non-spotted) photospheric temperature and the temperature that would be measured from its integrated (spot $+$ photsphere) spectrum."
" For typical T Tauri star parameters (D,4,:71000 IK: Toe 3200 I fao; = Ud) the temperature measured from the conibined | photosphere Tbaud spectiuui would be 3992 IK. oulv 8 EK different from the “truce” plotospheric teniperature."," For typical T Tauri star parameters $_{phot}$ =4000 K; $_{spot}$ = 3200 K; $_{spot}$ = 0.1), the temperature measured from the combined $+$ photosphere I–band spectrum would be 3992 K, only 8 K different from the `true' photospheric temperature."
" Given the small size of this effect. we assume the ""uncertainties im our derived temperatures are dominated bv errors in the observed colors. spectral types. aud reddening corrections."," Given the small size of this effect, we assume the uncertainties in our derived temperatures are dominated by errors in the observed colors, spectral types, and reddening corrections."
" We expect the temperature uncertainties iutroduced w these effects areroughly of,-T,Hfah R", We expect the temperature uncertainties introduced by these effects areroughly $\delta T_{eff} / T_{eff} \approx 5 \%$ .
ebulletal.(2006) calculated huuinosities for cluster iienibers from dereddeued photometry. assuiiug a distance to the cluster of τοῦ pe.," \citet{Rebull} calculated luminosities for cluster members from dereddened photometry, assuming a distance to the cluster of 760 pc."
 Stars without previously determuned spectral types were dereddened by the modal reddening value determined for stars with spectral types., Stars without previously determined spectral types were dereddened by the modal reddening value determined for stars with spectral types.
 Not taking iuto account the uucertaimtv, Not taking into account the uncertainty
and the analytic expressions of the following functions: and,and the analytic expressions of the following functions: and
"A semi-analytic solution for the thin-disc equation 34 with the boundary condition in equation 33 was derived by Tanaka(2011),, for a finite boundary R4>0 and a special viscosity prescription vοςR”.","A semi-analytic solution for the thin-disc equation \ref{eq:diff}
 with the boundary condition in equation \ref{eq:bc} was derived by \cite{Tanaka11}, for a finite boundary $R_{\lambda}>0$ and a special viscosity prescription $\nu\propto R^{n}$."
" The solution can be written in the form Rt, Ry) Q0)aR’, where G(R,R';t,RA) is the Green's function specific to the boundary condition and the chosen value of the viscosity power-law index n; and Yinit(R) is an arbitrary initial density profile."," The solution can be written in the form; t, ) ), where $G(R,R^{\prime}; t, R_{\lambda})$ is the Green's function specific to the boundary condition and the chosen value of the viscosity power-law index $n$; and $\Sigma_{\rm init}(R)$ is an arbitrary initial density profile."
" We find that our discs satisfy vxR°4 inside R«105GM/c?, and thus adopt n= 0.4."," We find that our discs satisfy $\nu\propto R^{0.4}$ inside $R<10^{3}GM/c^{2}$, and thus adopt $n=0.4$ ."
" In order to model a thin accretion disc around a driven SMBH binary, we modify the solution of Tanaka(2011) in two ways."," In order to model a thin accretion disc around a GW-driven SMBH binary, we modify the solution of \cite{Tanaka11} in two ways."
" First, we derive a more general Green's function to allow for a boundary condition with nonzero mass flux across the inner boundary: (Ra,t) (Ry,t)."," First, we derive a more general Green's function to allow for a boundary condition with nonzero mass flux across the inner boundary: ,t) ,t)."
" Above, Ms;=3xv3 is the standard accretion rate for expected of a steady-state disc, and 0Xficak<lisa numerical factor representing the incomplete suppression of gas inflow into the cavity."," Above, $\dot{M}_{\rm ss}=3\pi\nu\Sigma$ is the standard accretion rate for expected of a steady-state disc, and $0\le f_{\rm leak}<1$ is a numerical factor representing the incomplete suppression of gas inflow into the cavity."
 The case fieax=0 corresponds to total suppression of accretion by the binary’s tidal torques., The case $f_{\rm leak}=0$ corresponds to total suppression of accretion by the binary's tidal torques.
" The boundary condition in equation 36 is motivated by results from numerical simulations, which show that in general the binary torques do not completely prevent accretion into the gap, but rather allow some gas to leak into the centre of the disc with a suppressed mass flux fieakS0.1 (eg.,Artymowicz&Lubow1996;Günther,Fadyen&Milosavljevié 2008)."," The boundary condition in equation \ref{eq:bc2} is motivated by results from numerical simulations, which show that in general the binary torques do not completely prevent accretion into the gap, but rather allow some gas to leak into the centre of the disc with a suppressed mass flux $f_{\rm leak}\ltsim 0.1$ \citep[e.g.,
][]{ArtLub96, Gunther+04, Ochi+05, MM08}."
. We choose ficak=0.1 as our fiducial value.," We choose $f_{\rm
 leak}=0.1$ as our fiducial value."
 The long-term behavior of the gas is to pile up near the cavity and satisfy the power-law vXοςRtfieak—1)/? in the vicinity of the boundary., The long-term behavior of the gas is to pile up near the cavity and satisfy the power-law $\nu\Sigma \propto R^{(f_{\rm leak}-1)/2}$ in the vicinity of the boundary.
" In comparison, a state around a solitary central mass disc satisfies v3i=constant, and a boundary condition imposing zero inward mass flux satisfies vicR'!/? at the boundary."," In comparison, a steady-state around a solitary central mass disc satisfies $\nu\Sigma={\rm constant}$ , and a boundary condition imposing zero inward mass flux satisfies $\nu\Sigma\propto R^{-1/2}$ at the boundary."
" Note, however, that the value of ficak does not have a strong effect on the mass profile (and hence the luminosity produced) outside the cavity-opening radius R); the fractional surface density enhancement due to secondary-dominated migration is typically of order unity."," Note, however, that the value of $f_{\rm leak}$ does not have a strong effect on the mass profile (and hence the luminosity produced) outside the cavity-opening radius $R_{\lambda}$; the fractional surface density enhancement due to secondary-dominated migration is typically of order unity."
" The gas that enters the cavity does so in nearly radial orbits MacFadyen&Milosavljevié(2008),, and so is dynamically decoupled from the circumbinary disc."," The gas that enters the cavity does so in nearly radial orbits \cite{MM08}, and so is dynamically decoupled from the circumbinary disc."
" Thus, the surface density and mass flux inside R can consistently be disregarded in the Green’s function formalism."," Thus, the surface density and mass flux inside $R_{\lambda}$ can consistently be disregarded in the Green's function formalism."
" The leaked gas can presumably form accretion discs around one or both SMBHs (Hayasaki,Mineshige&Sudou2007),, presumably at the usual AGN radiative efficiency ~0.1."," The leaked gas can presumably form accretion discs around one or both SMBHs \citep{Hayasaki+07}, presumably at the usual AGN radiative efficiency $\sim 0.1$."
" The mass supply rate of such circum-secondary (or -primary) discs will be modulated by freakMes, which decreases as the binary outruns the circumbinary gas."," The mass supply rate of such circum-secondary (or -primary) discs will be modulated by $f_{\rm leak}\dot{M}_{\rm ss}$, which decreases as the binary outruns the circumbinary gas."
" Because the viscous time at the outer edge of such discs are shorter than at R), they will be nearly steady-state, with the surface density profile at any given time being determined by the instantaneous mass flux into the cavity."," Because the viscous time at the outer edge of such discs are shorter than at $R_{\lambda}$, they will be nearly steady-state, with the surface density profile at any given time being determined by the instantaneous mass flux into the cavity."
" Thus, the bolometric luminosities of the discs around each disc may be roughly expressed as L«0.1fieakMes(R)c?, and would be Eddington-limited by the potential of the individual black holes they orbit."," Thus, the bolometric luminosities of the discs around each disc may be roughly expressed as $L< 0.1~f_{\rm leak}\dot{M}_{\rm ss}(R_{\lambda})c^{2}$, and would be Eddington-limited by the potential of the individual black holes they orbit."
" This suggests that if the quantity fieakm exceeds unity, the region inside the cavity would develop radiation-driven outflow winds."," This suggests that if the quantity $f_{\rm leak}\dot{m}$ exceeds unity, the region inside the cavity would develop radiation-driven outflow winds."
" Thus, the parameter ficak affects the energetic output due to accretion inside the cavity far more than it does that of the circumbinarydisc."," Thus, the parameter $f_{\rm leak}$ affects the energetic output due to accretion inside the cavity far more than it does that of the circumbinarydisc."
 The Green’s function for the boundary condition in, The Green's function for the boundary condition in
neutral absorption. together with a Gaussian-shaped narrow (oc=0.01 keV) Fe Ix line at 6.4 keV. in the rest. frame of the quasar.,"neutral absorption, together with a Gaussian-shaped narrow $\sigma=0.01$ keV) Fe K line at 6.4 keV in the rest frame of the quasar."
 “Phe absorption column that was fitted has two components: (1) the Galactic column density which was fixec in value in the z=0 frame. and (ii) the intrinsic absorption of the quasar in its rest frame. which was allowed to vary in value.," The absorption column that was fitted has two components; (i) the Galactic column density which was fixed in value in the z=0 frame, and (ii) the intrinsic absorption of the quasar in its rest frame, which was allowed to vary in value."
 The Galactic column clensities have been obtainec from the program. run under the system. which uses data from Stark (1992) arc other surveys: absorption cross-sections have been taken from Morrison MeCanmmon (1983).," The Galactic column densities have been obtained from the program, run under the system, which uses data from Stark (1992) and other surveys; absorption cross-sections have been taken from Morrison McCammon (1983)."
 Where possible more accurate values of (towards a particular quasar have been taken from individual observations using cither Elvis (1989) or Alurphy (1996). which use higher spatial resolution.," Where possible more accurate values of towards a particular quasar have been taken from individual observations using either Elvis (1989) or Murphy (1996), which use higher spatial resolution."
 The results from the spectral fitting are presented. in table 2., The results from the spectral fitting are presented in table 2.
 Phe llux is given over the range 0.5 keV to 10 keV (in the observed frame and extrapolated. where necessary). with luminosity quoted from 2-10. keV. (corrected. for absorption) in the quasar rest-frame.," The flux is given over the range 0.5 keV to 10 keV (in the observed frame and extrapolated where necessary), with luminosity quoted from 2-10 keV (corrected for absorption) in the quasar rest-frame."
 The value quoted for iis the measured value of intrinsic absorption in the equasar rest frame. in excess of anv Galactic absorption.," The value quoted for is the measured value of intrinsic absorption in the quasar rest frame, in excess of any Galactic absorption."
 The equivalent width G2QW) and lux of the iron I[uorescence line are given. with the EQW being in the quasar rest frame.," The equivalent width (EQW) and flux of the iron fluorescence line are given, with the EQW being in the quasar rest frame."
 The line parameters are those corresponding to a narrow line (σ=0.01 keV): the line energy is included as a [ree parameter in the model Gt unless otherwise stated in the table., The line parameters are those corresponding to a narrow line $\sigma=0.01$ keV); the line energy is included as a free parameter in the model fit unless otherwise stated in the table.
 Note also that the photon indices E are those of the underlying power-law continuum. with spectral features such as low energy absorption (cold or ionised). soft. N-rav excess. iron line emission or Compton reflection. taken into account where these features are significant.," Note also that the photon indices $\Gamma$ are those of the underlying power-law continuum, with spectral features such as low energy absorption (cold or ionised), soft X-ray excess, iron line emission or Compton reflection, taken into account where these features are significant."
 When the fitted spectrum is particularly complex. the 2-10 keV. band index has been used.," When the fitted spectrum is particularly complex, the 2-10 keV band index has been used."
 Llowever in most cases there is little dillerence between theundertiing 0.610 keV index and the 2-10 keV index., However in most cases there is little difference between the 0.6-10 keV index and the 2-10 keV index.
 All errors in the table are quoted at confidence. [or the appropriate number of interesting parameters.," All errors in the table are quoted at confidence, for the appropriate number of interesting parameters."
 For cach quasar the best fitting value for reduced. chi-squared. (x7 fdeerees of freedom) is given., For each quasar the best fitting value for reduced chi-squared $\chi^{2}$ /degrees of freedom) is given.
 “Phe values of ANA? quoted in table 2 are for adding an additional spectral feature to the fit. such as an intrinsic absorption column or Fe line.," The values of $\Delta\chi^{2}$ quoted in table 2 are for adding an additional spectral feature to the fit, such as an intrinsic absorption column or Fe line."
 For adding an absorption column. 1 additional parameter is added to the fit and for an iron line. 1 or 2 parameters are added. depending on whether the line energy is fixed at 6.4 keV or not (i.e. the line normalization and line cnerey).," For adding an absorption column, 1 additional parameter is added to the fit and for an iron line, 1 or 2 parameters are added, depending on whether the line energy is fixed at 6.4 keV or not (i.e. the line normalization and line energy)."
 As a roughe ὃνguide. Ay?=2.7 corresponds to significance for the addition of1 interesting parameter. with AA?=4.6 corresponding to significance for 2 additional interesting parameters.," As a rough guide, $\Delta\chi^{2}=2.7$ corresponds to significance for the addition of 1 interesting parameter, with $\Delta\chi^{2}=4.6$ corresponding to significance for 2 additional interesting parameters."
 Only. measurements of aand the iron line which are at significance or better have been quoted as best fit values in table 3.2: otherwise upper-limits only are quoted., Only measurements of and the iron line which are at significance or better have been quoted as best fit values in table 3.2; otherwise upper-limits only are quoted.
 F-tests have also been performed. (see Bevington Robinson 1992). both on the intrinsic absorption and the iron. line. to test the significance οἱ these features.," F-tests have also been performed (see Bevington Robinson 1992), both on the intrinsic absorption and the iron line, to test the significance of these features."
 If the value. obtained for ‘3.0. fortheadditionoflerirainlterestingparamoceter.thentheresullhassigni | ," If the value obtained for $>3.0$, for the addition of 1 extra interesting parameter, then the result has significance greater than $\sim$."
This is used. as a criterion for including line or absorption features in the spectral fitting., This is used as a criterion for including line or absorption features in the spectral fitting.
 Actual spectral features such as an Le line or an absorption column are only regarded as significant if these F-test criteria are met., Actual spectral features such as an Fe line or an absorption column are only regarded as significant if these F-test criteria are met.
 The X-ray emission in the sample covers a wide range of photon index. from hard. (e.g. PISS 22511113. EP= 0.95) to soft (e.g. ΙΔ. P— 2.31).," The X-ray emission in the sample covers a wide range of photon index, from hard (e.g. PKS 2251+113, $\Gamma=0.95$ ) to soft (e.g. IZWI, $\Gamma=2.37$ )."
 The distribution of 2-10 photon index for the quasars in the sample is illustrated in, The distribution of 2-10 photon index for the quasars in the sample is illustrated in
This is the degeneracy identified by Gouldetal.(1994).,This is the degeneracy identified by \citet{gmb}.
. Taking account of the second term in equation (13)) leaves a linear continuous (ie.. line-like) degeneracy in the πι: plane. bul rotates this line bv an angle. Since (vpically ya;«I. this angle is also usually very small. but it can be significant 1h some cases (see G.2)).," Taking account of the second term in equation \ref{eqn:C3}) ) leaves a linear continuous (i.e., line-like) degeneracy in the $\bpi_\e$ plane, but rotates this line by an angle, Since typically $\eta\pi_j\ll 1$, this angle is also usually very small, but it can be significant in some cases (see \ref{sec:implications}) )."
 IPC) can be measured. (hen (he continuous degeneracy is broken. but it is replaced by a [our-fold cliserete degeneracy.," If $C_4$ can be measured, then the continuous degeneracy is broken, but it is replaced by a four-fold discrete degeneracy."
 In this case. all four solutions lie along the line of the continuous degeneracy described in (he previous paragraph.," In this case, all four solutions lie along the line of the continuous degeneracy described in the previous paragraph."
 The transition from (he continuous to the discrete clegeneracy is itself continuous., The transition from the continuous to the discrete degeneracy is itself continuous.
" If C, is measured. but with only modest significance. then (he error ellipses for each of (he four solution will be extremely elongated along the line ol continuous clegeneracy."," If $C_4$ is measured, but with only modest significance, then the error ellipses for each of the four solution will be extremely elongated along the line of continuous degeneracy."
 In (his ease. the two solutions with wy)>0 may merge.," In this case, the two solutions with $u_0>0$ may merge."
 Similarly for the (wo solutions with wy«0., Similarly for the two solutions with $u_0<0$.
 However. the uy20 and wy<0 solutions cannot merge with each other unless uj is consistent with zero.," However, the $u_0>0$ and $u_0<0$ solutions cannot merge with each other unless $u_0^2$ is consistent with zero."
 To break the discrete degeneracy requires sensitivity to the higher order terms. which is (wpically obtained only in relatively long events.," To break the discrete degeneracy requires sensitivity to the higher order terms, which is typically obtained only in relatively long events."
 For which microlensing events is (the jerk-parallax degeneracy likely to be important?, For which microlensing events is the jerk-parallax degeneracy likely to be important?
" To address (his question. it is best to think in terms of the ""inverse projected velocity. A. (1 choose “AY for this quantity because it looks like an inverted v.)"," To address this question, it is best to think in terms of the “inverse projected velocity”, $\bLambda$ (I choose $\bLambda$ ” for this quantity because it looks like an inverted $\bv$ ”.)"
 Since /;j is virtually the same for all solutions. A is basically just a linear rescaling of πι.," Since $t_\e$ is virtually the same for all solutions, $\bLambda$ is basically just a linear rescaling of $\bpi_\e$."
 Its usefulness derives from the [act that (he projected velocity (ancl so its inverse) is related solely to the kinematics of the event and (hus is independent of the mass., Its usefulness derives from the fact that the projected velocity (and so its inverse) is related solely to the kinematics of the event and thus is independent of the mass.
 Expressed in terms of Chis «quantity. equation (21)) becomes where. (Note (hat since e andj are 2-dimensional vectors.the denominator is a signed scalar.," Expressed in terms of this quantity, equation \ref{eqn:pisols}) ) becomes where, (Note that since $\balpha$ and $\bj$ are 2-dimensional vectors,the denominator is a signed scalar."
 See also eq. |13]], See also eq. \ref{eqn:C3}]
 See also eq. |13]]., See also eq. \ref{eqn:C3}]
 See also eq. |13]].), See also eq. \ref{eqn:C3}]
~1.5 ~LO0 (£2>25 (<L 2=0.835 ~300 (<LOO (8 op=0.05 σῃ=0.01 z15 517. ,"$\sim$ $\sim$ $R
\age 25$ $\ale 1$ $z=0.835$ $\sim$ $\ale 100$ $R$ $\sigma_B = 0.05$ $\sigma_R = 0.01$ $\approx 15$ $51^\circ$ "
start to merge into a single system (e.g.?)..,start to merge into a single system \citep[e.g.][]{karl08}.
 The new FIR observations. along with the available radio and CO data. suggest that the gravitational forces have shuffled a huge amount of interstellar material along an are connecting the two galaxy nuclei.," The new FIR observations, along with the available radio and CO data, suggest that the gravitational forces have shuffled a huge amount of interstellar material along an arc connecting the two galaxy nuclei."
 Several kpe-sized. large ISM complexes house the most active star formation sites and evolve independently.," Several kpc-sized, large ISM complexes house the most active star formation sites and evolve independently."
 With respect to their position on the optical image (Fig. 2)).," With respect to their position on the optical image (Fig. \ref{pacs100contoursonHSTBVHalpha}) ),"
 knots K2 and K3 appear to be located directly at the collision front of the two disks embedded in dense dust lanes but already associated with bright star clusters., knots K2 and K3 appear to be located directly at the collision front of the two disks embedded in dense dust lanes but already associated with bright star clusters.
 Knot K1 is located more in the disk of NGC4039. while knot K4 is located in the disk of 44038.," Knot K1 is located more in the disk of NGC4039, while knot K4 is located in the disk of 4038."
 Knot K4 is a star formation area burning on a relatively low level., Knot K4 is a star formation area burning on a relatively low level.
 No CO emission is detected and there is no noticeable dust lane. so that it appears to hàve consumed its fuel and will start to extinguish soon.," No CO emission is detected and there is no noticeable dust lane, so that it appears to have consumed its fuel and will start to extinguish soon."
 Knot K2 reaches 2/3 of the luminosity of ΚΙ. however the derived star formation. densities. and efficiencies are considerably lower than for ΚΙ.," Knot K2 reaches 2/3 of the luminosity of K1, however the derived star formation densities and efficiencies are considerably lower than for K1."
 The spatially resolved SEDs obtained from the PACS maps clarify that it is more active tha was suggested by ? on the basis of the ISO and the 450 anc 850 um SCUBA maps., The spatially resolved SEDs obtained from the PACS maps clarify that it is more active than was suggested by \citet{haas00} on the basis of the ISO and the 450 and 850 $\mu$ m SCUBA maps.
 If K2 suffered from higher obscuratiot than ΚΙ (see Fig. 4)).," If K2 suffered from higher obscuration than K1 (see Fig. \ref{knotanalysis}) ),"
 the intrinsic difference between K2 anc KI would become smaller., the intrinsic difference between K2 and K1 would become smaller.
 There seems. however. to be a dichotomy between K2a and K2b-c. While the latter appears to have a normal radio-to-FIR ratio. Κα shows an excess i radio emission. as already noted by ?..," There seems, however, to be a dichotomy between K2a and $+$ c. While the latter appears to have a normal radio-to-FIR ratio, K2a shows an excess in radio emission, as already noted by \citet{chyzy04}."
 It i$ very inconspicuous in the optical and is lacking X-ray emission (see?.Fig.9).., It is very inconspicuous in the optical and is lacking X-ray emission \citep[see][Fig. 9]{fabbiano01}.
 K2a might therefore be in a younger evolutionary stage thar K2b-+e. Knot KI is currently the most active area by far., K2a might therefore be in a younger evolutionary stage than $+$ c. Knot K1 is currently the most active area by far.
 This is manifested by the highest individual luminosity. which itself is às high as the Milky Way luminosity but concentrated in an area of kkpe diameter.," This is manifested by the highest individual luminosity, which itself is as high as the Milky Way luminosity but concentrated in an area of kpc diameter."
 It exhibits the highest star formation rate and efficiency. and the highest star formation density that consequently raises dust temperatures to a maximum.," It exhibits the highest star formation rate and efficiency, and the highest star formation density that consequently raises dust temperatures to a maximum."
 The high density of hot heating sources can explain the very high MIR emission., The high density of hot heating sources can explain the very high MIR emission.
 Its radio-to-FIR luminosity ratio deviates from the average in that there is a lack of radio emission. and at the peak position. there is no X-ray emission either (see 9)..," Its radio-to-FIR luminosity ratio deviates from the average in that there is a lack of radio emission, and at the peak position, there is no X-ray emission either \citep[see][Fig. 9]{fabbiano01}."
 These findings classify it as a very young complex., These findings classify it as a very young complex.
 New Herschel-PACS scan maps of the Antennae at 70. 100. and um provide a high spatial resolution complement to earlier MIR maps (ISOCAM. Spitzer-MIPS). high-resolution CO QO). and radio cem maps and thus allow detailed study of the star formation state in individual emission knots.," New -PACS scan maps of the Antennae at 70, 100, and $\mu$ m provide a high spatial resolution complement to earlier MIR maps (ISOCAM, -MIPS), high-resolution CO (1-0), and radio cm maps and thus allow detailed study of the star formation state in individual emission knots."
 We confirm that the highest star formation activity takes place in the overlap area concentrated in the two emission complexes ΚΙ and K2., We confirm that the highest star formation activity takes place in the overlap area concentrated in the two emission complexes K1 and K2.
 Our star formation diagnostics strongly indicates that these knots differ in evolutionary stage., Our star formation diagnostics strongly indicates that these knots differ in evolutionary stage.
wo sources is1,two sources is.
.154... Leeetal.(2009) found that IIo tends to uuderpredict he SER when compared with ΕΤ estimates., \citet{Lee09} found that $\alpha$ tends to underpredict the SFR when compared with FUV estimates.
" The fiud hat the ratio is a factor of 2 by SER ~0.00337,sx.1. around the range of the lowest SFR galaxics considered rere."," The find that the ratio is a factor of 2 by SFR $\sim 0.003 M_{\odot} \rm \, yr^{-1}$, around the range of the lowest SFR galaxies considered here."
 We examined the FUV SER for our sample of BCDs and fud that it is slightly lower than the Ta SFR., We examined the FUV SFR for our sample of BCDs and find that it is slightly lower than the $\alpha$ SFR.
 Thus. he effect found by Leeetal.(2009) does not explain he overproduction of X-ray. binarics in BC'Ds.," Thus, the effect found by \citet{Lee09} does not explain the overproduction of X-ray binaries in BCDs."
 Alirabeletal.(2011) have receutlv proposed that Hel-mass X-ray binaries in star-forming galaxies were an muportant source of heating aud reionization of the iuterealactie wedi at the epoch of reiouization., \citet{Mirabel11} have recently proposed that high-mass X-ray binaries in star-forming galaxies were an important source of heating and reionization of the intergalactic medium at the epoch of reionization.
 If DCDs are analogs to the unevolved ealaxics in the carly universe. then eulhauced X-ray binary production in BCDs would sugeest au enhanced impact of Nav Dinarics on the early thermal history of the universe over that predicted using X-ray versus SER relations determined using near-solar metallicity ealaxies.," If BCDs are analogs to the unevolved galaxies in the early universe, then enhanced X-ray binary production in BCDs would suggest an enhanced impact of X-ray binaries on the early thermal history of the universe over that predicted using X-ray versus SFR relations determined using near-solar metallicity galaxies."
 We thank the referee for verv useful conuueuts., We thank the referee for very useful comments.
 We acknowledge the usage of the IbvperLeda. database (http://leda.univ-Ivoul.fr)., We acknowledge the usage of the HyperLeda database (http://leda.univ-lyon1.fr).
 This research has made use of data obtained. from the Chandra Data Archive and software provided by the Chandra N-vay Center (CAC) in the application packages CIAO. ChIPS. aud Sherpa.," This research has made use of data obtained from the Chandra Data Archive and software provided by the Chandra X-ray Center (CXC) in the application packages CIAO, ChIPS, and Sherpa."
for fiu=10|.,"for $f_{\ast,\rm crit}=10^{-4}$."
 This range is consistent with the WAIAP results. also shown in the Figure.," This range is consistent with the WMAP results, also shown in the Figure."
 The effect of changing fui. is onlv to shift the 7. constraint up or down.," The effect of changing $f_{\ast,\rm crit}$, is only to shift the $\tau_{\rm es}$ constraint up or down."
 For instance. the constraint for fy=3xLO? is T=0.12tis. with no change in the n constraint.," For instance, the constraint for $f_{\ast,\rm crit}=3\times 10^{-5}$ is $\tau_{\rm es} = 0.12^{+0.02}_{-0.03}$, with no change in the $n$ constraint."
" The resulis for increasing f.,44 are in the opposite direction. decreasing the constraint on Te, bv e0.02 for a factor of 10 increase in. fig."," The results for increasing $f_{\ast,\rm crit}$ are in the opposite direction, decreasing the constraint on $\tau_{\rm es}$ by $\sim 0.02$ for a factor of 10 increase in $f_{\ast,\rm crit}$."
 In order to understand (he quasar constraints. let us retur to the constant efficiency case.," In order to understand the quasar constraints, let us return to the constant efficiency case."
" Consider (he following ""subset"" of the quasar data: Figure 1. also shows the effects of combining these three constraints.", Consider the following “subset” of the quasar data: Figure \ref{fig:mainresult} also shows the effects of combining these three constraints.
 The intersection (Q1+Q2+Q3) subset of data appears to adequately depict the upper limit of (he combined dataset., The intersection (Q1+Q2+Q3) subset of data appears to adequately depict the upper limit of the combined dataset.
" Ir particular. these three sets of data contain much of the ""information"" in the full V treatment."," In particular, these three sets of data contain much of the “information” in the full $\chi^2$ treatment."
 The lower limit of the combined dataset actually also contained within the constraint data in Q3. but changes as a function of zo445.," The lower limit of the combined dataset actually also contained within the constraint data in Q3, but changes as a function of $z_{\rm overlap}$."
 Since Q3 is (he union of all these constraints [or all values of zogp. the lower limit does not appear.," Since Q3 is the union of all these constraints for all values of $z_{\rm overlap}$, the lower limit does not appear."
 This illustration shows that given the three degrees for freedom O0. oy. and eis. the quasar data constrains to essentially a line in (his space.," This illustration shows that given the three degrees for freedom $\Omega_0$, $\sigma_8$, and $\epsilon_{\rm UV}$ , the quasar data constrains to essentially a line in this space."
 This is because for each £3. the data (Q1+Q2) at z6 and the data (Q3) al z4.5 provide strong. essentially independent. constraints on (he remaining parameters σε and eiy.," This is because for each $\Omega_0$, the data (Q1+Q2) at $z\sim 6$ and the data (Q3) at $z\sim 4.5$ provide strong, essentially independent, constraints on the remaining parameters $\sigma_8$ and $\epsilon_{\rm UV}$."
" Of course. the ""length and. ""width of the the line (i.e.. (he exact shape. as illustrated in Figure 1)) depend on the likelihood unction in detail."," Of course, the “length” and “width” of the the line (i.e., the exact shape, as illustrated in Figure \ref{fig:mainresult}) ) depend on the likelihood function in detail."
 Now let us return to the time-varving efficiency. case., Now let us return to the time-varying efficiency case.
" For each value of f,44. we have ol parameters n. T. eryo. and 21."," For each value of $f_{\ast,\rm crit}$, we have four parameters $n$, $\tau_{\rm es}$, $\epsilon_{\rm UV,0}$, and $A$."
" Ieuristicallv. we now have (hvee constraints on the output: the data (QI+Q2) al z~ 6. the data (Q3) al z~4.5. aud the sell-consisteney of τι, (value “in” = value ""out"")."," Heuristically, we now have three constraints on the output: the data (Q1+Q2) at $z\sim 6$ , the data (Q3) at $z\sim 4.5$, and the self-consistency of $\tau_{\rm es}$ (value “in” = value “out”)."
 Therefore. once again. we expect a “line” in the 4-dimensional parameter space (for fixed. f.i).," Therefore, once again, we expect a “line” in the 4-dimensional parameter space (for fixed $f_{\ast,\rm crit}$ )."
" As before. the""length and ""width"" of the the line will depend on the likelihood functionin detail."," As before, the“length” and “width” of the the line will depend on the likelihood functionin detail."
constants of amorphous corundum (Al»O3; porous) obtained by Begemann et al. (,constants of amorphous corundum $_2$ $_3$; porous) obtained by Begemann et al. (
1997).,1997).
 The radii of spherical dust grains have been assumed to be 0.1 wm uniformly., The radii of spherical dust grains have been assumed to be 0.1 $\mu$ m uniformly.
" We choose 10 um as the fiducial wavelength that sets the scale of the optical depth (710) and compute models for eleven optical depths (719 = 0.005, 0.01, 0.05, 0.1, 0.5, 1, 3, 7, 15, 30 and 40)."," We choose 10 $\mu$ m as the fiducial wavelength that sets the scale of the optical depth $\tau_{10}$ ) and compute models for eleven optical depths $\tau_{10}$ $=$ 0.005, 0.01, 0.05, 0.1, 0.5, 1, 3, 7, 15, 30 and 40)."
" For the central star, we assume that the luminosity is 104 Lo and the stellar blackbody temperature is 2500 K for Tio €3 and 2000 K for Tio >3."," For the central star, we assume that the luminosity is $10^4$ $L_{\odot}$ and the stellar blackbody temperature is 2500 K for $\tau_{10}$ $\leq 3$ and 2000 K for $\tau_{10}$ $> 3$."
" Also, we use the warm silicate dust grains for LMOA stars (7 models with 719 € 3) and the cold grains for HMOA stars (4 models τιο > 3)."," Also, we use the warm silicate dust grains for LMOA stars (7 models with $\tau_{10}$ $\leq 3$ ) and the cold grains for HMOA stars (4 models $\tau_{10}$ $> 3$ )."
 The track line of the theoretical model results for silicate dust with T. = 1000 K (TLS) on a 2CD can show various interesting behaviors because the 10 jm silicate feature changes from emission to absorption when the dust optical depth becomes larger (7:10 > 3)., The track line of the theoretical model results for silicate dust with $T_c$ = 1000 K (TLS) on a 2CD can show various interesting behaviors because the 10 $\mu$ m silicate feature changes from emission to absorption when the dust optical depth becomes larger $\tau_{10}$ $> 3$ ).
" The TLS for increasing optical depths show wide variations in the inclination (see Figs 4, 5, 6 and 7) including backward effects (see Figs 5 and 7)."," The TLS for increasing optical depths show wide variations in the inclination (see Figs 4, 5, 6 and 7) including backward effects (see Figs 5 and 7)."
 On the upper panel of Fig., On the upper panel of Fig.
" 5, the TLS in the horizontal direction ([9]—[12] colour) initially moves backward when the 10 yum silicate emission feature becomes more prominent then moves forward when the 10 um silicate feature changes from emission to absorption."," 5, the TLS in the horizontal direction $-$ [12] colour) initially moves backward when the 10 $\mu$ m silicate emission feature becomes more prominent then moves forward when the 10 $\mu$ m silicate feature changes from emission to absorption."
 On the upper panel of Fig., On the upper panel of Fig.
" 7, the TLS in the horizontal direction (Z—[9] colour) initially moves forward rapidly when the 10 um silicate emission feature becomes more"," 7, the TLS in the horizontal direction $L-$ [9] colour) initially moves forward rapidly when the 10 $\mu$ m silicate emission feature becomes more"
Figure 14. shows the mass accretion rate for the central black hole for simulations 81-894.,Figure \ref{fig:accretionhistory} shows the mass accretion rate for the central black hole for simulations S1-S4.
 Formally. the Figure indicates that our simulations provide a sustained. super-Ecelington accretion rate (the Ecldineton accretion rate for Ser A* being 0.03M. vr L1). in some cases for between 10 to 107 ves.," Formally, the Figure indicates that our simulations provide a sustained super-Eddington accretion rate (the Eddington accretion rate for Sgr A* being $\sim 0.03 \msun$ $^{-1}$ ), in some cases for between $10^{4}$ to $^5$ yrs."
 However it must. be remembered. that we do not resolve σας cdvnamices inside the accretion radius., However it must be remembered that we do not resolve gas dynamics inside the accretion radius.
 We expect that material will form a dise there. ancl accretion will proceed. viscouslv.," We expect that material will form a disc there, and accretion will proceed viscously."
 The viscous time scale. Ao. depends on the temperature in the dise midplane and the viscosity parameter. a (2)..," The viscous time scale, $t_{\rm visc}$, depends on the temperature in the disc midplane and the viscosity parameter, $\alpha$ \citep{Shakura73}. ."
 Phe midplane temperature in the inner areseconed is around 107A in both the standard (aonself-eravitating.e.g.7) andthe sclf-eravitating regimes (?).. vielding 44/1?~0.01.," The midplane temperature in the inner arcsecond is around $10^3 K$ in both the standard \citep[non self-gravitating,
e.g.,][]{NC05} andthe self-gravitating regimes \citep{Nayakshin06a}, yielding $H/R \sim 0.01$."
 Llence the viscous tine is where ay.=oaf/0.]1 and # is the radius of the cise in areseconds., Hence the viscous time is where $\alpha_{0.1} = \alpha/0.1$ and $R$ is the radius of the disc in arcseconds.
 With these [fiducial numbers. the viscous timescale coincides with the age of the voung massive stars in the GC (e.g...22).," With these fiducial numbers, the viscous timescale coincides with the age of the young massive stars in the GC \citep[e.g.,][]{KrabbeEtal95,PaumardEtal06}."
 As o is highly uncertain. we shall then consider the two opposite cases.," As $\alpha$ is highly uncertain, we shall then consider the two opposite cases."
 Loa=0.1. Foe a few million vears (zzLO? vears at 0.27). and we expect that gas would have mainly accreted onto bby now.," If $\alpha = 0.1$, $t_{\rm visc} \simlt$ a few million years $\approx 10^6$ years at 0.2”), and we expect that gas would have mainly accreted onto by now."
 This accretion rate would be a significant [fraction of the Ededington aceretion rate., This accretion rate would be a significant fraction of the Eddington accretion rate.
 Standard: disc. accretion would then generate as much as 2107 erg Mosa of radiative energv. where ου is the gas mass accreted. by iin unitsof 107 AL.," Standard disc accretion would then generate as much as $2\times 10^{56}$ erg $M_{\rm acc, 3}$ of radiative energy, where $M_{\rm
acc, 3}$ is the gas mass accreted by in unitsof $10^3 \msun$ ."
 A similar amount of enerey could have been released: as energetic outllows., A similar amount of energy could have been released as energetic outflows.
 There is currently no, There is currently no
((2000) have reported. the discovery οἱ oscillations in photometric data from à UMa observed. with theà star nrocameraapt on theà WIR.2|. satellite.,"(2000) have reported the discovery of oscillations in photometric data from $\alpha\,$ UMa observed with the star camera on the WIRE satellite."
.Abre PPhey uwinterpretο theB observed oscillations as racial modes. and cautiously suggest that the modes may be excited by the mechanism similar to that responsible for solar oscillations.," They interpret the observed oscillations as radial modes, and cautiously suggest that the modes may be excited by the mechanism similar to that responsible for solar oscillations."
 The star. however. is very dillerent. from the Sun.," The star, however, is very different from the Sun."
 Is spectral type is KO 1., Its spectral type is $\;$ III.
 Models of the star’s internal structure. and. its pulsation properties. suggest that the star is a red giant.," Models of the star's internal structure, and its pulsation properties, suggest that the star is a red giant."
 Thus. even if the oscillations are stochastically excited. by. turbulence in the outer convective zone. as they are in the Sun. some important. cdillerencessee between the oscillation. properties. of a UMa and the Sun should be expected.," Thus, even if the oscillations are stochastically excited by turbulence in the outer convective zone, as they are in the Sun, some important differences between the oscillation properties of $\alpha\,$ UMa and the Sun should be expected."
 Variability is à common feature of red. giants., Variability is a common feature of red giants.
 “Phere are strong observationally based arguments that. at least in," There are strong observationally based arguments that, at least in"
and pop I galaxies of type A can at certain ages venture blueward of this limit.,and pop I galaxies of type A can at certain ages venture blueward of this limit.
" In the case of our pop II model (red line), this happens because of a short-lived evolutionary phase at ages of &10—20 Myr, where the ionizing flux has dropped sufficiently to make the nebular contribution irrelevant, yet the UV continuum of the stars remains very blue."," In the case of our pop II model (red line), this happens because of a short-lived evolutionary phase at ages of $\approx 10-20$ Myr, where the ionizing flux has dropped sufficiently to make the nebular contribution irrelevant, yet the UV continuum of the stars remains very blue."
" In the case of our pop I model (Z—0.020; black solid line), this happens because very young 5 populations with high metallicities also display (Svery Myr)weak [ΟΠΠ emission lines (e.g.Pana-gia2003;Nagaoetal. 2006)."," In the case of our pop I model $Z=0.020$; black solid line), this happens because very young $\lesssim 5$ Myr) populations with high metallicities also display very weak [OIII] emission lines \citep[e.g.][]{Panagia,Nagao et al. b}."
" In an instantaneous burst scenario, the timescales over which which Pop I/II systems display these very blue colours are shorter than the corresponding timescales for pop III systems, which may limit the risk of pop I/II interlopers in surveys for pop III galaxies."," In an instantaneous burst scenario, the timescales over which which Pop I/II systems display these very blue colours are shorter than the corresponding timescales for pop III systems, which may limit the risk of pop I/II interlopers in surveys for pop III galaxies."
" On the other hand, the typical star formation histories may be different in the two cases, and adopting a more extended star formation history for pop I/II could make these objects display pop III-like colours for a longer period of time."," On the other hand, the typical star formation histories may be different in the two cases, and adopting a more extended star formation history for pop I/II could make these objects display pop III-like colours for a longer period of time."
" Dust extinction, which is likely to be associated with star formation in high-metallicity environments, may of course render the colours of young pop I/II systems redder and possibly prevent metal-enriched galaxies from entering selected pop III samples."," Dust extinction, which is likely to be associated with star formation in high-metallicity environments, may of course render the colours of young pop I/II systems redder and possibly prevent metal-enriched galaxies from entering colour-selected pop III samples."
" While the details of the IMF determines the longevity of pop III galaxies, the m277—1444 colours predicted for type A, pop III objects of age <10 Myr are not sensitive to the exact pop III IMF, as illustrated by the largely overlapping cyan, blue and green lines in Fig. 6.."," While the details of the IMF determines the longevity of pop III galaxies, the $m_{277}-m_{444}$ colours predicted for type A, pop III objects of age $\lesssim 10$ Myr are not sensitive to the exact pop III IMF, as illustrated by the largely overlapping cyan, blue and green lines in Fig. \ref{typeA_Inoue_crit}."
" In Fig. 7,,"," In Fig. \ref{typeA_MIRIcrit},"
" we demonstrate that that cleaner selection criteria for type A, pop III galaxy candidates at z+8 can in principle be derived by combining JWST/NIRCam data in the F444W filter with observations in two JWST/MIRI filters: F560W and F770W. The lack of [O1II]A5007 in F444W and the presence of Ha in F560W makes young pop III galaxies redder in m444—ms560 (Fig."," we demonstrate that that cleaner selection criteria for type A, pop III galaxy candidates at $z\approx8$ can in principle be derived by combining JWST/NIRCam data in the F444W filter with observations in two JWST/MIRI filters: F560W and F770W. The lack of $\lambda$ 5007 in F444W and the presence of $\alpha$ in F560W makes young pop III galaxies redder in $m_{444}-m_{560}$ (Fig."
" 7aa), yet bluer in mgeo—M770 (Fig."," \ref{typeA_MIRIcrit}a a), yet bluer in $m_{560}-m_{770}$ (Fig."
 7bb) thanother types of galaxies., \ref{typeA_MIRIcrit}b b) thanother types of galaxies.
" Due to the [OIII] degeneracy between pop III and pop I galaxies, both can attain similar m444—meo colours."," Due to the [OIII] degeneracy between pop III and pop I galaxies, both can attain similar $m_{444}-m_{560}$ colours."
" However, since there is no similar degeneracy in m560—™770, pop III objects appear in a corner of the mago—m70 Vs. m444—™560 diagram (Fig. 8))"," However, since there is no similar degeneracy in $m_{560}-m_{770}$, pop III objects appear in a corner of the $m_{560}-m_{770}$ vs. $m_{444}-m_{560}$ diagram (Fig. \ref{typeA_MIRIcolcol}) )"
 where no other galaxies should appear., where no other galaxies should appear.
" Pop III objects remain in the upper left corner (grey region) up tozz15 Myr after formation in the case of an instantaneous burst, and longer in the case of more extended star formation histories (in principle up to ~100 Myr in the case of a constant SFR, although such high ages are unlikely to apply, as discussed in Sect. 3))."," Pop III objects remain in the upper left corner (grey region) up to$\approx 15$ Myr after formation in the case of an instantaneous burst, and longer in the case of more extended star formation histories (in principle up to $\sim 100$ Myr in the case of a constant SFR, although such high ages are unlikely to apply, as discussed in Sect. \ref{scenarios}) )."
 'The colours of pop III galaxies with different IMFs are almost identical at young ages., The colours of pop III galaxies with different IMFs are almost identical at young ages.
" We have verified that these criteria hold for objects up to an age of z15 Myr as long as the Lyman-continuum escape fraction is feseS0.5, irrespective of the pop III IMF (but higher fes; are allowed in the case of younger ages or more top-heavy IMFs)."," We have verified that these criteria hold for objects up to an age of $\approx 15$ Myr as long as the Lyman-continuum escape fraction is $f_\mathrm{esc}\lesssim 0.5$, irrespective of the pop III IMF (but higher $f_\mathrm{esc}$ are allowed in the case of younger ages or more top-heavy IMFs)."
 Fig., Fig.
 8 also illustrates the effect of dust reddening on these colour criteria., \ref{typeA_MIRIcolcol} also illustrates the effect of dust reddening on these colour criteria.
" Even though substantial dust reddening (if at all possible in pop III galaxies) could conceal pop III galaxies by making them appear as newborn, high-metallicity objects, there is no risk of false positives - i.e. no other type of galaxies in our model grid can spuriously appear as pop III objects at this redshift."," Even though substantial dust reddening (if at all possible in pop III galaxies) could conceal pop III galaxies by making them appear as newborn, high-metallicity objects, there is no risk of false positives - i.e. no other type of galaxies in our model grid can spuriously appear as pop III objects at this redshift."
" One major obstacle forany colour criteria applied to SEDs dominated by nebular emission, is that they evolve very quickly as a function of redshift."," One major obstacle for colour criteria applied to SEDs dominated by nebular emission, is that they evolve very quickly as a function of redshift."
 In Fig., In Fig.
" 9 we plot the redshift dependence of the moz;—m444 and fhseo—m79 colours for young (1 Myr old) pop III, II and I galaxies of type A. As seen, both colours evolve dramatically as a function of redshift."," \ref{typeA_zevol} we plot the redshift dependence of the $m_{277}-m_{444}$ and $m_{560}-m_{770}$ colours for young (1 Myr old) pop III, II and I galaxies of type A. As seen, both colours evolve dramatically as a function of redshift."
" By contrast, the colours of equally young type C galaxies (i.e. with SEDs dominated by direct starlight; dashed lines) hardly evolve at all."," By contrast, the colours of equally young type C galaxies (i.e. with SEDs dominated by direct starlight; dashed lines) hardly evolve at all."
" Even though the moz;—m444«0 criterion picks up on pop III galaxies at z7 6.7—8.5, it also selects young pop I and II galaxies of both type A and C at z«5 and at z>6.5."," Even though the $m_{277}-m_{444}<0$ criterion picks up on pop III galaxies at $z\approx 6.7$ –8.5, it also selects young pop I and II galaxies of both type A and C at $z<5$ and at $z\geq 6.5$."
" The risk of confusion can be limited, but not completely removed, by adding photometry in other NIRCam filters and applying drop-out criteria."," The risk of confusion can be limited, but not completely removed, by adding photometry in other NIRCam filters and applying drop-out criteria."
" For instance, additional observations in the F090W and F115W filters should effectively limit the redshift range of the sample to ze 6.5-9.0 since objects in this redshift range are expected to appear as drop-outs in F090W, but not in F115W, due to IGM absorption shortward of Lya."," For instance, additional observations in the F090W and F115W filters should effectively limit the redshift range of the sample to $z\approx 6.5$ –9.0 since objects in this redshift range are expected to appear as drop-outs in F090W, but not in F115W, due to IGM absorption shortward of $\alpha$."
 A MIRI-based pop III criterion like msg9—m779<—1.3 (shaded region in Fig. 7)), A MIRI-based pop III criterion like $m_{560}-m_{770}\leq -1.3$ (shaded region in Fig. \ref{typeA_MIRIcrit}) )
 is not likely to generate nearly as many interlopers., is not likely to generate nearly as many interlopers.
 T'his colour is useful as a pop III diagnostic over the limited range z~ 7-8 (shaded region in Fig. 9)), This colour is useful as a pop III diagnostic over the limited range $z\approx 7$ –8 (shaded region in Fig. \ref{typeA_zevol}) )
" only, but within that range, pop III galaxies produce bluer colours than all other kinds of Z« 0.020, z>1 galaxies in our model grid."," only, but within that range, pop III galaxies produce bluer colours than other kinds of $Z\leq 0.020$ , $z>1$ galaxies in our model grid."
" Even though we have here only plotted the colours of very young galaxies (1 Myr old), this holds for all interloper ages (up to the"," Even though we have here only plotted the colours of very young galaxies (1 Myr old), this holds for all interloper ages (up to the"
aceretion and feedback processes during galaxy mergers. have shown that black hole growth ts determined by the gas supply and terminates abruptly when significant gas 1s expelled as a result of the coupling of radiative feedback energy from black hole accretion to the surrounding gas.,"accretion and feedback processes during galaxy mergers, have shown that black hole growth is determined by the gas supply and terminates abruptly when significant gas is expelled as a result of the coupling of radiative feedback energy from black hole accretion to the surrounding gas."
 This determines a timescale for a strong accretion phase - re. a quasar phase for the black hole - of ~10°yr (for a Milky Way mass system).," This determines a timescale for a strong accretion phase - i.e. a quasar phase for the black hole - of $\sim 10^{8}\,{\rm yr}$ (for a Milky Way mass system)."
 This timescale is approximately equal to the gas inflow timescale. set by the time when strong torques are present in the central region. and is substantially shorter than the merger timescale of ~2«10?yr (Hernquist1989;Barnes&Hern-quist1991. 1996).," This timescale is approximately equal to the gas inflow timescale, set by the time when strong torques are present in the central region, and is substantially shorter than the merger timescale of $\sim2\times10^{9}\,{\rm yr}$ \citep{Hernquist89,BH91,BH96}."
. The timescale for associated starbursts can be similarly understood as dependent on large densities and thus related to the timescale of gas inflow (Mihos&Hernquist1994. 1996).," The timescale for associated starbursts can be similarly understood as dependent on large densities and thus related to the timescale of gas inflow \citep{MH94,MH96}."
. In order to compare the quasar lifetime predicted in these simulations with the observational estimates of the quasar lifetime. mostly inferred from visible wavelength observations. Is necessary to take into account additional physical mechanisms.," In order to compare the quasar lifetime predicted in these simulations with the observational estimates of the quasar lifetime, mostly inferred from visible wavelength observations, is necessary to take into account additional physical mechanisms."
 Previous theoretical models of supermassive black hole and quasar evolution assume that the observed quasar lifetime and the lifetime of the accretion phase are the same., Previous theoretical models of supermassive black hole and quasar evolution assume that the observed quasar lifetime and the lifetime of the accretion phase are the same.
 However. the aceretion rate is tied to the surrounding density of gas. and quasar activity occurs when material is efficiently channeled to the center of the galaxy onto the supermassive black hole.," However, the accretion rate is tied to the surrounding density of gas, and quasar activity occurs when material is efficiently channeled to the center of the galaxy onto the supermassive black hole."
 Therefore. the time(s) of greatest accretion activity are likely to also have the largest obscuring column density. especially if the line-of-sight column to the black hole 1s dominated by the density in the central. most dense regions of the galaxy.," Therefore, the time(s) of greatest accretion activity are likely to also have the largest obscuring column density, especially if the line-of-sight column to the black hole is dominated by the density in the central, most dense regions of the galaxy."
 The corresponding attenuation may render the quasar unobservable in some bands or lower its observed luminosity below the quasar threshold. giving a shorter observed lifetime than intrinsic lifetime.," The corresponding attenuation may render the quasar unobservable in some bands or lower its observed luminosity below the quasar threshold, giving a shorter observed lifetime than intrinsic lifetime."
 Modeling this effect is necessary to estimate the intrinsic quasar lifetimes from observations. as well as for using theoretically motivated accretion models to predict the quasar luminosity funetion and space density of present-day supermassive black holes.," Modeling this effect is necessary to estimate the intrinsic quasar lifetimes from observations, as well as for using theoretically motivated accretion models to predict the quasar luminosity function and space density of present-day supermassive black holes."
 Further. this effect can account for the presence of an obscured population of quasars which are missed by optical. UV. or soft X-ray surveys but significantly contribute to the cosmic X-ray background (Brandt&Hasinger2005.andreferencestherein)..," Further, this effect can account for the presence of an obscured population of quasars which are missed by optical, UV, or soft X-ray surveys but significantly contribute to the cosmic X-ray background \citep[and references therein]{BH05}."
 Similarly. re-processing of quasar radiation by dust in a large obscuring column ean account for and model observations of luminous and ultra-Iuminous infrared galaxies. which show evidence of both merger activity and obscured AGN in their cores Sanders&Mirabel1996:Komossaetal. 2003).," Similarly, re-processing of quasar radiation by dust in a large obscuring column can account for and model observations of luminous and ultra-luminous infrared galaxies, which show evidence of both merger activity and obscured AGN in their cores \citep[e.g.,][]{SM96,Komossa03}."
. In this paper. we propose a model of quasar lifetimes in which quasars are ereated and fed in major mergers. but for à large part of the accretion lifetime they are heavily obscured by the large gas densities powering accretion.," In this paper, we propose a model of quasar lifetimes in which quasars are created and fed in major mergers, but for a large part of the accretion lifetime they are heavily obscured by the large gas densities powering accretion."
 Eventually. feedback from the accretion energy drives away gas. creating a brief window in which the central object is observable as a quasar. until accretion levels drop below quasar thresholds.," Eventually, feedback from the accretion energy drives away gas, creating a brief window in which the central object is observable as a quasar, until accretion levels drop below quasar thresholds."
 To test this. we determine quasar lifetimes ffrom a simulation of à major merger of gas-rich galaxies. calculating the effects of obscuration.," To test this, we determine quasar lifetimes from a simulation of a major merger of gas-rich galaxies, calculating the effects of obscuration."
 Using a model for black hole aceretion and associated feedback developed by Springel.DiMatteo.&Hernquist(20058):Springel.&Hernquist (2005).. the simulation allows us to determine unambiguously the intrinsic lifetime of a strong aceretion phase in the merger.," Using a model for black hole accretion and associated feedback developed by \citet{SDH05b, DSH05}, the simulation allows us to determine unambiguously the intrinsic lifetime of a strong accretion phase in the merger."
 By simultaneously tracing the column densities along different lines-of-sight from the simulation supermassive black hole. here we calculate the observed Juminosity in visible wavelengths. and determine a lifetime over which the AGN would be observed as a quasar. with luminosity above some threshold.," By simultaneously tracing the column densities along different lines-of-sight from the simulation supermassive black hole, here we calculate the observed luminosity in visible wavelengths, and determine a lifetime over which the AGN would be observed as a quasar, with luminosity above some threshold."
 We find that there is a significant difference between the intrinsic lifetime and the observable value. and find that our predicted observed lifetimes agree well with the observational estimates. with a substantially longer intrinsic lifetime.," We find that there is a significant difference between the intrinsic lifetime and the observable value, and find that our predicted observed lifetimes agree well with the observational estimates, with a substantially longer intrinsic lifetime."
 The simulations used the GADGET-2 code. a new version of the parallel TreeSPH code GADGET (Springel.Yoshida.&White2001).," The simulations used the GADGET-2 code, a new version of the parallel TreeSPH code GADGET \citep{SYW01}."
. It uses an entropy-conserving formulation of SPH (Springel&Hernquist2002).. and includes radiative cooling. heating by à UV background. and à sub-resolution model of a two-phase structure of the dense ISM to describe star formation and supernova feedback (Springel&Hertquist2003).," It uses an entropy-conserving formulation of SPH \citep{SH02}, and includes radiative cooling, heating by a UV background, and a sub-resolution model of a two-phase structure of the dense ISM to describe star formation and supernova feedback \citep{SH03}."
. This sub-resolution model gives an effective equation of state which includes pressure feedback from supernovae heating. and allows us to stably evolve even massive pure gaseous disks (see. e.g. Robertson et al.," This sub-resolution model gives an effective equation of state which includes pressure feedback from supernovae heating, and allows us to stably evolve even massive pure gaseous disks (see, e.g. Robertson et al."
 2004)., 2004).
 The methodology of accretion. feedback. and galaxy generation is described in detail in Springel.DiMatteo.&Hernquist(2005b).," The methodology of accretion, feedback, and galaxy generation is described in detail in \citet{SDH05b}."
. In particular. in this model. supermassive black holes (BHs) are represented by “sink” particles that accrete gas from their local environment. with an accretion rate M estimated from the local eas density and sound speed using a Bondi-Hoyle-Lyttleton parameterization with an imposed upper limit equal to the Eddington rate.," In particular, in this model, supermassive black holes (BHs) are represented by “sink” particles that accrete gas from their local environment, with an accretion rate $\Mdot$ estimated from the local gas density and sound speed using a Bondi-Hoyle-Lyttleton parameterization with an imposed upper limit equal to the Eddington rate."
 The bolometric luminosity of the BH particle is then Ly.)=eMc. with e=0.1 the accretion efficiency.," The bolometric luminosity of the BH particle is then $\Lbol=\dEdt$, with $\epsilon=0.1$ the accretion efficiency."
 We further assume that a small fraction (5%) of L&4 couples dynamically to the surrounding gas. and this feedback is injected as thermal energy.," We further assume that a small fraction $5\%$ ) of $\Lbol$ couples dynamically to the surrounding gas, and this feedback is injected as thermal energy."
 This fraction is a free parameter. determined in DiMatteo.Springel.&Hernquist(2005). by fitting to the Mgy- a relation.," This fraction is a free parameter, determined in \citet{DSH05} by fitting to the $M_{\rm BH}$ $\sigma$ relation."
 We do not attempt to resolve the small-scale accretion dynamics near the black hole. but instead assume that the time-averaged accretion can be estimated from the gas properties on the scale of our spatial resolution (=100 ppc).," We do not attempt to resolve the small-scale accretion dynamics near the black hole, but instead assume that the time-averaged accretion can be estimated from the gas properties on the scale of our spatial resolution $\lesssim100$ pc)."
 We generate two stable. isolated disk galaxies. each with an extended dark matter halo with a Hernquist(1990). profile. an exponential gas disk. and a bulge.," We generate two stable, isolated disk galaxies, each with an extended dark matter halo with a \citet{Hernquist90} profile, an exponential gas disk, and a bulge."
" Our simulation is one of the series described in detail in Springel.DiMatteo.&Hern-quist (20052). with virial velocity Vj,2I60kms''. a fiducial choice with a rotation curve and mass similar to the Milky Way."," Our simulation is one of the series described in detail in \citet{SDH05a}, with virial velocity $V_{\rm vir}=160\,{\rm km\,s^{-1}}$, a fiducial choice with a rotation curve and mass similar to the Milky Way."
 We begin our simulation with pure gaseous disks. which may better correspond to the high-redshift galaxies in which most quasars are observed.," We begin our simulation with pure gaseous disks, which may better correspond to the high-redshift galaxies in which most quasars are observed."
 Each galaxy is initially composed of 168000 dark matter halo particles. 8000 bulge particles. 24000 gaseous disk particles. and one BH particle. with a small initial seed mass of 10?M...," Each galaxy is initially composed of 168000 dark matter halo particles, 8000 bulge particles, 24000 gaseous disk particles, and one BH particle, with a small initial seed mass of $10^{5}M_{\sun}$."
 Given these choices. the dark matter. gas. and star particles are all of roughly equal mass. and central cusps in the dark matter and bulge profiles (Hernquist1990) are reasonably well resolved.," Given these choices, the dark matter, gas, and star particles are all of roughly equal mass, and central cusps in the dark matter and bulge profiles \citep{Hernquist90} are reasonably well resolved."
 The galaxies are then set to collide in a pure prograde encounter with zero orbital energy and a pericenter separation of 7.1kpe.," The galaxies are then set to collide in a pure prograde encounter with zero orbital energy and a pericenter separation of $7.1\,{\rm kpc}$."
 We calculate the column density between a black hole and a hypothetical observer from the simulation outputs spaced every 10 Myr before and after the merger and every 5 Myr during the merger (1.14 - 1.71 Gyr)., We calculate the column density between a black hole and a hypothetical observer from the simulation outputs spaced every 10 Myr before and after the merger and every 5 Myr during the merger (1.14 - 1.71 Gyr).
 We first generate ~1000 radial lines-of-sight (rays). each withits origin at the black hole particle location and with directions uniformly spaced in solid angle dcos@do.," We first generate $\sim1000$ radial lines-of-sight (rays), each withits origin at the black hole particle location and with directions uniformly spaced in solid angle $d\cos{\theta}\,d\phi$ ."
 For each ray. we then begin at the origin. calculate and record the local gas properties using," For each ray, we then begin at the origin, calculate and record the local gas properties using"
could become an issue if the absorbing region is actually far enough from the ΑΝ core to avoid severe line broacening by radiation dampiug.,could become an issue if the absorbing region is actually far enough from the AGN core to avoid severe line broadening by radiation damping.
 If saturation aud radiatiou damping were not a problem. the uoidetection of H65a could be ied. along with the measured EM. to place some coustraiuts ou the plysical couditious of the ree-[ree absorbing gas.," If saturation and radiation damping were not a problem, the non-detection of $\alpha$ could be used, along with the measured EM, to place some constraints on the physical conditions of the free-free absorbing gas."
 For completeness. we will show those constraints.," For completeness, we will show those constraints."
 The coustraints would be oL interest if the absorbing region is actually farther from the continuum radiation source than we yelieve or if for some reason tlie radiation damping is not as severe as we believe., The constraints would be of interest if the absorbing region is actually farther from the continuum radiation source than we believe or if for some reason the radiation damping is not as severe as we believe.
 Usiug the measured EM aud the assumed temperature. the optical depth of maser-like recombination ine emission that would be expected from the absorbing gas ean be calculated lor a rauge of values oL deusity aud line width.," Using the measured EM and the assumed temperature, the optical depth of maser-like recombination line emission that would be expected from the absorbing gas can be calculated for a range of values of density and line width."
 The lower panel of Figure L shows the results of such caleulations [or an emission measure Chat is midway between the values obtaiued from the continuum absorption., The lower panel of Figure \ref{recomb10000} shows the results of such calculations for an emission measure that is midway between the values obtained from the continuum absorption.
 The caleulatious are based ou departure coelficieuts very similar to those of aud are for the case when external radiation does not allect the level populations or line width., The calculations are based on departure coefficients very similar to those of \citet{SB79} and are for the case when external radiation does not affect the level populations or line width.
 Optical depths are plotted [or several different line widths., Optical depths are plotted for several different line widths.
 The figure also shows the upper LE.nit to the liue-to-contiuuum ratio of 0.15 that was derived above from the data., The figure also shows the upper limit to the line-to-continuum ratio of 0.15 that was derived above from the data.
 Iu the absence of saturation or radiation dampine. Figure L shows that the recombination line non-detectiou would iiply that the absorbing [n]gas must have au appropriate combination of high density or high line width. presumably due to turbulence or other dynamic ellects.," In the absence of saturation or radiation damping, Figure \ref{recomb10000} shows that the recombination line non-detection would imply that the absorbing gas must have an appropriate combination of high density or high line width, presumably due to turbulence or other dynamic effects."
 Or it could be at a velocity outside he observed window., Or it could be at a velocity outside the observed window.
 Note that the thermal line width for the assumed temperature is lower than he sinallest. velocity width shown., Note that the thermal line width for the assumed temperature is lower than the smallest velocity width shown.
" To put the velocity raugee and line widths observed into some perspective. consider that the orbital velocity at about 3.3 pe [rom a. £x10A. black hole is about 700ὃν, or about half ol the total observed velocity rauge."," To put the velocity range and line widths observed into some perspective, consider that the orbital velocity at about 3.3 pc from a $4\times10^8 M_\sun$ black hole is about 700, or about half of the total observed velocity range."
 That mass is clearly uucertain but is what was estimated by Wu&Han(2001) using the correlation between black hole mass and bulge velocity dispersion., That mass is clearly uncertain but is what was estimated by \citet{WH01} using the correlation between black hole mass and bulge velocity dispersion.
 The 3.3 pe distauce is just the observed offset of 2.5 pe deprojected aloug the clisk for the geometry described earlier., The 3.3 pc distance is just the observed offset of 2.5 pc deprojected along the disk for the geometry described earlier.
 That orbital velocity shows that it is possible for gravitationally bound material iu the region of interest to be moving [ast enough to produce line widths high enough to explain our non-detection., That orbital velocity shows that it is possible for gravitationally bound material in the region of interest to be moving fast enough to produce line widths high enough to explain our non-detection.
 But it also shows that. ifthe absorbing material is moving with the disk. it could be observable.," But it also shows that, if the absorbing material is moving with the disk, it could be observable."
 This would depeud ou the reasonable assumptions that the disk is moving trausversely where it crosses the sight lines to the far-sice jet. and that it has au internal velocity spread tliat is much sinaller than the orbital velocity.," This would depend on the reasonable assumptions that the disk is moving transversely where it crosses the sight lines to the far-side jet, and that it has an internal velocity spread that is much smaller than the orbital velocity."
 The Broad Line Region (BLR) line widths in NGC 1275 are of the order of our total frequency span (see.forexample.Nelson&Whittle1999).," The Broad Line Region (BLR) line widths in NGC 1275 are of the order of our total frequency span \citep[see, for example,][]{N99}."
. A recombination liue that wide would both not be sufficiently intense to observe. according to Figure L. and would be washed out by our primary data reduction technique.," A recombination line that wide would both not be sufficiently intense to observe, according to Figure \ref{recomb10000}, and would be washed out by our primary data reduction technique."
 But typical BLR sizes. determiued from reverberation mapping. are much stnaller than the region probed by the sight lines through tle free-DIree absorbing material.," But typical BLR sizes, determined from reverberation mapping, are much smaller than the region probed by the sight lines through the free-free absorbing material."
 Therefore the BLR line widths would ouly be a problem if the high velocities are maintained well bevond the region that emits the optical lines., Therefore the BLR line widths would only be a problem if the high velocities are maintained well beyond the region that emits the optical lines.
"have redshifts provided by ?,, ?,, ?,, ? and the NASA Extragalactic Database.","have redshifts provided by \citet{fal99}, , \citet{huc99}, , \citet{weg03}, \citet{jon09} and the NASA Extragalactic Database."
" For 170 galaxies we have redshift independent distances from Extragalactic Distance Database (?),, while for the remaining galaxies we use a Hubble constant of 73kms!Mpc""! (?).."," For 170 galaxies we have redshift independent distances from Extragalactic Distance Database \citep{tul09}, while for the remaining galaxies we use a Hubble constant of $73~{\rm km~s^{-1}~Mpc^{-1}}$ \citep{spe07}."
" The luminosity distances of the galaxies span from 0.7 to 10?Mpc Our principal radio imaging is the 1.4GHz NRAO VLA Sky Survey (?),, which has an angular resolution of 45"" and an RMS noise of ~0.45mJy."," The luminosity distances of the galaxies span from $0.7$ to $10^2~{\rm Mpc}$ Our principal radio imaging is the $1.4~{\rm GHz}$ NRAO VLA Sky Survey \citep{con98}, which has an angular resolution of $45^{\prime\prime}$ and an RMS noise of $\simeq 0.45~{\rm mJy}$."
" As the NVSS can underestimate the flux densities of powerful and extended radio sources, we also measure flux densities using lower resolution (~12’) imaging from the 300-ft Green Bank and the 64-m Parkes radio telescopes (???,Calabrettainprep.).."," As the NVSS can underestimate the flux densities of powerful and extended radio sources, we also measure flux densities using lower resolution $\sim 12^\prime$ ) imaging from the 300-ft Green Bank and the 64-m Parkes radio telescopes \citep[][Calabretta in prep.]{con85,con86,bar01}."
 We employed the following method to measure flux densities for each galaxy., We employed the following method to measure flux densities for each galaxy.
 At the 2MASS position of each galaxy we measured flux density per beam directly from the NVSS , At the 2MASS position of each galaxy we measured a flux density per beam directly from the NVSS images.
"For thosea galaxies with a flux density per beam greater than images.2mJy, we searched for the nearest counterpart in the NVSS catalog within 2’ and used the NVSS catalog deconvolved flux density."," For those galaxies with a flux density per beam greater than $2~{\rm mJy}$, we searched for the nearest counterpart in the NVSS catalog within $2^{\prime}$ and used the NVSS catalog deconvolved flux density."
" To account for powerful extended sources, we measured the flux density per beam using Green Bank 300-ft and Parkes single dish data, and utilized this flux density measurement if it was greater than 0.6Jy."," To account for powerful extended sources, we measured the flux density per beam using Green Bank 300-ft and Parkes single dish data, and utilized this flux density measurement if it was greater than $0.6~{\rm Jy}$."
" If the flux density per beam measured with single dish imagery was greater than 5Jy, we utilized the integrated flux density rather than the flux density per beam."," If the flux density per beam measured with single dish imagery was greater than $5~{\rm Jy}$, we utilized the integrated flux density rather than the flux density per beam."
" Our flux density measurements are imperfect, but our principal conclusions remain unchanged unless gross systematic errors are present."," Our flux density measurements are imperfect, but our principal conclusions remain unchanged unless gross systematic errors are present."
" To verify that our data was free of such errors, we visually inspected the radio imagery and compared our flux densities with those from the literature."," To verify that our data was free of such errors, we visually inspected the radio imagery and compared our flux densities with those from the literature."
 Visual inspection also revealed that few galaxies in our sample are extended ? class II radio sources., Visual inspection also revealed that few galaxies in our sample are extended \citet{far74} class II radio sources.
" The flux densities of 4 galaxies neighboring bright radio sources could not be reliably measured, and excluding these galaxies reduced our final to 396 "," The flux densities of 4 galaxies neighboring bright radio sources could not be reliably measured, and excluding these galaxies reduced our final sample to 396 early-type galaxies."
The properties of the sample samplegalaxies are early-typesummarized in galaxies.Table 1.., The properties of the sample galaxies are summarized in Table \ref{table:summary}.
 The 1.4GHz flux densities of the 396 K<9 early-type galaxies are plotted as a function of K-band absolute magnitude in Figure 1.., The $1.4~{\rm GHz}$ flux densities of the 396 $K<9$ early-type galaxies are plotted as a function of $K$ -band absolute magnitude in Figure \ref{fig:fluxes}.
" For the lowest mass galaxies, with My—22.5, only a small fraction have counterparts in the NVSS catalogs and most galaxies have flux densities consistent with zero plus a random measurement error."," For the lowest mass galaxies, with $M_K\sim -22.5$, only a small fraction have counterparts in the NVSS catalogs and most galaxies have flux densities consistent with zero plus a random measurement error."
" In contrast, all 59 of the My«—25.5 galaxies (corresponding to stellar masses of >2.5early-typex10!! M.) have measured flux densities greater than zero and just two of the 61 —25.0«My—25.5 have measured flux densities below zero."," In contrast, all 59 of the $M_K<-25.5$ early-type galaxies (corresponding to stellar masses of $\gtrsim 2.5\times 10^{11}~M_\odot$ ) have measured flux densities greater than zero and just two of the 61 $-25.0<M_K<-25.5$ early-type galaxies have measured flux densities below zero."
" As early-typerandom galaxieserrors will result in measured flux densities scattered above and below the true flux densities, it is highly likely that all My«—25.5 early-type galaxies in our sample are radio continuum sources, with flux densities of >1mJy."," As random errors will result in measured flux densities scattered above and below the true flux densities, it is highly likely that all $M_K<-25.5$ early-type galaxies in our sample are radio continuum sources, with flux densities of $\gtrsim 1~{\rm mJy}$."
 In Figure 2 we present the radio power of early-type galaxies as a function of K-band absolute magnitude., In Figure \ref{fig:lums} we present the radio power of early-type galaxies as a function of $K$ -band absolute magnitude.
" As noted by others (e.g.??), the radio powers of early-type galaxies are a strong function of absolute magnitude."," As noted by others \citep[e.g.,][]{fab89,sad89}, the radio powers of early-type galaxies are a strong function of absolute magnitude."
" Although our K29 magnitude limit excludes some extremely radio sources Cygnus A), the radio power of powerfulMy«—24 galaxies at (e.g.,fixed K-band magnitude (or stellar mass) still spans 4 orders of magnitude."," Although our $K=9$ magnitude limit excludes some extremely powerful radio sources (e.g., Cygnus A), the radio power of $M_K<-24$ galaxies at fixed $K$ -band magnitude (or stellar mass) still spans 4 orders of magnitude."
 The most massive galaxies can have 1.4GHz powers as largeas thatof M 87 (~7x10W Hz!) or less than the Milky Way Hz»., The most massive galaxies can have $1.4~{\rm GHz}$ powers as largeas thatof M 87 $\simeq 7\times 10^{24}~{\rm W~Hz^{-1}}$ ) or less than the Milky Way $\simeq 4\times 10^{21}~{\rm W~Hz^{-1}}$ ).
"to compute the non-linear matter power spectrum: where Phalofit,s9) is the power spectrum from applying on the no-wiggle power spectrum and P,,;(k) is the non-linear power spectrum.","to compute the non-linear matter power spectrum: where $P_{\emph{halofit},nw}(k)$ is the power spectrum from applying on the no-wiggle power spectrum and $P_{nl}(k)$ is the non-linear power spectrum."
 We compute the theoretical two-point correlation function by Fourier transforming the non-linear power spectrum., We compute the theoretical two-point correlation function by Fourier transforming the non-linear power spectrum.
" We show an example of the effect of applying dewiggle and to the correlation function in yreffig:dw;, alofit.", We show an example of the effect of applying dewiggle and to the correlation function in \\ref{fig:dw_halofit}.
".Clearly, thedampingofthebaryonacousticoscillation(dháf2)peakisaccwnabely dtsty addi dlivatecorrelation funct"," Clearly, the damping of the baryon acoustic oscillation (BAO) peak is accurately described by the dewiggled linear correlation function."
 , Additional nonlinear effects are only important on very small scales.
"The parameter set we use to compute the theoretical correlation function is {Dv(2),Qmh?,Ah?,ns,k,}, where Qn and Ων are the density fractions of matter and baryons, Τις is the power law index of the primordial matter power spectrum, h is the dimensionless Hubble constant (Ho=1005 km s-'Mpe~*), and Dy(2) is defined by where H(z) and Da(z) are the Hubble parameter and the angular diameter distance at the redshift, z."," The parameter set we use to compute the theoretical correlation function is $\{D_V(z), \Omega_mh^2, \Omega_bh^2, n_s, k_\star\}$, where $\Omega_m$ and $\Omega_b$ are the density fractions of matter and baryons, $n_s$ is the power law index of the primordial matter power spectrum, $h$ is the dimensionless Hubble constant $H_0=100h$ km $^{-1}$ $^{-1}$ ), and $D_V(z)$ is defined by where $H(z)$ and $D_A(z)$ are the Hubble parameter and the angular diameter distance at the redshift, $z$."
 We set h=0.7 while calculating the non-linear power spectra., We set $h=0.7$ while calculating the non-linear power spectra.
" Assuming the true model of our universe is not far from a ACDM model, the dark energy and curvature dependence are absorbed by the effective distance, (2)."," Assuming the true model of our universe is not far from a $\Lambda$ CDM model, the dark energy and curvature dependence are absorbed by the effective distance, $D_V(z)$."
 Thus we are able to extract constraints from data without assumingDv a dark energy model and cosmic curvature., Thus we are able to extract constraints from data without assuming a dark energy model and cosmic curvature.
" We use the mock catalogs from the LasDamas (McBride et al.,"," We use the mock catalogs from the LasDamas (McBride et al.,"
 in preparation) to estimate the covariance matrix of the observed correlation function., in preparation) to estimate the covariance matrix of the observed correlation function.
 LasDamas provides mock catalogs matching SDSS main galaxy and LRG samples., LasDamas provides mock catalogs matching SDSS main galaxy and LRG samples.
" We use the LRG mock catalogs from the LasDamas gamma release with the same cuts as the SDSS LRGfull sample, —23.2<M,<—21.2 and 0.16«z0.44."," We use the LRG mock catalogs from the LasDamas gamma release with the same cuts as the SDSS LRGfull sample, $-23.2<M_g<-21.2$ and $0.16<z<0.44$."
 We have diluted the mock catalogs to match the radial selection function of the observed data by randomly selecting the mock galaxies according to the number density of the data sample., We have diluted the mock catalogs to match the radial selection function of the observed data by randomly selecting the mock galaxies according to the number density of the data sample.
" We calculate the spherical-averaged correlation functions of the mock catalogs and construct the covariance matrix as basa ge mean of the m*” bin of the mock correlation functions, and €*, is the value of mf"" bin of the k*” mock correlation function."," We calculate the spherical-averaged correlation functions of the mock catalogs and construct the covariance matrix as where $N$ isthe number of the mock catalogs, $\bar{\xi}_m$ is the mean of the $m^{th}$ bin of the mock correlation functions, and $\xi_m^k$ is the value of $m^{th}$ bin of the $k^{th}$ mock correlation function."
 We test the normality of the correlation functions from the LasDamas mock catalogs and find that they are well described by normal distributions (see refsec:normality;est))., We test the normality of the correlation functions from the LasDamas mock catalogs and find that they are well described by normal distributions (see \\ref{sec:normality_test}) ).
 The mock catalogs derived from N-body simulations require long computing times and are very limited in availability., The mock catalogs derived from N-body simulations require long computing times and are very limited in availability.
" It is interesting to investigate whether there is an easier, faster, and cheaper way to construct mock catalogs which could work as well as those derived from N-boday simulation."," It is interesting to investigate whether there is an easier, faster, and cheaper way to construct mock catalogs which could work as well as those derived from N-boday simulation."
" Towards this end, we have created 500 lognormal(LN) mock catalogs (Coles& 2004),, and computed the spherically-averaged correlation functions from these."," Towards this end, we have created 500 lognormal(LN) mock catalogs \citep{Coles:1991if,Percival:2003pi}, , and computed the spherically-averaged correlation functions from these."
 The details, The details
"To assess whether two stars a physical or are just a chance alignment of random comprisestars, the most binaryreliable method used is to check for a common systemic velocity.","To assess whether two stars comprise a physical binary or are just a chance alignment of random stars, the most reliable method used is to check for a common systemic velocity."
" However, J0130—4445B is very faint, even in the infrared, and has not been detected in any earlier epoch; hence, we do not have proper motions for the secondary nor radial velocities for either component."," However, B is very faint, even in the infrared, and has not been detected in any earlier epoch; hence, we do not have proper motions for the secondary nor radial velocities for either component."
" In the absence of kinematic information, we employed two other tests to examine whether J0130—4445AB is a physical pair: (1) the heliocentric distances of the two components and (2) the probability that the sources are a chance alignment based on the surface distribution of stars on the sky."," In the absence of kinematic information, we employed two other tests to examine whether AB is a physical pair: (1) the heliocentric distances of the two components and (2) the probability that the sources are a chance alignment based on the surface distribution of stars on the sky."
 The spectrophotometric distances to each component of J0130—4445AB were derived using the M ;/spectral type relations from ? based on the combined-light 2MASS and our relative SpeX (Table 3)). The photo," The spectrophotometric distances to each component of AB were derived using the $M_J$ /spectral type relations from \citet{Cruz2003}
 based on the combined-light 2MASS photometry and our relative SpeX photometry (Table \ref{Tab: properties}) )."
"metryderived distances are 34.5+3.2 pc for the photometryprimary and 45.8+13.6 pc for the secondary, where the errors are from the uncertainties in the NIR spectral types (see Sec. ??))."," The derived distances are $\pm$ 3.2 pc for the primary and $\pm$ 13.6 pc for the secondary, where the errors are from the uncertainties in the NIR spectral types (see Sec. \ref{Sec: NIRanalysis}) )."
 These distances are consistent with each other within their associated errors., These distances are consistent with each other within their associated errors.
" Next, we calculated the that J0130—4445AB is a random chance probabilityalignment along our line-of-sight based on its three-dimensional position in the Galaxy."," Next, we calculated the probability that AB is a random chance alignment along our line-of-sight based on its three-dimensional position in the Galaxy."
" We followed ?,, constructing a three-component Galactic model with the thin disk, thick disk, and halo, constrained with empirical stellar density profiles (??).."," We followed \citet{Dhital2010}, constructing a three-component Galactic model with the thin disk, thick disk, and halo, constrained with empirical stellar density profiles \citep{Juric2008,
 Bochanski2010}."
" The model recreates a 30'x30' region in the sky, centered around the coordinates of the given binary system, and out to heliocentric distances of 2500 pc."," The model recreates a $\arcmin\times$ $\arcmin$ region in the sky, centered around the coordinates of the given binary system, and out to heliocentric distances of 2500 pc."
" As all the simulated stars are single and non-associated, any visual binary is a random chance alignment."," As all the simulated stars are single and non-associated, any visual binary is a random chance alignment."
" In 10"" Monte Carlo realizations, we found, on average, 0.0285 chance alignments per realization on the sky, within the 3/3 angular separation of "," In $10^7$ Monte Carlo realizations, we found, on average, 0.0285 chance alignments per realization on the sky, within the $\farcs$ 3 angular separation of AB."
"J0130—4445AB. More none of these chance alignments were withinimportantly, the range of spectrophotometric distances (40414 pc) estimated for J0130—4445AB."," More importantly, none of these chance alignments were within the range of spectrophotometric distances $\pm$ 14 pc) estimated for AB."
" As such, we conservatively infer a κ107* probability of positional coincidence."," As such, we conservatively infer a $\lesssim 10^{-7}$ probability of positional coincidence."
 We therefore conclude that J0130—4445AB is a physically-bound binary and not a chance alignment of two unassociated stars., We therefore conclude that AB is a physically-bound binary and not a chance alignment of two unassociated stars.
" The NIR spectral types of J0130—4445A and J0130—4445B correspond to effective temperatures,Teg, of 2400+110 K and 14504100 K, respectively, based on the T.g/spectral type relation of ?.."," The NIR spectral types of A and B correspond to effective temperatures, of $\pm$ 110 K and $\pm$ 100 K, respectively, based on the /spectral type relation of \citet{Stephens2009}."
 Uncertainties include scatter in the relation and uncertainties in the classifications (Sec. ??))," Uncertainties include scatter in the relation and uncertainties in the classifications (Sec. \ref{Sec:
 NIRanalysis}) )"
 added in quadrature., added in quadrature.
" Figure 4 shows the ?? evolutionary models, displaying as a function of mass and age; the Burrows mass tracks and the observed are shown as dotted and dashed lines and gray boxes, rangesrespectively."," Figure \ref{Fig: evol_track}
 shows the \citet{Burrows1993, Burrows1997} evolutionary models, displaying as a function of mass and age; the Burrows mass tracks and the observed ranges are shown as dotted and dashed lines and gray boxes, respectively."
" As the masses of the two components vary quite a bit with assumed age, it is imperative to constrain the age of J0130—4445AB."," As the masses of the two components vary quite a bit with assumed age, it is imperative to constrain the age of AB."
" The absence of in 10130--4445Α indicates that it is more massive than the predicted lithium depletion boundary (LDB) mass, ~0.06 (??) for field stars of solar metallicity."," The absence of in A indicates that it is more massive than the predicted lithium depletion boundary (LDB) mass, $\sim$ 0.06 \citep{Chabrier1996, Burrows2001} for field stars of solar metallicity."
" Using the for the primary based on its spectral type (including uncertainties) and assuming a mass 20.06Mo, the evolutionary models of ? indicate an age 2250 Myr (Figure 4))."," Using the for the primary based on its spectral type (including uncertainties) and assuming a mass $\gtrsim$ 0.06, the evolutionary models of \citet{Burrows1997}
 indicate an age $\gtrsim$ 250 Myr (Figure \ref{Fig:
 evol_track}) )."
" We note that recent work by ? has suggested that episodic accretion during the pre-main sequence stages causes central temperature of a star to increase up to 1 dex, with a sharp dependence on the frequency and magnitude of the episodic accretion."," We note that recent work by \citet{Baraffe2010} has suggested that episodic accretion during the pre-main sequence stages causes central temperature of a star to increase up to 1 dex, with a sharp dependence on the frequency and magnitude of the episodic accretion."
" This serves to depleteI earlier than in non-accreting stars of the same mass and effectively reduces the inferred LDB mass and, thus, the minimum allowable age."," This serves to deplete earlier than in non-accreting stars of the same mass and effectively reduces the inferred LDB mass and, thus, the minimum allowable age."
" Here, we have not taken episodic accretion into account."," Here, we have not taken episodic accretion into account."
 While we do not have, While we do not have
refclistortion)). the errors from the sparse sampling of the integration area may also be substantial in our application of the ¢-statistic. and. as 5; approaches unity those introduced. by an inaccurate reconstruction of the cluster mass cistribution will become important as well.,"), the errors from the sparse sampling of the integration area may also be substantial in our application of the $\zeta$ -statistic, and as $\kappa_i$ approaches unity those introduced by an inaccurate reconstruction of the cluster mass distribution will become important as well."
 The discussion so far assumed. a single redshift. plane for the xickeground: sources., 	 The discussion so far assumed a single redshift plane for the background sources.
 In the case of a redshift distribution. he equations are still valid for 5;ας if the surface mass density and the shear are interpreted as quantities referring o the redshift-averaged critical surface mass density. which means s=faery and 5=(uj.54.," In the case of a redshift distribution, the equations are still valid for $\kappa_i\ll 1$, if the surface mass density and the shear are interpreted as quantities referring to the redshift-averaged critical surface mass density, which means $\kappa=\wave\,\kappa_{\infty}$ and $\gamma=\wave\,\gamma_{\infty}$."
 But in the non-incar regime. the treatment should be generalized. because hen the expectation value for observed. image cllipticitics also depends on higher moments of the redshift distribution.," But in the non-linear regime, the treatment should be generalized, because then the expectation value for observed image ellipticities also depends on higher moments of the redshift distribution."
 Seitz Schneider (1997) derived the approximation which is quite accurate for ny0.5 and generic redshift distributions., Seitz Schneider (1997) derived the approximation which is quite accurate for $\kappa_{\infty}\la 0.8$ and generic redshift distributions.
 In analogy to the calculations above. this ollers an estimate for the tangential shear in terms of the observed. image ellipticities (and the reconstructed. cluster mass distribution). and expressions equivalent to equations (13)) and (14)) can be computed.," In analogy to the calculations above, this offers an estimate for the tangential shear in terms of the observed image ellipticities (and the reconstructed cluster mass distribution), and expressions equivalent to equations \ref{zeta-sum}) ) and \ref{zeta-sig}) ) can be computed."
 In the application to the simulations we emploved this slightly generalized formalism., In the application to the simulations we employed this slightly generalized formalism.
 The racial mass profile of the cluster galaxies can be probed ov caleulating the C-statistic as a function. of the inner racius of the annulus., The radial mass profile of the cluster galaxies can be probed by calculating the $\zeta$ -statistic as a function of the inner radius of the annulus.
 The minimal inner racius for applying his method. is limited by the ability to measure reliable cllipticities for background: galaxy. images in the vicinity of the typically much brighter cluster galaxies., The minimal inner radius for applying this method is limited by the ability to measure reliable ellipticities for background galaxy images in the vicinity of the typically much brighter cluster galaxies.
 Another. heoretical complication for images located very close to cluster galaxies is that those also contribute to the surface mass density at the image positions. and this should. in xinciple be included. in the corrective factor Hi) ds well. which requires specilving a model for the galaxy mass distribution.," Another, theoretical complication for images located very close to cluster galaxies is that those also contribute to the surface mass density at the image positions, and this should in principle be included in the corrective factor $1-\kappa_i$ ) as well, which requires specifying a model for the galaxy mass distribution."
 However. this problem can be neglected in view ofthe observational limitations mentioned above. because at useful racial distances the surface mass density of the cluster galaxies should already have dropped to insignificant values.," However, this problem can be neglected in view of the observational limitations mentioned above, because at useful radial distances the surface mass density of the cluster galaxies should already have dropped to insignificant values."
 The outer radius of the annulus must not be so Large jw the referencee-term Ανο). picks up variations of 10 cluster mass distribution.," The outer radius of the annulus must not be so large that the reference-term $\overline{\kappa}(x_1,x_2)$ picks up variations of the cluster mass distribution."
 In. practice. a more serious imitation for its extent is the presence of neighbouring cluster galaxies. which must not be located within the annulus in order to keep &CGrj..rc2) as small as possible.," In practice, a more serious limitation for its extent is the presence of neighbouring cluster galaxies, which must not be located within the annulus in order to keep $\overline{\kappa}(x_1,x_2)$ as small as possible."
 As ong as these constraints are satisfied. sri.) and therefore also QGri.c9) are nearly independent of the outer radius re.," As long as these constraints are satisfied, $\overline{\kappa}(x_1,x_2)$ and therefore also $\zeta(x_1,x_2)$ are nearly independent of the outer radius $x_2$."
 ‘This means that the outer radius can be chosen individually or each cluster galaxy. ancl still the c-estimates from each of those can be combined afterwards to achieve significant results., This means that the outer radius can be chosen individually for each cluster galaxy and still the $\zeta$ -estimates from each of those can be combined afterwards to achieve significant results.
 In order to use the available information effectively. we therefore adopted the following strategy for calculating the mass estimates for a given value of the inner radius.," In order to use the available information effectively, we therefore adopted the following strategy for calculating the mass estimates for a given value of the inner radius."
 Cluster galaxies were included in the analysis. if inner-racdius circles centred on them did not intersect. the inner-racdius circle crawn around any other cluster galaxy. (," Cluster galaxies were included in the analysis, if inner-radius circles centred on them did not intersect the inner-radius circle drawn around any other cluster galaxy. ("
Vhis leads to a bias of the positions of the included cluster galaxies away from the cluster centre.),This leads to a bias of the positions of the included cluster galaxies away from the cluster centre.)
 For each of the cluster galaxies used. the outer radius was then specified. as the maximal radius possible without intersecting the zmmer--radius circle around any other cluster galaxy.," For each of the cluster galaxies used, the outer radius was then specified as the maximal radius possible without intersecting the -radius circle around any other cluster galaxy."
 With this. prescription. some background galaxy images located between cluster ealaxies have to be included in the analysis twice or several times. with reference to dillerent cluster galaxies.," With this prescription, some background galaxy images located between cluster galaxies have to be included in the analysis twice or several times, with reference to different cluster galaxies."
 Note that the method implicitly takes into account the shear effects of more than one cluster galaxy on individual background images., Note that the method implicitly takes into account the shear effects of more than one cluster galaxy on individual background images.
luminosity and halo mass.,luminosity and halo mass.
 Phe essential idea is that any scatter increases the contribution from lower mass (and less highly biased) halos. so that a measurement of large clustering amplitude limits the contribution from lower mass objects.," The essential idea is that any scatter increases the contribution from lower mass (and less highly biased) halos, so that a measurement of large clustering amplitude limits the contribution from lower mass objects."
 This constraint is as strong as direct. observational constraints at lower redshift on the amount of scatter in the various relationships mentioned above.," This constraint is as strong as direct, observational constraints at lower redshift on the amount of scatter in the various relationships mentioned above."
 We shall work throughout in the . CDM framework ancl adopt the following cosmological parameters: Qua= 0.25. OX=0.75. h=0.72 and ax=0.8.," We shall work throughout in the $\Lambda$ CDM framework and adopt the following cosmological parameters: $\Omega_{\rm mat}=0.25$ , $\Omega_\Lambda=0.75$, $h=0.72$ and $\sigma_8=0.8$."
 Where appropriate we shall comment on how our results depend on these particular choices., Where appropriate we shall comment on how our results depend on these particular choices.
 The next section outlines our formalism and applies it to the data of(2007)... while relsce:dis¢ussion interprets our results and discusses future observations.," The next section outlines our formalism and applies it to the data of, while \\ref{sec:discussion} interprets our results and discusses future observations."
 We will focus our attention on the space density and. [arge-scale. bias of QSOs. as these are the easiest to interpret. observationally measurable properties of the population.," We will focus our attention on the space density and large-scale bias of QSOs, as these are the easiest to interpret, observationally measurable properties of the population."
" Lf the mean number of QSOs in à halo of mass My is IN(M,) then their number density and large-scale bias are ∖∖⊽↓↥∢⊾↓⋅⋖⋅↙∣⊔∕∣↙∣⇀∪∕∣⊲↓⊳∖↿↓↕∢⊾↿∖≼⇍∪⊔↓∪∖⋰↓⊔⋏∙≟∃⊔⊔⊔↓∣⋈⊾↓⋅∠⇂∢⋅⊔≱∖⋠↓↿∙∖⇁∪⇂∎ ↓↕⋜↧⇂∪⊳∖↓≻∢⊾↓⋅⊔⋯⊳∖⊳∖⊲↓⊔∩⊾↓⋅∖⇁⋜↧↓⋜⋯∠⇂∣↗∕∣∟∪∕∣∃⊲↓⊳∖↿↓↕⋖⋅∣⋡⋠↓⋜↧≱∖⋜↧≱∖≱∖⋯∷↓", If the mean number of QSOs in a halo of mass $M_h$ is $N(M_h)$ then their number density and large-scale bias are where $dn_h/dM_h$ is the (comoving) number density of halos per mass interval and $b_h(M_h)$ is the bias associated with halos of that mass.
⋜⋯⋅∠⇂ ∖∖⋰∐↓↥↓↥⋜↧↓∪⊳∖∪⇂⋅∣↓⋯↿⊔↓⋜↧⊳∖⊳∖⊳↓⊲∖∪↓⋅↙∣⊔∕∣↙∣⇀∪∕∣∖∖⊽⋖⊾⊔⊳∖∢⊾↿↓↕∢⊾∐∏⊀↓⊔⋏∙≟ function from(1999)., For $dn_h/dM_h$ we use the fitting function from.
". Phe large-scale bias is slightly more problematic. as our results depend: upon massive halos being biased and cillerent fits to 6,CÀ,) from simulations have appeared in the literature."," The large-scale bias is slightly more problematic, as our results depend upon massive halos being biased and different fits to $b_h(M_h)$ from simulations have appeared in the literature."
 derived the bias appropriate to their mass function fit using the peak background split., derived the bias appropriate to their mass function fit using the peak background split.
 A slightly higher bias at fixed mass comes from assuming ellipsoidal collapse1999)., A slightly higher bias at fixed mass comes from assuming ellipsoidal collapse.
. The largest bias for halos of a given mass is given by the bias of the mass function. as computed by(1989).. which is very similar to the fit of for the masses of interest.," The largest bias for halos of a given mass is given by the bias of the mass function, as computed by, which is very similar to the fit of for the masses of interest."
 These two forms give biases larger. in the mass and redshift range of interest. than the lowest fit1999).," These two forms give biases larger, in the mass and redshift range of interest, than the lowest fit."
. The fit of lies vctween that of ancl(1999)., The fit of lies between that of and.
. For mass-thresholeled samples of dark matter halos in a large N-bodysimulation*.. we found the and orms bracketed the halo correlation function so this should represent the level of current uncertainty.," For mass-thresholded samples of dark matter halos in a large N-body, we found the and forms bracketed the halo correlation function so this should represent the level of current uncertainty."
 As it provides we most Conservative estimate of the maximum allowed scatter. and the numerical data lav closer to the form of(1974). we adopt this as our fiducial bias.," As it provides the most conservative estimate of the maximum allowed scatter, and the numerical data lay closer to the form of, we adopt this as our fiducial bias."
 llowever. we caution that the statistics in the simulation are »oor and. we also illustrate the elect of choosing the form below.," However, we caution that the statistics in the simulation are poor and we also illustrate the effect of choosing the form below."
" In several recent simulations at other redshifts and masses the orm appears to be a worse fit2005).. but unfortunately there are few direct N-body calibrations of b,(CAL) at the number densities and recshilts of interest to us."," In several recent simulations at other redshifts and masses the form appears to be a worse fit, but unfortunately there are few direct N-body calibrations of $b_h(M)$ at the number densities and redshifts of interest to us."
 One could imagine that for extremely rare peaks. whose collapse is much closer to sphericalLOSG).. these newer fits actually perform worse than the original formulation2008).," One could imagine that for extremely rare peaks, whose collapse is much closer to spherical, these newer fits actually perform worse than the original formulation."
. This challenging numerical problem deserves further investigation., This challenging numerical problem deserves further investigation.
" To interpret. the large-scale. clustering measurements of we then must specify INNCA,).", To interpret the large-scale clustering measurements of we then must specify $N(M_h)$.
" We wish to choose as simple a model as possible to illustrate the effect of scatter on the £bn relation. so we assume that the probability that a QSO is seen with instantaneous luminosity £ is log-normally around a central value Lo(Ad,) with a width σι Ifa given halo hosts a QSO. the probability that it is above some limiting luminosity. Lai. is then the integral of this expression from Livin to infinity. which can be expressed as an error function."," We wish to choose as simple a model as possible to illustrate the effect of scatter on the $\beff-\bar{n}$ relation, so we assume that the probability that a QSO is seen with instantaneous luminosity $L$ is log-normally around a central value $L_0(M_h)$ with a width $\sigma_L$ If a given halo hosts a QSO, the probability that it is above some limiting luminosity, $L_{\rm min}$, is then the integral of this expression from $L_{\rm min}$ to infinity, which can be expressed as an error function."
" H£ we assume that at most a fraction μι of halos host active QSOs at anyepoch and that LoxAZ; in the region of the turn-on. the mean number of QSOs above Lain in halos of mass My, is"," If we assume that at most a fraction $f_{\rm on}$ of halos host active QSOs at anyepoch and that $L_0\propto M_h^\alpha$ in the region of the turn-on, the mean number of QSOs above $L_{\rm min}$ in halos of mass $M_h$ is"
"simular to those identified in the I& baud (Table 21). except that the old disk eroup is further divided iuto ""bright? (ODb) aud “faint” (ODL) subgroups.","similar to those identified in the K band (Table \ref{tab_estiK}) ), except that the old disk group is further divided into “bright” (ODb) and “faint” (ODf) subgroups."
 Let us remember here that the method allows us o distiuguish eroups with similar mean kinematics but cliffereut Ihunuimosities (OD) and ODf) or groups with a similar huninosity distribution but different Xiaicnaties (D and ED for instance)., Let us remember here that the method allows us to distinguish groups with similar mean kinematics but different luminosities (ODb and ODf) or groups with a similar luminosity distribution but different kinematics (D and ED for instance).
 Moreover. if is nüportant to remark that D. OD) aud ED have. on average. a siuilu color iudex 25]2-—0.6 mag corresponding to a thick eircuustellar envelope. while ODf has a mean index of 0.1 mag that suggests that the majority of he stars in this last eroup have thin cuvelopes.," Moreover, it is important to remark that D, ODb and ED have, on average, a similar color index $25-12 = -0.6$ mag corresponding to a thick circumstellar envelope, while ODf has a mean index of 0.1 mag that suggests that the majority of the stars in this last group have thin envelopes."
 Let us finally poiut out tha the kincmatic parameters (Uy.Vo. Wu) associated with cach of the four eroups are very siuülar for both the 12 and 25 calibrations.," Let us finally point out that the kinematic parameters $U_0,V_0,W_0$ ) associated with each of the four groups are very similar for both the 12 and 25 calibrations."
 Once the parameter estimation aud eroup discrimination is conmipleted. cach star in our initial sample is attributed to one of the LPV exoups identified iu each xuidpass. following the method described iu Sect. 3..," Once the parameter estimation and group discrimination is completed, each star in our initial sample is attributed to one of the LPV groups identified in each bandpass, following the method described in Sect. \ref{Sect:statistics}."
 This allows us to estimate the most probable individual distance and absolute magnitude in cach band according o the observed astrometric. Kinematic aud plotometiric data and attributed eroup.," This allows us to estimate the most probable individual distance and absolute magnitude in each band according to the observed astrometric, kinematic and photometric data and attributed group."
 Due to the probabilistic nature of the Bayesian oocedure. some iisclassification is unavoidable.," Due to the probabilistic nature of the Bayesian procedure, some misclassification is unavoidable."
" To check and imiprove individual star assignations m each waveleugthn. we compare the calculated color iuclices cal-AM,Mx, (obtained fro the estimated individual absolute maeuitucdes deduced by the Davesiau assignations iu the Ay and/or A» wavelengths) with the observed color indices (obs=in, nx)."," To check and improve individual star assignations in each wavelength, we compare the calculated color indices $cal=M_{\lambda_1}-M_{\lambda_2}$ (obtained from the estimated individual absolute magnitudes deduced by the Bayesian assignations in the $\lambda_1$ and/or $\lambda_2$ wavelengths) with the observed color indices $obs=m_{\lambda_1}-m_{\lambda_2}$ )."
 of the stars are re-assigned to eroups reduciug the differences calobs for their iudices 25-12 aud/or K-12., of the stars are re-assigned to groups reducing the differences $cal-obs$ for their indices 25-12 and/or K-12.
 Figure 2 shows the histograms of difference of eof for the iudices 25-12 aud EK-12 of all the stars in the sample., Figure \ref{fig_omoinsc} shows the histograms of difference of $cal-obs$ for the indices 25-12 and K-12 of all the stars in the sample.
 These distributions are related to the errors of the individual estimated. huninosities aud of the observed magnitudes., These distributions are related to the errors of the individual estimated luminosities and of the observed magnitudes.
 We cau deduce that the accuracies of our estimated inclivicua huninosities are distributed according to a gaussian rule of standard error 0.3 nag., We can deduce that the accuracies of our estimated individual luminosities are distributed according to a gaussian rule of standard error 0.3 mag.
 and 0.1 mae., and 0.1 mag.
 respectively in I& aud IRAS bands., respectively in K and IRAS bands.
 The lower accuracy in the IRAS bands is consistent with the fact that IRAS photometry is more homogeneous than the Ik photometry. and that the variability amplitude of LPVs is stualler in the IRAS bauds than in Kk. As previously stated. he LA method has allowed us to take advantage of all the available information. cading to better estimations of the individual absolute magnitudes.," The lower accuracy in the IRAS bands is consistent with the fact that IRAS photometry is more homogeneous than the K photometry, and that the variability amplitude of LPVs is smaller in the IRAS bands than in K. As previously stated, the LM method has allowed us to take advantage of all the available information, leading to better estimations of the individual absolute magnitudes."
 Furthermore. the LM method has provided at he same time the statistical distribution of these individual imaguitudes (see Sect.," Furthermore, the LM method has provided at the same time the statistical distribution of these individual magnitudes (see Sect."
 aud d)., \ref{Sect:statistics} and \ref{Sect:groups}) ).
 The individual estimates of Ik. 12 aud 25 huninosities are eivou in a able available in electronic form at the CDS aud are inchided in the specialized ASTRID database," The individual estimates of K, 12 and 25 luminosities are given in a table available in electronic form at the CDS and are included in the specialized ASTRID database."
 The LM method gives unbiased calibrations for the basepopulation., The LM method gives unbiased calibrations for the base.
 It also gives individual kincmatic and photometric estimates for cach star of thesemple., It also gives individual kinematic and photometric estimates for each star of the.
 The dcistribution of these individual estimates (Sect. 5.1)), The distribution of these individual estimates (Sect. \ref{sec_discri}) )
 is. of course. biased by the sample selection criteria. contrarv to the group characteristics derived in Sect. L.," is, of course, biased by the sample selection criteria, contrary to the group characteristics derived in Sect. \ref{Sect:groups}."
 A comparison of the statistical properties of thesemigple with the calibrated parameters for the allows us to check thepopulation., A comparison of the statistical properties of the with the calibrated parameters for the allows us to check the.
 Let us analyze the representativity of our sample stars with respect to the kinematics., Let us analyze the representativity of our sample stars with respect to the kinematics.
 The observed proper notions and racial velocities. together with the estimated individual distances. allow us to compute the three velocity components (0VI) and the distance Z above 1ο ealactic plane for cach star in our sample.," The observed proper motions and radial velocities, together with the estimated individual distances, allow us to compute the three velocity components $(U,V,W)$ and the distance $Z$ above the galactic plane for each star in our sample."
 The mean sinelatical properties of our seuple derived from these individual estimates are shown in Table 45 for cach eroup of the IK and 12 xuxls., The mean kinematical properties of our sample derived from these individual estimates are shown in Table \ref{tab_kinbias} for each group of the K and 12 bands.
 Thev are verv simular to 1ο persineters deseribiug the groups (Tables 2 aud 3)). showing that," They are very similar to the parameters describing the groups (Tables \ref{tab_estiK} and \ref{tab_esti12}) ), showing that"
measurements and plots αἱ the Center for Astrophysics.,measurements and plots at the Center for Astrophysics.
 We also thank Alceeste Bonanos. Lica Oskinova. Myron Smith. Rick White. and an anonvinous referee for useful comments.," We also thank Alceste Bonanos, Lida Oskinova, Myron Smith, Rick White, and an anonymous referee for useful comments."
 This research was supported by erants GO5-6006 to the authors from the X-Ray Observatory Center. in the context of Cycle 6 Proposal 66200204 (PI. WLW).," This research was supported by grants GO5-6006 to the authors from the X-Ray Observatory Center, in the context of Cycle 6 Proposal 6200204 (PI WLW)."
 WLW also acknowledges partial support from SAO grants AR&-9003A and GOS-9021D. is operated by the Smithsonian Astrophysical Observatory under NASA contract NAS8-03060. which also provided publication support for this paper.," WLW also acknowledges partial support from SAO grants AR8-9003A and GO8-9021B. is operated by the Smithsonian Astrophysical Observatory under NASA contract NAS8-03060, which also provided publication support for this paper."
Therelore Tig) = oU). ASCFoc,"Therefore ) = ,T) + K ]."
al] In equilibrium O/V=—P., In equilibrium $\Omega/V = - P$.
 Comparing this with theexpression (3.2)) for the pressure. one can deduce that A=V and “Toed) V (fu dn my," Comparing this with theexpression \ref{pressure}) ) for the pressure, one can deduce that $K=V$ and ,T) = V ( f_0 - n_c )."
 Ilence we have obtained (he expansion around Che equilibrium states Tig) = VC ο Wilh One.PoleT) givengiv above.," Hence we have obtained the expansion around the equilibrium states ) = ,T) + V ] with $\Omega_0(\mu,T)$ given above."
 The probability (1j) to find the svstem with a particular value of 7 al eiven values of µ and Tis D)ec/T., The probability ${\cal P}(\eta)$ to find the system with a particular value of $\eta$ at given values of $\mu$ and $T$ is ).
 Along the coexistence curve mgs=fy(1)., Along the coexistence curve $n_c \mu = f_1(t)$.
 Then Tig) — =\'[fon? f," Then ) - ,T) = V ]."
ollall7||(84) mince /<Q. ο}«0 and the thermodynamic potential has (wo equal minima at the densities of the liquid and gas phases. of course.," Since $t<0$, $f_2(t)<0$ and the thermodynamic potential has two equal minima at the densities of the liquid and gas phases, of course."
 This potential is shown in Fie., This potential is shown in Fig.
 7 lor temperatures both below and above 7;., 7 for temperatures both below and above $T_c$.
 I1 this figure the volume was taken (ο be 400 [n., In this figure the volume was taken to be 400 $^3$ .
 Obviously the potentia scales proportionatelv with this volume., Obviously the potential scales proportionately with this volume.
 The value of 400 [m? is really quite optimistic for high energy nuclear collisions., The value of 400 $^3$ is really quite optimistic for high energy nuclear collisions.
 Considering that the critica densitv is estimated to be about 5»20.15 barvons/fn'. this would mean that about 300J) barvons participate in the fluctuation.," Considering that the critical density is estimated to be about $5n_0 \approx 0.75$ $^3$, this would mean that about 300 baryons participate in the fluctuation."
 That is a substaial fraction of the total of 394 in Au+AU. 116 in Ph+Ph. and 476 in U+U collisions.," That is a substantial fraction of the total of 394 in Au+AU, 416 in Pb+Pb, and 476 in U+U collisions."
" Even then. the potential for the low and high density phases are only 5 MeV below the unstable mid-point when TT,=0.0: it is even less as T,. is approached."," Even then, the potential for the low and high density phases are only 5 MeV below the unstable mid-point when $T/T_c = 0.6$; it is even less as $T_c$ is approached."
" lt is interesting to find 1he probability (hat (he svstem has some value of 7other than 5, or i along the curve of phasecoexistence.", It is interesting to find the probability that the system has some value of $\eta$other than $\eta_g$ or $\eta_l$ along the curve of phasecoexistence.
 The relative probability is (Pup, The relative probability is _l) =
with where the factors &; cusure that the ICIAIF is continuous where the power changes aud À is a normalization constant.,with where the factors $k_i$ ensure that the IGIMF is continuous where the power changes and $k$ is a normalization constant.
 &conim) equals 0 if οOAD. or in ο mPuuaxs Is the anit stellar mass.," $\xi_{\rm IGIMF}(m)$ equals 0 if $m<0.1 \, {\rm M}_{\odot}$ or $m> m_{\rm max *}$, where $m_{\rm max *}$ is the maximum stellar mass."
 Thus. the IGIME defined here is equal to the IME defined by Equation (1)). except for the high-iass slope aud the upper mass linüt.," Thus, the IGIMF defined here is equal to the IMF defined by Equation \ref{eq:IMF}) ), except for the high-mass slope and the upper mass limit."
 The case of a canonical IME in all star-clusters. be. a3=23. implies (7) cmp23.," The case of a canonical IMF in all star-clusters, i.e. $\alpha_3=2.3$, implies \citep{weidner2005a} $\alpha_{\rm IGIMF} \gtrsim 3$."
 The expectation value for the nuuber of SNIT per year would then be =0.2 per vear., The expectation value for the number of SNII per year would then be $\lesssim 0.2$ per year.
 On the other haud. OlcpuE&2 is possible. if a varying IME that beconies i1nore top-heavy with star-cluster mass Is considered (27)..," On the other hand, $\alpha_{\rm IGIMF} \lesssim 2$ is possible, if a varying IMF that becomes more top-heavy with star-cluster mass is considered \citep{weidner2010a,kroupa2012a}."
 This implies that the probability to actually observe four now SNIT in a eiveu year (2?) is essentially zero if the IME is canonical in all star-clusters; but it can still be about LO per cout if the IME becomes top-heavy in massive UCD-tvpe stir-cluscrs.," This implies that the probability to actually observe four new SNII in a given year \citep{lonsdale2006a} is essentially zero if the IMF is canonical in all star-clusters, but it can still be about 10 per cent if the IMF becomes top-heavy in massive UCD-type star-clusters."
 The remnants produced by SNII are neutron stars and blaek holes., The remnants produced by SNII are neutron stars and black holes.
 The SNILrates thereby are an 1idicator for how many mergers of such remnants can be detected by searching for eravitational waves., The SNII-rates thereby are an indicator for how many mergers of such remnants can be detected by searching for gravitational waves.
 Comparing the SN-rate for ogapng=3 to he SN-rate for OIGIME=2 thus sueeestsBS that about an order of magnitude more of such eveuts nay be expeced if the IME ln llassive star-clusters is uot canonical. bu top-leavy.," Comparing the SN-rate for $\alpha_{\rm IGIMF}=3$ to the SN-rate for $\alpha_{\rm IGIMF}=2$ thus suggests that about an order of magnitude more of such events may be expected if the IMF in massive star-clusters is not canonical, but top-heavy."
 Thus. he utherto predicted detection rate of about 30 uergers of dark remnants per wear (?) for the upcoming adLICO-experinen could be too low wean order of mmaenitude. as au invariant IME has οσο used for this estimate.," Thus, the hitherto predicted detection rate of about 30 mergers of dark remnants per year \citep{banerjee2010a} for the upcoming adLIGO-experiment could be too low by an order of magnitude, as an invariant IMF has been used for this estimate."
 Further evidence for a top-reavy IMIF in star-παιο galaxies is found by ? in Arp 299., Further evidence for a top-heavy IMF in star-bursting galaxies is found by \citet{anderson2011a} in Arp 299.
 They study uunibers of different types of πήραπονας in Arp 299 and couclude from the mass of the appropriate progenitor stars that the IAIF is probably top-heavy in that system., They study numbers of different types of supernovae in Arp 299 and conclude from the mass of the appropriate progenitor stars that the IMF is probably top-heavy in that system.
 Thus. ? qualitatively come to the same conclusion for," Thus, \citet{anderson2011a} qualitatively come to the same conclusion for"
The growth of supermassive black holes (SMBH) through mass accretion is accompanied by the release of enormous amounts of energy which. if it is not advected directly into the hole (see eg. Narayan Yi 1905). can be either radiated away. as in Quasars and bright Seyferts. or disposed of in kinetic form through powerful. collimated outflows or jets. as observed. for example. in Radio Galaxies.,"The growth of supermassive black holes (SMBH) through mass accretion is accompanied by the release of enormous amounts of energy which, if it is not advected directly into the hole (see e.g. Narayan Yi 1995), can be either radiated away, as in Quasars and bright Seyferts, or disposed of in kinetic form through powerful, collimated outflows or jets, as observed, for example, in Radio Galaxies."
 The feedback exerted by such a powerful energy release on the surrounding gas and stars is imprinted into the observed correlations between black hole mass and galaxy properties (Gebhardtetal.2000:Ferrarese&MerrittMar-coni&Hunt 2003). as well as into the disturbed morphology of the hot. X-ray emitting atmospheres of groups and clusters harboring active SMBH in their centers (Bóhringeretal.1993:Fabian2000:Churazovetal.2002:Birzan2004:Fabian 2006).," The feedback exerted by such a powerful energy release on the surrounding gas and stars is imprinted into the observed correlations between black hole mass and galaxy properties \cite{gebhardt:00,ferrarese:00,marconi:03}, as well as into the disturbed morphology of the hot, X-ray emitting atmospheres of groups and clusters harboring active SMBH in their centers \cite{boehringer:93,fabian:00,churazov:02,birzan:04,fabian:06}."
. Radiative and kinetic feedback differ not only on physical grounds. in terms of coupling efficiency with the ambient gas. but also on evolutionary grounds.," Radiative and kinetic feedback differ not only on physical grounds, in terms of coupling efficiency with the ambient gas, but also on evolutionary grounds."
 The luminosity dependent density evolution parameterization of the AGN luminosity functions (Hasinger.Miyaji&Schmidt2005:Hopkins.RichardsHernquist2007) implies a delay between the epoch of Quasar dominance and that of more sedate Radio Galaxies (Merloni2004:Merloni&Heinz 20060... such that kinetic energy feedback plays an increasingly important role in the epoch of cluster formation and virialization.," The luminosity dependent density evolution parameterization of the AGN luminosity functions \cite{hasinger:05,hopkins:07} implies a delay between the epoch of Quasar dominance and that of more sedate Radio Galaxies \cite{merloni:04,merloni:06h}, such that kinetic energy feedback plays an increasingly important role in the epoch of cluster formation and virialization."
 From the theoretical point of view. the idea that AGN were responsible for the additional heating needed to explain the cooling flow riddle (see e.g. Fabian et al.," From the theoretical point of view, the idea that AGN were responsible for the additional heating needed to explain the cooling flow riddle (see e.g. Fabian et al."
 2001 and reference therein) in, 2001 and reference therein) in
the trend is for haloes to become less spherical. and. more prolate.,"the trend is for haloes to become less spherical, and more prolate."
 This is in agreement with studies at low redshift and higher masses (Macció.Dutton.&vandenBosch2008:Betetal.2007:Alleood2006:Avila-Reesect 2005).," This is in agreement with studies at low redshift and higher masses \citep{Maccio08,Bett07,Allgood06,Avila05}."
. We also plot the two »oint. correlation function. €(r) of the haloes in Figure 9.. binning the haloes by their values of s anc T.," We also plot the two point correlation function, $\xi(r)$ of the haloes in Figure \ref{corr_shape}, binning the haloes by their values of $s$ and $T$."
 The correlation function represents the excess probability of finding a halo at a distance ro when compare to a random distribution of haloes., The correlation function represents the excess probability of finding a halo at a distance $r$ when compared to a random distribution of haloes.
 We use the same methoc as our previous work (Davis&Natarajan2009) to calculate £(r)., We use the same method as our previous work \citep{Davis09} to calculate $\xi(r)$.
 We calculate £(r) for haloes in four cdillerent bins. corresponding to Iarge anc small values of s (left column of Figure 9)) and Y (right column).," We calculate $\xi(r)$ for haloes in four different bins, corresponding to large and small values of $s$ (left column of Figure \ref{corr_shape}) ) and $T$ (right column)."
 The cuts were chosen so that one third of the haloes lie in cach bin., The cuts were chosen so that one third of the haloes lie in each bin.
 We also separate out haloes by mass so tha we eliminate any mass ellect on the value of £67)., We also separate out haloes by mass so that we eliminate any mass effect on the value of $\xi(r)$.
" In the top row of Figure 9 we show £(r) for haloes with AJ=10°M. and in the bottom row we show 10""M. haloes.", In the top row of Figure \ref{corr_shape} we show $\xi(r)$ for haloes with $M = 10^7 \Msun$ and in the bottom row we show $10^6 \Msun$ haloes.
 We find that nearly spherical haloes (values of s close to 1) are more clustered than the more aspherical haloes., We find that nearly spherical haloes (values of $s$ close to 1) are more clustered than the more aspherical haloes.
 This trencl is stronger for haloes in the LOOM. mass range than in the LO?M. mass range., This trend is stronger for haloes in the $10^7 \Msun$ mass range than in the $10^6 \Msun$ mass range.
 In the higher mass bin. we find an increase of 50% in the correlation of nearly spherical haloes when compared to the extremely aspherical halo sample.," In the higher mass bin, we find an increase of $50\%$ in the correlation of nearly spherical haloes when compared to the extremely aspherical halo sample."
 La the lower mass range. we report a smaller increase of only 20%.," In the lower mass range, we report a smaller increase of only $20\%$."
 Thus we conclude that our results at high. redshift ollow trends found at lower redshift (Paltenbacher&White2010:Bettetal.2007:Avila-Reese 2005).," Thus we conclude that our results at high redshift follow trends found at lower redshift \citep{Faltenbacher09,Bett07,Avila05}."
. Unlike the sphericity parameter. when we separate 1aloes into high and low values of Z. we see little change in £(r7).," Unlike the sphericity parameter, when we separate haloes into high and low values of $T$, we see little change in $\xi(r)$."
 This is somewhat surprising. as Faltenbacher&White(2010) lind at low redshift an olfset in €(r) when binning aaloes by 7.," This is somewhat surprising, as \citet{Faltenbacher09} find at low redshift an offset in $\xi(r)$ when binning haloes by $T$."
 Finally. we note tha the clustering streneth is weaker if an unnormalized moment of inertia tensor is used o caleulate s and 2 of the haloes.," Finally, we note that the clustering strength is weaker if an unnormalized moment of inertia tensor is used to calculate $s$ and $T$ of the haloes."
" In Figure 10.. we show the bias parameter, ἕμη. às a Function of the jio peak height vill.o.£(oCVE AL)D(S)). where D(z) is t10 growth function (Alo&White 2002)."," In Figure \ref{bias}, we show the bias parameter, $b = \sqrt{\xi_{\rm{MM}}/\xi_{\rm{HH}}},$ as a function of the halo peak height $\nu(M,z) = \delta_c / (\sigma(M) D(z))$ , where $D(z)$ is the growth function \citep{mo02}."
. This allows us to compare results from. two redshifts (2=I0 and z= 6) anc from 5r mass bins at each redshift., This allows us to compare results from two redshifts $z=10$ and $z=6$ ) and from 5 mass bins at each redshift.
 We split our halocs acc'ording to four properties: A. 8. d. and Css.," We split our haloes according to four properties: $\lambda$, $s$, $T$, and $C_{178}$."
 We report. that the spin parameter has the largest elfect on the bias. wule the triaxiality has the least.," We report that the spin parameter has the largest effect on the bias, while the triaxiality has the least."
 Fhis implies that angular momentum has the strongest dependence on environment of all these considered. variables and will be the source of the largest systematics due to dillering environments between haloes., This implies that angular momentum has the strongest dependence on environment of all these considered variables and will be the source of the largest systematics due to differing environments between haloes.
 Therefore. barvonic properties that depend. on the angular momentum of the host halo should have svstematic olfsets due to their local environment.," Therefore, baryonic properties that depend on the angular momentum of the host halo should have systematic offsets due to their local environment."
 This may plav a large role in the formation of the earliest. galactic disks. aswell as allect semi-analytic models which relate the barvonic spin to the dark matter halo spin (Crotonctal.2006:Benson&Bower2010. e," This may play a large role in the formation of the earliest galactic disks, aswell as affect semi-analytic models which relate the baryonic spin to the dark matter halo spin \citep[][e.g.]{Croton06,Benson10}, ,"
 This may plav a large role in the formation of the earliest. galactic disks. aswell as allect semi-analytic models which relate the barvonic spin to the dark matter halo spin (Crotonctal.2006:Benson&Bower2010. e.," This may play a large role in the formation of the earliest galactic disks, aswell as affect semi-analytic models which relate the baryonic spin to the dark matter halo spin \citep[][e.g.]{Croton06,Benson10}, ,"
 This may plav a large role in the formation of the earliest. galactic disks. aswell as allect semi-analytic models which relate the barvonic spin to the dark matter halo spin (Crotonctal.2006:Benson&Bower2010. e.9," This may play a large role in the formation of the earliest galactic disks, aswell as affect semi-analytic models which relate the baryonic spin to the dark matter halo spin \citep[][e.g.]{Croton06,Benson10}, ,"
 This may plav a large role in the formation of the earliest. galactic disks. aswell as allect semi-analytic models which relate the barvonic spin to the dark matter halo spin (Crotonctal.2006:Benson&Bower2010. e.9.," This may play a large role in the formation of the earliest galactic disks, aswell as affect semi-analytic models which relate the baryonic spin to the dark matter halo spin \citep[][e.g.]{Croton06,Benson10}, ,"
 This may plav a large role in the formation of the earliest. galactic disks. aswell as allect semi-analytic models which relate the barvonic spin to the dark matter halo spin (Crotonctal.2006:Benson&Bower2010. e.9.)," This may play a large role in the formation of the earliest galactic disks, aswell as affect semi-analytic models which relate the baryonic spin to the dark matter halo spin \citep[][e.g.]{Croton06,Benson10}, ,"
 This may plav a large role in the formation of the earliest. galactic disks. aswell as allect semi-analytic models which relate the barvonic spin to the dark matter halo spin (Crotonctal.2006:Benson&Bower2010. e.9.).," This may play a large role in the formation of the earliest galactic disks, aswell as affect semi-analytic models which relate the baryonic spin to the dark matter halo spin \citep[][e.g.]{Croton06,Benson10}, ,"
 This may plav a large role in the formation of the earliest. galactic disks. aswell as allect semi-analytic models which relate the barvonic spin to the dark matter halo spin (Crotonctal.2006:Benson&Bower2010. e.9.)..," This may play a large role in the formation of the earliest galactic disks, aswell as affect semi-analytic models which relate the baryonic spin to the dark matter halo spin \citep[][e.g.]{Croton06,Benson10}, ,"
the SEIT is known.,the SFH is known.
 The best-fit metallicity is ereater thai solar but for other cosmic SFIIs 9=2.5 aud a= 2). the best fit is around solar.," The best-fit metallicity is greater than solar but for other cosmic SFHs $\beta=2.5$ and $\alpha=2$ ), the best fit is around solar."
 Iu general. there is πμαι degeneracy between DP aud ietallicity for the data. tthe choice of metallicity does not sienificautly affect the constraints on the IME slope.," In general, there is minimal degeneracy between $\Gamma$ and metallicity for the data, the choice of metallicity does not significantly affect the constraints on the IMF slope."
 Note that there is an upper luit on the metallicity as a function of D due to the closed-box model (Fig. 7))., Note that there is an upper limit on the metallicity as a function of $\Gamma$ due to the closed-box model (Fig. \ref{fig:Z-imf+0}) ).
 Tusufficieut metals can be produced in the available time to raise the average metallicity above a certain lait and this lit decreases as the fraction of higl-ass stars in the IME decreases., Insufficient metals can be produced in the available time to raise the average metallicity above a certain limit and this limit decreases as the fraction of high-mass stars in the IMF decreases.
 This chemical evolution lit is derived from the PECGASE code using the upper vields of (1995)., This chemical evolution limit is derived from the PEGASE code using the upper yields of \cite{WW95}.
. The pattern shown in Figure 7 is similar as > is varied except the best-fit E shifts., The pattern shown in Figure \ref{fig:Z-imf+0} is similar as $\beta$ is varied except the best-fit $\Gamma$ shifts.
 To illustrate this. we now fix Z equal to 0.02 for the σαν ft aud show joiut confidence levels in versus P iu Figure 8..," To illustrate this, we now fix $\bar{Z}$ equal to 0.02 for the SBav fit and show joint confidence levels in $\beta$ versus $\Gamma$ in Figure \ref{fig:beta-imf}."
 We plot the result for three values of a., We plot the result for three values of $\alpha$.
 All the plots show a broad deseueracy with a slope of 97£1., All the plots show a broad degeneracy with a slope of $-7\pm1$.
 In other words. the best-fit E decreases by about O.LL for each |1 iucrease in 2.," In other words, the best-fit $\Gamma$ decreases by about $-0.14$ for each $+1$ increase in $\beta$."
 Note that the metallicity was chosen to be cousisteut with the solar ucieghborlood aud so that the closed-box model was valid out to P of 1.5 with only 1ininal extrapolation (see Fig. 7))., Note that the metallicity was chosen to be consistent with the solar neighborhood and so that the closed-box model was valid out to $\Gamma$ of 1.8 with only minimal extrapolation (see Fig. \ref{fig:Z-imf+0}) ).
 The choice of high redshift star formation (defined bv a) makes little difference on the best-fit D (Fie. 8))., The choice of high redshift star formation (defined by $\alpha$ ) makes little difference on the best-fit $\Gamma$ (Fig. \ref{fig:beta-imf}) ).
 Mavreinalizing over SEIT (./ à). we obtain a best fit for D in the range 0.851.3 conf.)," Marginalizing over SFH $\beta$ $\alpha$ ), we obtain a best fit for $\Gamma$ in the range 0.85–1.3 conf.)"
 aud a stroug upper lit of P<1.7 couf.)., and a strong upper limit of $\Gamma<1.7$ conf.).
 This is our principal result., This is our principal result.
 Iu other words. asstunine a declining SER from τ=1 to the present day (.J> 0.5) then the preseut luuinosity densities. in particular. the UV to optical colors mican that the upperanass IMFE slope cannot be steeper than 1.7.," In other words, assuming a declining SFR from $z=1$ to the present day $\beta\ge0.5$ ) then the present luminosity densities, in particular, the UV to optical colors mean that the upper-mass IMF slope cannot be steeper than 1.7."
 The caveats are: (1) PEGASE evolutionary tracks aud stellar spectra are suffücientle accurate: Gi) cosmic SEIT can be, The caveats are: (i) PEGASE evolutionary tracks and stellar spectra are sufficiently accurate; (ii) cosmic SFH can be
ihe mean direction of our SALA polarization field is consistent with the lower resolution single-dish data. which do not resolve the structure of FIR 5 and (traces a larger plivsical scale associated with the diffuse gas found at (he core envelope.,"the mean direction of our SMA polarization field is consistent with the lower resolution single-dish data, which do not resolve the structure of FIR 5 and traces a larger physical scale associated with the diffuse gas found at the core envelope."
 Our SATA CO (3 — 2) maps reveal a verv complex morphology possibly related to multiple outflows., Our SMA CO (3 $\rightarrow$ 2) maps reveal a very complex morphology possibly related to multiple outflows.
 Figure 6. shows the channel maps of the CO (3 — 2) emission with a velocity resolution of ~2.1IF..., Figure \ref{CO32} shows the channel maps of the CO (3 $\rightarrow$ 2) emission with a velocity resolution of $\sim 2.1$.
 At blueshifted velocities the emission arises fron an elongated but wiggling structure in the East-West direction., At blueshifted velocities the emission arises from an elongated but wiggling structure in the East-West direction.
 This blueshilted component appears (o be associated with FIR 6., This blueshifted component appears to be associated with FIR 6.
 The cistvibution of the molecular gas al the cloud systemic velocity cLO 1)) is basically associated with the FIR. 5 main core., The distribution of the molecular gas at the cloud systemic velocity $\simeq 10$ ) is basically associated with the FIR 5 main core.
 At redshiftecl velocities >|Hkms !)) there are two main elongated features almost parallel extending in the NorthSouth direction and observed over a wide range of velocities (up to £30 19)., At redshifted velocities $\ga 14$ ) there are two main elongated features almost parallel extending in the North–South direction and observed over a wide range of velocities (up to $\simeq 30$ ).
" One of these features is associated with the well-know outflow powered by FIR 5A (Sanders&Willner1935;Richer1996) and the other one is located about 10"" to the west ancl seems to arise from FIR 5-sw."," One of these features is associated with the well-know outflow powered by FIR 5A \citep{Sanders85,Richer92,Chandler96} and the other one is located about $10''$ to the west and seems to arise from FIR 5-sw."
 These two lobes have (heir brightest emission located near their associated dust components (FIR. 5À and sw)., These two lobes have their brightest emission located near their associated dust components (FIR 5A and sw).
" The emission presents a clumpy morphology. with an average angular size of ~5"". corresponding to a phvsical size of 0.01. parsees."," The emission presents a clumpy morphology, with an average angular size of $\sim 5''$, corresponding to a physical size of 0.01 parsecs."
 Ht is worth noting that the three possible CO high velocity features have no counterpart at the opposite flow velocities., It is worth noting that the three possible CO high velocity features have no counterpart at the opposite flow velocities.
 Thus. the North-South redshilted lobes have no blueshilted counterpart. and the EastWest blueshifted lobe does not have a redshifted counterpart.," Thus, the North-South redshifted lobes have no blueshifted counterpart, and the East–West blueshifted lobe does not have a redshifted counterpart."
 Figure 7 shows the Position-Velocity. (PV) diagram centered in FIR. 5À with a PA = (0*9 (along the brightest redshifted lobe)., Figure \ref{pvplot_co32} shows the Position-Velocity (PV) diagram centered in FIR 5A with a PA = $\degr\!$ .9 (along the brightest redshifted lobe).
 The outflowpowered bv FIR 5A has a wide distribution of velocities which prevails until 30 |..., The outflowpowered by FIR 5A has a wide distribution of velocities which prevails until $\sim 30$ .
 An extended spatial distribution, An extended spatial distribution
"Sclunidt(1959) X4, spr. (Sanduleak1969:Tlartwick1971).. Πα, (Ixeunicut X","\cite{1959ApJ...129..243S} $\Sigma_{\rm gas}$ $\Sigma_{\rm SFR}$ \citep{1969AJ.....74...47S, 1971ApJ...163..431H}, $\alpha$ \citep{1989ApJ...344..685K, 1998ApJ...498..541K}."
 Msp (IT2) (Wong 2011j).. ," $\Sigma_{\rm gas}$ $\Sigma_{\rm SFR}$ $\H2$ \citep{2002ApJ...569..157W, 2008AJ....136.2846B}. ."
The shape of the relation is also predicted to evolve strouelv with redshift due to the build-p of metallicity in the iuterstellar medimm (ISM) over cosmic listory and the importance of dust in the formation of Ἡυ and its shicldine from Lyman-Werner raciation (INruiholzetal.2009a:Cinedin&Iravtsov 2010)..," The shape of the relation is also predicted to evolve strongly with redshift due to the build-up of metallicity in the interstellar medium (ISM) over cosmic history and the importance of dust in the formation of $\H2$ and its shielding from Lyman-Werner radiation \citep{2009ApJ...693..216K, 2010ApJ...714..287G}."
 Tn contrast. the relation is often assuned to evolve little and be relatively insensitive to changes iu metallicity and interstellar radiation field. althoush this has not vet been confirmed obscrvationally.," In contrast, the relation is often assumed to evolve little and be relatively insensitive to changes in metallicity and interstellar radiation field, although this has not yet been confirmed observationally."
" The asswuption ou which this “universality” is based is that the efficiency with which clouds of moleculay hydrogen convert their Il, into stars is not a stroug fiction of the average ISM metallicity or the iterstellar radiation field. at least nuder conditions typical for spiral galaxies (INxiuubiolz&Tan 2007).."," The assumption on which this “universality” is based is that the efficiency with which clouds of molecular hydrogen convert their $\H2$ into stars is not a strong function of the average ISM metallicity or the interstellar radiation field, at least under conditions typical for spiral galaxies \citep{2007ApJ...654..304K}."
 The scale at which this conversion takes place is the scale of (eiaut) molecular clouds. i.c.. 100 pe or less.," The scale at which this conversion takes place is the scale of (giant) molecular clouds, i.e., 100 pc or less."
 However. there are a couple of complications.," However, there are a couple of complications."
" First. there is crowing observational evidence suggesting that the scatter In the relation increases if one goes to snaller aud sinaller —scalesNy, (sec. e$. Onoderaetal.2010:Sclxubaetal. 2010))."," First, there is growing observational evidence suggesting that the scatter in the relation increases if one goes to smaller and smaller scales (see, e.g., \citealt{2010arXiv1009.1971O, 2010arXiv1009.1651S}) )."
 Taken at face value. this seems to contradict a tight simall-scale coupling between molecular hydrogen surface density and star formation.," Taken at face value, this seems to contradict a tight small-scale coupling between molecular hydrogen surface density and star formation."
 Secoud. the relation is typically measured on ~ kpc scales —and the spatial averaging may lead to changes in the slope. intercept aud scatter compared with those on small scales (INxavtsov.2003)...," Second, the relation is typically measured on $\sim{}$ kpc scales and the spatial averaging may lead to changes in the slope, intercept and scatter compared with those on small scales \citep{2003ApJ...590L...1K}."
 Third. observationally determined ΕΙΝ are time-averaged over theeffectivelifetime of the specific star formation tracer aud may thus differ frou iustautauncous SFRs.," Third, observationally determined SFRs are time-averaged over theeffectivelifetime of the specific star formation tracer and may thus differ from instantaneous SFRs."
 À straightforward observational check of. the, A straightforward observational check of the
on-source spectra are selected based on the contiuuua nage bv setting a Lieh threshold. typically of the coutiuuun peak. aud includius iu the average ouly spectra toward pixels whose contiumunu brightuess is above this high threshold.,"on-source spectra are selected based on the continuum image by setting a high threshold, typically of the continuum peak, and including in the average only spectra toward pixels whose continuum brightness is above this high threshold."
 These ou-source spectra are averaged with weighting factor equal to the continua xiehtuess in each pixel. which optimizes the sigual-to-jolse ratio in the resulting absorption spectimm.," These on-source spectra are averaged with weighting factor equal to the continuum brightness in each pixel, which optimizes the signal-to-noise ratio in the resulting absorption spectrum."
 Similarly. he offsource spectra are selected toward pixels whose continui brightuess is below a low threshold. typically of the contimun peak.," Similarly, the off-source spectra are selected toward pixels whose continuum brightness is below a low threshold, typically of the continuum peak."
 The off-source spectra are rot simply averaged. but interpolated to eive a better xediction of the emission spectriun in the direction of the contimmun peak.," The off-source spectra are not simply averaged, but interpolated to give a better prediction of the emission spectrum in the direction of the continuum peak."
 This iuterpolation is based ou a simple d-linear fit (least squares fitting a linear function of two dimensions) doue independently for each spectral channel., This interpolation is based on a simple bi-linear fit (least squares fitting a linear function of two dimensions) done independently for each spectral channel.
" The outer boundary for which spectra are included iu he offsource interpolation is typically7"".. but for the extended supernova remnant C327.1)0.2 we extend the outer boundary to15%."," The outer boundary for which spectra are included in the off-source interpolation is typically, but for the extended supernova remnant G327.4+0.2 we extend the outer boundary to."
 The high aud low thresholds also have to be adjusted im some cases depeuding ou the coutimmun fux. down to for the ligh threshold of the fattest sources; and in the range to for the low threshold depeudiug ou the augular size of the contiuuun distribution.," The high and low thresholds also have to be adjusted in some cases depending on the continuum flux, down to for the high threshold of the faintest sources, and in the range to for the low threshold depending on the angular size of the continuum distribution."
 This iuterpolation process does uot change the fundamental angular resolution of the survey. nt is nof an extrapolation ou thewe plaue.," This interpolation process does not change the fundamental angular resolution of the survey, it is not an extrapolation on the plane."
 So the spectra derived for both the absorption aud the expected cussion still correspond. to one bea area of ronehly deliaineter., So the spectra derived for both the absorption and the expected emission still correspond to one beam area of roughly diameter.
 The absorption spectrum is then determuned by subtracting the interpolated off-sourece spectrum from the averaged on-source spectrum., The absorption spectrum is then determined by subtracting the interpolated off-source spectrum from the averaged on-source spectrum.
 The resultant optical depth spectrun. ©ὃν is the absorption spectrum divided bv the ουπα flux averaged over the pixels above the hieh threshold.," The resultant optical depth spectrum, $e^{-\tau}$, is the absorption spectrum divided by the continuum flux averaged over the pixels above the high threshold."
 The optical depth spectra towards ROW 91 and 0326.65|0.59 ave shown in Fig., The optical depth spectra towards RCW 94 and G326.65+0.59 are shown in Fig.
 6 (bottom plot) with corresponding interpolated offsource eniission spectra (upper plot)., \ref{fig:rcw94_abs} (bottom plot) with corresponding interpolated off-source emission spectra (upper plot).
" The errors in the absorption spectra are generally dominated by uucertaintv in the iuterpolated Cluission. particularly at low latitudes where the emission is not smoothly distributed on scales of a few arciuimnutes,"," The errors in the absorption spectra are generally dominated by uncertainty in the interpolated emission, particularly at low latitudes where the emission is not smoothly distributed on scales of a few arcminutes."
 To estimate the error in the absorption we compute the predictions for the off-source spectra based ou the bilinear fit to the emissiou and take the difference between these predictions and the actual spectra in each off-source pixel., To estimate the error in the absorption we compute the predictions for the off-source spectra based on the bilinear fit to the emission and take the difference between these predictions and the actual spectra in each off-source pixel.
 The ius average of these differences gives au error envelope ou the interpolated cussion iu the direction of the contimmun peak., The rms average of these differences gives an error envelope on the interpolated emission in the direction of the continuum peak.
 Dividing by the coutimuun flux averaged over the on-source pixels gives the optical depth error spectrum (clo). shown iu dotted lines ou Figs.," Dividing by the continuum flux averaged over the on-source pixels gives the optical depth error spectrum $\pm 1
\sigma$ ), shown in dotted lines on Figs."
 6 7.., \ref{fig:rcw94_abs}~ \ref{fig:g327abs}.
 For the absorption feature with most extreme velocity (nost negative or most positive) we define Vi. the lower lit on the distance.," For the absorption feature with most extreme velocity (most negative or most positive) we define $V_{\rm L}$ , the lower limit on the distance."
 Following Frail&Weisbere(1990) we define the velocity corresponding to the upper cistance lit. Vb. as the first cussion peak (Z5>35 K) bevoud Ty which docs not correspoud to any absorption features.," Following \citet{frail90} we define the velocity corresponding to the upper distance limit, $V_{\rm U}$, as the first emission peak $T_b > 35$ K) beyond $V_{\rm
L}$ which does not correspond to any absorption features."
 Upper distance limits are estimates ouly. since tle absecuce of an absorption line at the higher velocity is not conclusive evidence that the contiuuun source is ucarer than the HI enussion.," Upper distance limits are estimates only, since the absence of an absorption line at the higher velocity is not conclusive evidence that the continuum source is nearer than the HI emission."
 A region of 21-c1i enuüission may not show absorption because the cool gas may lave a covering factor less than one (ic. there may be gaps between absorbing clouds. filled oulv with warm. gas which docs not show detectable absorption).," A region of 21-cm emission may not show absorption because the cool gas may have a covering factor less than one (i.e. there may be gaps between absorbing clouds, filled only with warm gas which does not show detectable absorption)."
" This is unlikely for an emission region with colui deusity of 3107""cur? or more. which is iuplied by a line of brightucss temperature ereater than 35 Ix. Enuüssiou lines stronger than this alhiost ablwavs show some absorption. so our upper guit distauces should be inostlv valid."," This is unlikely for an emission region with column density of $3 \times 10^{20}~{\rm cm^{-2}}$ or more, which is implied by a line of brightness temperature greater than 35 K. Emission lines stronger than this almost always show some absorption, so our upper limit distances should be mostly valid."
 We asstume that the velocity errors are dominated by random cloud motious on the order of Glins (Dickey1997)., We assume that the velocity errors are dominated by random cloud motions on the order of 6 \citep{dickey97}.
 Velocity hits aud correspouding kinematic distances are given in Table 2. for sources brighter than ~SOOmJybeam+.," Velocity limits and corresponding kinematic distances are given in Table \ref{tab:abs} for sources brighter than $\sim 800~{\rm
mJy~beam^{-1}}$."
 Included in the table for comparison are the radio recombination lime velocities from CIIST., Included in the table for comparison are the radio recombination line velocities from CH87.
 Some individual sources are discussed in detail later., Some individual sources are discussed in detail later.
 The absorption velocities in Table 2. clearly lie in two doniünanut distributions. one centered around e= aand another ceutered around ¢=90.," The absorption velocities in Table \ref{tab:abs} clearly lie in two dominant distributions, one centered around $v=-50$ and another centered around $v=-90$."
 Calculated isovelocity contours from the Fich.Blitz.&Stark(1989) rotation curve are plotted in Fie.," Calculated isovelocity contours from the \citet{fich89}
 rotation curve are plotted in Fig."
 2 on the Tavlor&Cordes(1993) inodel of Galactic spiral aris to show the velocities covered by individual spiral aris aud the range of expected velocities for any given line of sight.," \ref{fig:diag} on the \citet{taylor93}
 model of Galactic spiral arms to show the velocities covered by individual spiral arms and the range of expected velocities for any given line of sight."
" The Test Region line of sight is marked by the ποσο,", The Test Region line of sight is marked by the wedge.
 Fig., Fig.
 2 shows that gas between ¢=20) aad eo=30 lis located in the Sagittarius-Carina aria., \ref{fig:diag} shows that gas between $v=-20$ and $v=-30$ is located in the Sagittarius-Carina arm.
 The position of this arii in velocity space is well traced by ireeions (Ceorechn&Ceoreclin1976:Caswell 1987).," The position of this arm in velocity space is well traced by regions \citep{georgelin76,caswell87}."
. Gas at velocities between e=50 aad οThkins iis located within the Scutum-Crus ari., Gas at velocities between $v=-50$ and $v=-75$ is located within the Scutum-Crux arm.
 As noted in CIIST there are many features around e=90 ffor which the correspoudence to a spiral feature is unclear., As noted in CH87 there are many features around $v=-90$ for which the correspondence to a spiral feature is unclear.
 Ceoreclin&Georeclin(1976). assign these rroegionus to the Norma arm., \citet{georgelin76} assign these regions to the Norma arm.
 At 7=3277 the line of sieht is nearly tanecut to this arm. which accounts for the large munber of sources seen there.," At $l\approx 327\arcdeg$ the line of sight is nearly tangent to this arm, which accounts for the large number of sources seen there."
 The distribution centered at c=ου hhas a large spread in velocity space. exteuding as far asthe terminal velocity near e=110Hl.," The distribution centered at $v=-90$ has a large spread in velocity space, extending as far asthe terminal velocity near $v=-110$."
 This laree velocity spread can also be explained by the line of sight renaming in the spiral aru for a significant distance near the arin tanegcut., This large velocity spread can also be explained by the line of sight remaining in the spiral arm for a significant distance near the arm tangent.
 Usine regions aud diffuse Ta cinission. Georgelinetal.(1991) similarly note velocity distributions at e=20th. e=dO hoande-2 65launs!.," Using regions and diffuse $\alpha$ emission, \citet{georgelin94} similarly note velocity distributions at $v=-20$, $v=-40$, and $v=-65$."
 They do not. however. detect Ila euissiou near ο90 bbecause it is bevoud the extinction lit.," They do not, however, detect $\alpha$ emission near $v=-90$ because it is beyond the extinction limit."
 Frou. the aabsorptiou we cetermuue new kinematic distances to two supernova remauts and confirm distances to a further one SNR aud nine regions., From the absorption we determine new kinematic distances to two supernova remnants and confirm distances to a further one SNR and nine regions.
 For SNR C328.1/022 our absorption spectruui looks verv similar to Gaensleretal. (2000)... who also found au extreme velocity of e=28|.," For SNR G328.4+0.2 our absorption spectrum looks very similar to \citet{gaensler00b}, , who also found an extreme velocity of $v=28$."
 The irceion velocities all correspoudwith the velocities given iu CIIST., The region velocities all correspondwith the velocities given in CH87.
 The new kinematic distances for SNR 6327.110.1, The new kinematic distances for SNR G327.4+0.4
For the seeing condition. we choose a typical value for Mauna Wea. rg=0.2 im at 0.5jnn. which correspouds to a hall aresecoud seeing in tlie visible.,"For the seeing condition, we choose a typical value for Mauna Kea, $r_0 = 0.2$ m at $0.5 \mu$ m, which corresponds to a half arcsecond seeing in the visible."
 Since the wavefront. varlauce scales as (0.2/rg)7. it is sullicient to present the results ouly for a single value of ry.," Since the wavefront variance scales as $(0.2/r_0)^{5/3}$, it is sufficient to present the results only for a single value of $r_0$."
 However. when we need to estimate the variauce for a site other than Mauna Ixea. this scaling does necessarily hold. since the altitucles and streugthis of individual turbulent layers determine the overall wavelront variance.," However, when we need to estimate the variance for a site other than Mauna Kea, this scaling does necessarily hold, since the altitudes and strengths of individual turbulent layers determine the overall wavefront variance."
" Results are plotted as the zenith angle vs. log),67. for Asc=1.25.1.65. and 2.2 jin respectively in Figs 1. 2. aud 3."," Results are plotted as the zenith angle vs. $\log_{10} \sigma^2$ , for $\lambda_{SCI} = 1.25, 1.65$, and 2.2 $\mu$ m respectively in Figs 1, 2, and 3."
" For another seeing coucition characterized by ry. variauces are obtained by shifting individual curves by (5/3)log,)(0.2/rq)."," For another seeing condition characterized by $r_0$, variances are obtained by shifting individual curves by $(5/3) \log_{10} (0.2/r_0)$."
 Since the Strehl ratio. S. is given by or for a small o7. (Tyson.R.Ik...1991).," Since the Strehl ratio, $S$, is given by or for a small $\sigma^2$, \citep{Tyson1991}."
". Therefore. logy,o?=—2 and -1. correspond to $=0.99 and 0.90 respectively."," Therefore, $\log_{10} \sigma^2 = -2$ and -1, correspond to $S = 0.99$ and 0.90 respectively."
 Another way of presenting our results is to show the behavior of the maximum allowed zenith distauce as a function of Aqpes. lor a given set of Asc aud the tolerauce level of tlie wavefront variance.," Another way of presenting our results is to show the behavior of the maximum allowed zenith distance as a function of $\lambda_{WFS}$, for a given set of $\lambda_{SCI}$ and the tolerance level of the wavefront variance."
 In Fies [. 2 and 6. such plots are given for Age; = 1.25. 1.65. ancl 2.2 jan. for the cases of o? = 0.01 and 0.05.," In Figs 4, 5 and 6, such plots are given for $\lambda_{SCI}$ = 1.25, 1.65, and 2.2 $\mu$ m, for the cases of $\sigma^2$ = 0.01 and 0.05."
 To see the relative contribution of each laver. its altitude relative to the summuuit of Mauna Ixea (1.2 kin). the beam separation r; in cm. the pliase variance. στ in rac? are given for Agey=1.65 jan. Ayes=0.7 qun. and 2=50° in Table 1.," To see the relative contribution of each layer, its altitude relative to the summit of Mauna Kea (4.2 km), the beam separation $r_i$ in cm, the phase variance, $\sigma_i^2$ in $^2$ are given for $\lambda_{SCI}=1.65$ $\mu$ m, $\lambda_{WFS} = 0.7$ $\mu$ m, and $z=50^\circ$ in Table 1."
 The total wavefront variance for this case is about 0.01., The total wavefront variance for this case is about 0.01.
 Two factors that determiue the relative significance of each laver. are the fractional integrated >C5. E]Fi. and beam separation. r;. since. o?>xjerE]2.," Two factors that determine the relative significance of each layer, are the fractional integrated $C_n^2$, $F_i$, and beam separation $r_i$, since $\sigma_i^2 \propto F_i \cdot r_i^{5/3}$."
) The largest contribution comes from the layer 8 kin above the Mauua Wea Observatory., The largest contribution comes from the layer 8 km above the Mauna Kea Observatory.
 rj; ata higli altitude is order of a lew cu., $r_i$ ata high altitude is order of a few cm.
 [If we cousider a Shack-Hartiman seusor with a 212x212 pixel CCD. attached to an 5 m," If we consider a Shack-Hartman sensor with a $\times$ 512 pixel CCD, attached to an 8 m"
and Stokes V signals at the blue and red wings (4280 from the line center) of the 6301.5 lline.,and Stokes V signals at the blue and red wings $\pm$ 280 from the line center) of the 6301.5 line.
" The Stokes V images at far wings have been used to identify Doppler shifts of Stokes V profiles (Ichimotoetal.,2007;Shimizuetal.,2008b;MartinezPillet 2009)."," The Stokes V images at far wings have been used to identify Doppler shifts of Stokes V profiles \citep{ichimoto2007,shimizu2008b,martinez2009}."
". In the sunspot penumbra, especially, many small patches having enhancements of Stokes V signals at either wing are commonly observed, as is shown in the middle and bottom panels in Fig. 1.."," In the sunspot penumbra, especially, many small patches having enhancements of Stokes V signals at either wing are commonly observed, as is shown in the middle and bottom panels in Fig. \ref{spmapclv}."
 It is realized that there are three kinds of patchy enhancements visible in the Stokes V maps of the penumbra., It is realized that there are three kinds of patchy enhancements visible in the Stokes V maps of the penumbra.
 The first one is the enhancement at the blue wing., The first one is the enhancement at the blue wing.
 It is easily seen that there are many patches with enhanced positive Stokes V signals in the middle panels of Fig. 1.., It is easily seen that there are many patches with enhanced positive Stokes V signals in the middle panels of Fig. \ref{spmapclv}.
 They are distributed all over the penumbra when the spot is located near the disk center on 5 Jan. while they are preferentially found in the center side penumbra when the spot is located away from the disk center on 6 and 7 Jan. Examples of Stokes profiles observed inside the patches indicated by the dashed arrows are shown in Figs., They are distributed all over the penumbra when the spot is located near the disk center on 5 Jan. while they are preferentially found in the center side penumbra when the spot is located away from the disk center on 6 and 7 Jan. Examples of Stokes profiles observed inside the patches indicated by the dashed arrows are shown in Figs.
 2 (a) and 3 (a)., \ref{sprof1} (a) and \ref{sprof2} (a).
" These Stokes V profiles are characterized by presence of humps at the blue wing with the same polarity as the spot, which indicates presence of enhanced blue shifts there."," These Stokes V profiles are characterized by presence of humps at the blue wing with the same polarity as the spot, which indicates presence of enhanced blue shifts there."
" The blue-shifted Stokes V profiles were studied by Rimmele&Marino(2006) and Ichimotoetal.(2007),, and it was found that they correspond to bright penumbral grains and are the inner footpoints of the Evershed flows where hot upflow occurs."," The blue-shifted Stokes V profiles were studied by \cite{rimmele2006} and \cite{ichimoto2007}, and it was found that they correspond to bright penumbral grains and are the inner footpoints of the Evershed flows where hot upflow occurs."
 The second kind of the patchy enhancements in the Stokes V maps is positive (white) enhancements at the red wing as, The second kind of the patchy enhancements in the Stokes V maps is positive (white) enhancements at the red wing as
eas and stars may be kinematically segregated in tically disturbed: disks.,gas and stars may be kinematically segregated in tidally disturbed disks.
 Reacercted eas seeks out. closed orbits. usually close to the disk plane. while reaccreted stars remain on highly eccentric orbits.," Reaccreted gas seeks out closed orbits, usually close to the disk plane, while reaccreted stars remain on highly eccentric orbits."
 This can produce dramatic cllects in further tidal interactions. possibly explaining some cases in which stellar ane tails fail to coincide. (Llibbearcl. Vacca. Yun 2000. Mihos 2001).," This can produce dramatic effects in further tidal interactions, possibly explaining some cases in which stellar and tails fail to coincide (Hibbard, Vacca, Yun 2000, Mihos 2001)."
 Figure 4 illustrates the behavior of gas and stars in the passage of encounter POL 1: €. The first frame shows two fairly relaxed gas disks (e&rev-scale). each embedded in a distended stellar disk (contours).," Figure \ref{fig04}
 illustrates the behavior of gas and stars in the passage of encounter POL 1:1 C. The first frame shows two fairly relaxed gas disks (grey-scale), each embedded in a distended stellar disk (contours)."
 In the second. frame the disks have moved. past each other. and a pronounced tail extends from the more direct. member of the pair.," In the second frame the disks have moved past each other, and a pronounced tail extends from the more direct member of the pair."
 By the third frame the eas and stellar distributions are quite dilferent the one is no longer a οσους predictor of the other., By the third frame the gas and stellar distributions are quite different – the one is no longer a good predictor of the other.
 For example. the gas in the lower tail crosses several contours of the stellar distribution: evidently the kinematically cold gas has produced a narrow tidal tail while the warmer stellar component has raised a broader feature.," For example, the gas in the lower tail crosses several contours of the stellar distribution; evidently the kinematically cold gas has produced a narrow tidal tail, while the warmer stellar component has raised a broader feature."
 This last frame also shows that the two galaxies are at the point of merging., This last frame also shows that the two galaxies are at the point of merging.
 In broad outline. the mergers resulting from these twelve encounters are similar to those described in earlier. studies (Negroponte White. 1983: Barnes LHlernquist. 1991. 1906).," In broad outline, the mergers resulting from these twelve encounters are similar to those described in earlier studies (Negroponte White 1983; Barnes Hernquist 1991, 1996)."
 Table 2. lists merger times Aene for all twelve encounters., Table \ref{tab:remnants} lists merger times $t_{\rm merger}$ for all twelve encounters.
 Gas usually has little effect. on. the overall cynamics of orbit decay and merging: of the four 1:1 € encounters. three merge within times Afeo:409£0.2 of the corresponding stellar-civnamical versions (Barnes 1998). (," Gas usually has little effect on the overall dynamics of orbit decay and merging; of the four 1:1 C encounters, three merge within times $\Delta t_{\rm
merger} \simeq \pm 0.2$ of the corresponding stellar-dynamical versions (Barnes 1998). ("
The exception is RET 1:1 €. where the mechanical friction of the gas hastens the merger by Λήμνος& 0.5.),"The exception is RET 1:1 C, where the mechanical friction of the gas hastens the merger by $\Delta t_{\rm
merger} \simeq 0.5$ .)"
 As mentioned in the introduction. gas which sullers strong shocks tends to fall into galactic nuclei before ancl during the merger. where it creates central concentrations of highly clissipatecdl gas.," As mentioned in the introduction, gas which suffers strong shocks tends to fall into galactic nuclei before and during the merger, where it creates central concentrations of highly dissipated gas."
 On the other hand. the gas which is involved in strong shocks retains a large fraction of its initia angular momentum. and is available to build extended disks like those shown in Figure 5..," On the other hand, the gas which is involved in strong shocks retains a large fraction of its initial angular momentum, and is available to build extended disks like those shown in Figure \ref{fig05}."
 Ideally. simulations would be run well past Zi.," Ideally, simulations would be run well past $t_{\rm merger}$."
 Bu as already noted. the later stages of these calculations are quite slow.," But as already noted, the later stages of these calculations are quite slow."
 Practical considerations forced me to stop a the times A44 listed in Table 2.., Practical considerations forced me to stop at the times $t_{\rm end}$ listed in Table \ref{tab:remnants}.
 Most remnants had fondaversos:+1 time units to relax: scaling the initial clisks to the Alilky Was. the galaxies in these experiments merge QU to 100vr before the end of the calculations.," Most remnants had $t_{\rm end} -
t_{\rm merger} \simeq 2 \pm 1$ time units to relax; scaling the initial disks to the Milky Way, the galaxies in these experiments merged $\sim 10^8$ to $10^9 {\rm\,yr}$ before the end of the calculations."
 These remnants may thus be compared. to. late-stage, These remnants may thus be compared to late-stage
absorption to the nucleus was much higher than that implied by the ον obtained by modelling the infrared. SED (for a Galactic dust:gas ratio). suggesting that much of the X-ray absorbing eas lies within the dust sublimation radius.,"absorption to the nucleus was much higher than that implied by the $A_{\rm{V}}$ obtained by modelling the infrared SED (for a Galactic dust:gas ratio), suggesting that much of the X-ray absorbing gas lies within the dust sublimation radius."
 In terms of its optical depth ancl racial extent. the obscuring dust in these two Compton-thin..| absorbed. X-ray background. sources more closely resembles the model of GD than that of Phy.," In terms of its optical depth and radial extent, the obscuring dust in these two Compton-thin, absorbed, X-ray background sources more closely resembles the model of GD than that of PK."
 Indeed. Cranato. Danese Lranceshini (1907) predicted. a sienificiant correlation between the LNB sources and those. appearing in7/50 surveys at2041... precisely because the GD models are (at the most) only moclerately optically thick over this wavelength range.," Indeed, Granato, Danese Franceshini (1997) predicted a significiant correlation between the HXB sources and those appearing in surveys at, precisely because the GD models are (at the most) only moderately optically thick over this wavelength range."
 An important dillerence. however. is that the GD tori have opening angles of 3545 degrees. whereas the covering fraction of any torus-like structure in the LIAB sources must be higher (c.g. Fabian lwawasa 1999 deduce that 85 per cent of the accretion power must be absorbed).," An important difference, however, is that the GD tori have opening angles of 35–45 degrees, whereas the covering fraction of any torus-like structure in the HXB sources must be higher (e.g. Fabian Iwawasa 1999 deduce that 85 per cent of the accretion power must be absorbed)."
 Note also that the DUSTY calculations assumed a spherical dust. geometry., Note also that the DUSTY calculations assumed a spherical dust geometry.
 Lt is thus appropriate to examine how such a high space covering obscuration could. be maintained against the dissipative forces which would tend to vield a Ilattened structure: at the sub-parsec scale inner radii found for sources A and D. Pier Ixrolik. (1992b) have shown that the pressure of the nuclear radiation on the grains can =jake the torus geometricallv thick. by both reducing the vertical component of gravity. ancl if the torus is sulliciently clumpsy. driving random motions of its constituent. clouds.," It is thus appropriate to examine how such a high space covering obscuration could be maintained against the dissipative forces which would tend to yield a flattened structure: at the sub-parsec scale inner radii found for sources A and B, Pier Krolik (1992b) have shown that the pressure of the nuclear radiation on the grains can make the torus geometrically thick, by both reducing the vertical component of gravity, and if the torus is sufficiently clumpy, driving random motions of its constituent clouds."
 Fabian et al. (, Fabian et al. (
1998) demonstrated: that a space-covering obscuration could be maintained with energy. input. from. a nuclear starburst withinpe.. but the distinct lack of emission lines in many of the newlv-discovered LINB sources iopears to rule this out.,"1998) demonstrated that a space-covering obscuration could be maintained with energy input from a nuclear starburst within, but the distinct lack of emission lines in many of the newly-discovered HXB sources appears to rule this out."
 More generally. the lack of ΑΝ emission lines in such objects may be due to cust within the narrow line region clouds.," More generally, the lack of AGN emission lines in such objects may be due to dust within the narrow line region clouds."
 Netzer Laor (1993) showec that its presence could significantly suppress line emission. through absorption of the photoionizing continuum anc destruction of line photons. and. thereby account for. the mismatch between the inferred covering factors of the broa and narrow line regions in classical AGN when dust is no considered. (the broad. line region lies just within the clus sublimation radius).," Netzer Laor (1993) showed that its presence could significantly suppress line emission, through absorption of the photoionizing continuum and destruction of line photons, and thereby account for the mismatch between the inferred covering factors of the broad and narrow line regions in classical AGN when dust is not considered (the broad line region lies just within the dust sublimation radius)."
 Our findings suggest the presence of a high covering fraction ofdusty gas in the new LNB sources., Our findings suggest the presence of a high covering fraction of gas in the new HXB sources.
 Alternatively. the lack of emission lines anc non-detections by could be used to argue that the new LNB sources are not actually obscured ACN. but. some other class of intrinsically hard. source (e.g. ΑΙ)Αν. as proposed for the ΠΗΝΙΟ by Di Matteo. Allen 1999): the present quality of the X-ray spectra is not high enough to discriminate between these two possibilities.," Alternatively, the lack of emission lines and non-detections by could be used to argue that the new HXB sources are not actually obscured AGN, but some other class of intrinsically hard source (e.g. ADAFs, as proposed for the HXB by Di Matteo Allen 1999); the present quality of the X-ray spectra is not high enough to discriminate between these two possibilities."
 Our present findings. however. strongly support. the obscured ACN hypothesis. by demonstrating that the energy. which is inferred. to be absorbed in the opticalUWX-ray. range is repadiated in the infrared.," Our present findings, however, strongly support the obscured AGN hypothesis, by demonstrating that the energy which is inferred to be absorbed in the optical–UV–X-ray range is reradiated in the infrared."
 La consequence. the results of deep. surveys of well-stucliecd fields (e.g. from the ELALS consortium: Rowan-Robinson ct al.," In consequence, the results of deep surveys of well-studied fields (e.g. from the ELAIS consortium; Rowan-Robinson et al."
 1999). are eagerly awaited.," 1999), are eagerly awaited."
 Lacdeed. as noted by 060. 5 of the 6 sources discovered. bv. Hornschemeier et al. (," Indeed, as noted by F00, 5 of the 6 sources discovered by Hornschemeier et al. ("
2000) in the LIDEN also haveJSOCAAL counterparts at listed by Aussel et al. (,2000) in the HDFN also have counterparts at listed by Aussel et al. (
1999) (CCNOLIIDEN 123648.1|621309 is the only source not detected),1999) (CXOHDFN 123648.1+621309 is the only source not detected).
 We are grateful to Z. Ivezié... M. Nenkova and. M. Elitzur for making DUSTY publically available. and likewise to Al. Bolzonella. 1. Pelló and J.-M. Miralles for HYPERZ.," We are grateful to Z. Ivezić,, M. Nenkova and M. Elitzur for making DUSTY publically available, and likewise to M. Bolzonella, R. Pelló and J.-M. Miralles for HYPERZ."
 Carolin Crawford. is thanked. for helpful. comments. and assistance., Carolin Crawford is thanked for helpful comments and assistance.
 RIW thanks PPARC for support. and P. the Isaac Newton Trust and the Overseas Research Trust.," RJW thanks PPARC for support, and PG the Isaac Newton Trust and the Overseas Research Trust."
 ACT thanks the Roval Society for support., ACF thanks the Royal Society for support.
A variety of surveys over the last two decades have revealed populations of galaxies that to be,"A variety of surveys over the last two decades have revealed populations of galaxies that emit the bulk of their radiation at infrared wavelengths \citep{Soifer84a,Soifer84b,Houck84,Houck85,Lefloch05,Perezgonzalez05}."
" responsible for most of the star formation activity in the distantUniverse (?).. while in the local Universe these galaxies are classified as Luminous (L;,>104 ΡΕ ) or Ultraluminous (Li,>107 L.) infrared galaxies (LIRGs/ULIRGs)."," Such objects are thought to be responsible for most of the star formation activity in the distantUniverse \citep{Lefloch05}, while in the local Universe these galaxies are classified as Luminous $_{ir}
> 10^{11}$ $_{\odot}$ ) or Ultraluminous $L_{ir} > 10^{12}$ $_{\odot}$ ) infrared galaxies (LIRGs/ULIRGs)."
 The morphologies of ULIRGs clearly sugeest that most of them are recent or ongoing mergers., The morphologies of ULIRGs clearly suggest that most of them are recent or ongoing mergers.
 The presence of an AGN in a subset of these objects is also well known. with some ULIRGs classified as QSOs at optical wavelengths.," The presence of an AGN in a subset of these objects is also well known, with some ULIRGs classified as QSOs at optical wavelengths."
 For a more complete review of the properties of LIRGs/ULIRGs see ?.., For a more complete review of the properties of LIRGs/ULIRGs see \cite{Sanders96}.
 Aluch attention has been paid recently to the so-called warn. ULIBGs (Gnid-inlrared colours of ο.0.2)..., Much attention has been paid recently to the so-called warm ULIRGs (mid-infrared colours of $_{25}$ $_{60}$$\ge$.
 These objects represent the 20-254 of the total ULIRG population discovered. bv the IRAS satellite., These objects represent the $\sim$ of the total ULIRG population discovered by the IRAS satellite.
" Most. of Chem have AGN optical spectra. very large molecular gas masses (Lj),~LOM M.: Sanders et al."," Most of them have AGN optical spectra, very large molecular gas masses $M_{H_{2}} \sim 10 ^{10}$ $_{\odot}$: Sanders et al."
 1955) and are found in an advanced merger state., 1988) and are found in an advanced merger state.
 In fact. ? found that of the warm ULIRGs in their sample of 12 objects were associated wilh single nucleus galaxies.," In fact, \cite{Surace99} found that of the warm ULIRGs in their sample of 12 objects were associated with single nucleus galaxies."
 These properties suggest (hat such objects represent a (ransitional stage between cool ULIRGs and Raclio Galaxies/QSOs (?):: cool ULIRGs evolve into warm ULIRGs on their wav (ο becoming QSOs/raclio galaxies., These properties suggest that such objects represent a transitional stage between cool ULIRGs and Radio Galaxies/QSOs \citep{Sanders88}: cool ULIRGs evolve into warm ULIRGs on their way to becoming QSOs/radio galaxies.
 Indeed. models predict that the tidal forces associated wilh mergers lead (0 large concentrations of eas and dust in the nuclear regions of the galaxies (22): such concentrations may trigger," Indeed, models predict that the tidal forces associated with mergers lead to large concentrations of gas and dust in the nuclear regions of the galaxies \citep{Mihos96,Barnes96}; ; such concentrations may trigger"
"where a,=a/0.1 is the disc viscosity. parameter and Ro=R,/LOOAU.","where $\alpha_{-1} = \alpha/0.1$ is the disc viscosity parameter and $R_2
=R_p/100$ AU."
"/ Tepe1 E migrationD is faster: We can improve these estimates by building a simple disc mocdel that would specify how cise properties change with lt. We define the stellar mass doubling time scale as fu,=M,Al where AJ is the mass accretion rate on the star. asstumed to be constant in time and radius inside the disc."," Type I migration is faster: We can improve these estimates by building a simple disc model that would specify how disc properties change with R. We define the stellar mass doubling time scale as $t_{\rm db} = M_*/\dot M$ where $\dot M$ is the mass accretion rate on the star, assumed to be constant in time and radius inside the disc."
 By the order of magnitude. fa. should be comparable to the free fall time of the host gaseous cloud from which the star forms. eg. ~107 ves for al ML. cloud at the typical interstellar temperature of 10 Ix (e.g..Larson1969).," By the order of magnitude, $t_{\rm db}$ should be comparable to the free fall time of the host gaseous cloud from which the star forms, e.g., $ \sim 10^5$ yrs for a 1 $\msun$ cloud at the typical interstellar temperature of $10$ K \citep[e.g.,][]{Larson69}."
. Using the standard Shakura&Sunvaey(1973) accretion disce formalism. we first find £202?) and. AL;C2) in the innermost disc.," Using the standard \cite{Shakura73} accretion disc formalism, we first find $H(R)$ and $M_d(R)$ in the innermost disc."
 We then find the radius beyond. which the Toomre(1964). parameter falls below unity., We then find the radius beyond which the \cite{Toomre64} parameter falls below unity.
 In the outer self-eravitating region of the cise we follow the treatment of Goodman(2003). for a disc with no internal energy sources and enforce (Qz(HH)(MM)=1., In the outer self-gravitating region of the disc we follow the treatment of \cite{Goodman03} for a disc with no internal energy sources and enforce $Q \approx (H/R)\;(M_*/M_d) = 1$.
" Finally. [or a constant accretion rate cise. Al=ἍπαΟν. where NacMaGRO)/ο the column depth of the disce. we can solve for The aceretion rate disc model built. in this wav is schematic but correctly predicts the expected. location of the sel-eravitational region at. &), few tens of AU to 100 AU."," Finally, for a constant accretion rate disc, $\dot M = 3 \pi \alpha H^2 \Omega \Sigma_d$, where $\Sigma_d \approx
M_d(R)/(\pi R^2)$ is the column depth of the disc, we can solve for The accretion rate disc model built in this way is schematic but correctly predicts the expected location of the self-gravitational region at $R_p >$ few tens of AU to 100 AU."
" Using this disc mocel we are now able to calculate the migration rate of a Gil at an arbitrary £2, fora given fii.", Using this disc model we are now able to calculate the migration rate of a GE at an arbitrary $R_p$ for a given $t_{\rm db}$.
 We follow the fit to numerical results by Dateetal.(2003). to smoothly join the type Land type H migration regimes., We follow the fit to numerical results by \cite{BateEtal03} to smoothly join the type I and type II migration regimes.
 The fastest migration> then occurs for embrvo mass Maui=AL., The fastest migration then occurs for embryo mass $M_{\rm emb} = M_t$.
" Phe Hills radius of the embryo is (for AZ,= 1M.) 1n general. raclial migration accelerates as the GE moves in. whereas the contraction of the envelope slows down with time (cf."," The Hill's radius of the embryo is (for $M_*=1 \msun$ ) In general, radial migration accelerates as the GE moves in, whereas the contraction of the envelope slows down with time (cf."
 equation 4))., equation \ref{r_ass}) ).
" Therefore. as planets migrate inwarels. there is a point where the embryos radius Zea, becomes comparable to the GE Lllls radius."," Therefore, as planets migrate inwards, there is a point where the embryo's radius $R_{\rm emb}$ becomes comparable to the GE Hill's radius."
 Phe tidal field rom the parent star at this point starts to peel olf the outer lavers of the embryo., The tidal field from the parent star at this point starts to peel off the outer layers of the embryo.
 The tidal racius. 4. is defined as the radius where Rau=opRy Gp Ld).," The tidal radius, $R_t$, is defined as the radius where $R_{\rm emb} = \eta_t R_{\rm H}$ $\eta_t \simlt 1$ )."
 Using the tvpe HL migration estimate as an example. we insert eq.," Using the type II migration estimate as an example, we insert eq."
 S. into eq. 4.," \ref{tII} into eq. \ref{r_ass},"
 and find In Appendix we show that irradiation bv the parent star can heat up the outer lavers of the envelope and even disrupt the whole GLE., and find In Appendix we show that irradiation by the parent star can heat up the outer layers of the envelope and even disrupt the whole GE.
 The effect appears less important than tidal disruption for Solar type stars. but may become dominant for higher mass stars.," The effect appears less important than tidal disruption for Solar type stars, but may become dominant for higher mass stars."
 Figure 1. shows the results of our giant embryo moclel described in 82 for embryos of 3. 6 and. LO Jupiter masses. shown with cillerent colours.," Figure \ref{fig:3panel_migr} shows the results of our giant embryo model described in 2 for embryos of 3, 6 and 10 Jupiter masses, shown with different colours."
" The parameters used for the figure are fai=10 ves. 0=OL. πο=0.01. M.=OAL. (as the star is assumed to be in the midst of its growth. rather than close to being completely assembled). Solar metalicitv. and the starting radial position is £2,=100 AU."," The parameters used for the figure are $t_{\rm db} = 10^5$ yrs, $\alpha=0.1$, $\kappa_0 = 0.01$, $M_* = 0.5 \msun$ (as the star is assumed to be in the midst of its growth, rather than close to being completely assembled), Solar metalicity, and the starting radial position is $R_p = 100$ AU."
 The upper panel shows racial migration of the embryos., The upper panel shows radial migration of the embryos.
 Note that initially all migrate at about the same rate (in the tvpe LE regime). but then the least massive embryo overtakes the other two.," Note that initially all migrate at about the same rate (in the type II regime), but then the least massive embryo overtakes the other two."
 The least massive one (black solid: curves) migrates the fastest as it happens to fall in the minimum of the migration time curve. between type L ancl type LE regimes (scefig.11inBateetal.2003).," The least massive one (black solid curves) migrates the fastest as it happens to fall in the minimum of the migration time curve, between type I and type II regimes \citep[see fig. 11
 in][]{BateEtal03}."
". The asterisks and the text ""core formation"" mark the time /—fay. when the solid core is assumed to form in the centre of the Gls (cef."," The asterisks and the text “core formation” mark the time $t=t_{\rm gr}$, when the solid core is assumed to form in the centre of the GEs (cf."
 eq. 6) , eq. \ref{tgr}) ).
The middle panel shows the embryo size. foun. and the respective Hill'sradius. for the three cases considered.," The middle panel shows the embryo size, $R_{\rm emb}$, and the respective Hill'sradius, for the three cases considered."
 We assume that when the two sets of curves intersect. the embryo is disrupted. leaving only the core. which migrates at a negligible rate during the calculation.," We assume that when the two sets of curves intersect, the embryo is disrupted, leaving only the core, which migrates at a negligible rate during the calculation."
" This is why the radial position curves. £,(/). approach a constant value in the upper panel at late times."," This is why the radial position curves, $R_p(t)$ , approach a constant value in the upper panel at late times."
the fit to the 7 spectra simultaneously was found to be 1.01 with 248 degrees of freedom.,the fit to the 7 spectra simultaneously was found to be 1.01 with 248 degrees of freedom.
 The PROJCT model fit to the spectra from all annuli is shown in Fig. 5.., The PROJCT model fit to the spectra from all annuli is shown in Fig. \ref{fig5}.
 Bolometrie luminosities and intrinsic absorption of the shells are presented in Table |., Bolometric luminosities and intrinsic absorption of the shells are presented in Table 1.
 The intrinsic absorption is consistent with zero in the outermost annuli. as expected if it is associated with the galaxy.," The intrinsic absorption is consistent with zero in the outermost annuli, as expected if it is associated with the galaxy."
" In order to derive the central gas density. a circular region with a radius of 7,=40 aresee was chosen and an APEC model was fit to the extracted spectrum."," In order to derive the central gas density, a circular region with a radius of $r_1=40$ arcsec was chosen and an APEC model was fit to the extracted spectrum."
" The APEC normalisation is defined as: n.ngdV.H where £2, is the angular distance to the source.", The APEC normalisation is defined as: n_e n_H dV where $D_a$ is the angular distance to the source.
 With the 9.7- fitted to the inner region we calculate the integral which is now written as the following: n.nydi322 8))dr., With the $\beta$ -model fitted to the inner region we calculate the integral which is now written as the following: n_e n_H r^2 )dr.
" where tan@=ry/r and assuming that n,= l.17ng.", where $\tan\theta=r_1/r$ and assuming that $n_e=1.17 n_H$.
" Knowing the APEC normalisation from the spectral tit. we found the central hydrogen nummber density. ng,=0.151x0.061 ."," Knowing the APEC normalisation from the spectral fit, we found the central hydrogen nummber density, $n_{H_{01}}=0.151\pm 0.061$ $^{-3}$."
" The value of ny,=0.010  (see equations 3) was then obtained from the continuity of the density at ρω=40 ηοσους.", The value of $n_{H_{02}}=0.010$ $^{-3}$ (see equations 3) was then obtained from the continuity of the density at $R_{cut}=40$ arcsec.
 The above calculations are based on the assumption that the emissivity of the gas simply scales as density squared., The above calculations are based on the assumption that the emissivity of the gas simply scales as density squared.
 Given the steep temperature gradient. this is not obviously justified.," Given the steep temperature gradient, this is not obviously justified."
 However. at these temperatures. emissivity is not a strong function of temperature.," However, at these temperatures, emissivity is not a strong function of temperature."
 In order to estimate the effect of temperature variations on the value of electron density. we calculate the electron number density directly from the emission of each shell (section 3.3) and compare with the model 3D electron number density obtained from the double «?-model fit.," In order to estimate the effect of temperature variations on the value of electron density, we calculate the electron number density directly from the emission of each shell (section 3.3) and compare with the model 3D electron number density obtained from the double $\beta$ -model fit."
 Fig 6 shows that the measured electron density is in reasonable agreement with that obtained from the -7-model analysis., Fig \ref{fig6} shows that the measured electron density is in reasonable agreement with that obtained from the $\beta$ -model analysis.
 Fig., Fig.
 7. shows the annular and de-projected temperature profiles out to 150 arcsec., \ref{fig7} shows the annular and de-projected temperature profiles out to 150 arcsec.
 The temperature profile shows a drop outward reaching 0.63 of its central value at à mean radius of 29 kpe., The temperature profile shows a drop outward reaching 0.63 of its central value at a mean radius of 29 kpc.
 It is useful to have an analytical temperature profile. which ean be used to derive the entropy profile.," It is useful to have an analytical temperature profile, which can be used to derive the entropy profile."
" We adopt a polytropic model: where r, and «7 are the same parameters as one obtains from a 3-model to the surface brightness profile."," We adopt a polytropic model:, where $r_c$ and $\beta$ are the same parameters as one obtains from a $\beta$ -model to the surface brightness profile."
 We fixed «7 at an average value of 0.53. from the inner and outer regions. and fitted the polytropic temperature profile. equation 6. with the core radius and  exponent as free parameters.," We fixed $\beta$ at an average value of 0.53, from the inner and outer regions, and fitted the polytropic temperature profile, equation 6, with the core radius and $\gamma$ exponent as free parameters."
" We found .=1.14 and +,=12.5 arcsec.", We found $\gamma=1.14$ and $r_c=12.5$ arcsec.
 The profile is presented in Fig 7.., The profile is presented in Fig \ref{fig7}.
 The innermost region gave a poor spectral fit., The innermost region gave a poor spectral fit.
 Here we investigate wo possible eauses., Here we investigate two possible causes.
 The temperature profile shows no sign of cooler gas in the central region. contrary to many X-ray groups of similar mass (Helsdon&Ponman2000a:Sunetal.2003:Tushotzkyetal... 2003).," The temperature profile shows no sign of cooler gas in the central region, contrary to many X-ray groups of similar mass \citep{hp00a,sun03,mush03}."
. Since the point source detection routine ound a point source at the centre of the galaxy in both low and vigh energy bands. we suspected that it might contain an active galactic nucleus (AGN).," Since the point source detection routine found a point source at the centre of the galaxy in both low and high energy bands, we suspected that it might contain an active galactic nucleus (AGN)."
 There is further motivation for this in the iterature — Goudfrooijetal.(1994). found an 446.|INL7] flux of IO.S(41.8)joH erg ?& + inthe core of 66482.," There is further motivation for this in the literature – \citet{gou94} found an $H\alpha+[NII]$ flux of $10.8(\pm1.8)\, \times 10^{-14}$ erg $^{-2}$ $^{-1}$ in the core of 6482."
 To check whether a central point source could be contaminating our spectral fit in the innermost region. which extends to r=d aresec. we excluded the central 2 aresee. which encompasses most of the point spread function of.. and monitored the tit quality.," To check whether a central point source could be contaminating our spectral fit in the innermost region, which extends to $r=4$ arcsec, we excluded the central 2 arcsec, which encompasses most of the point spread function of, and monitored the fit quality."
 We found no significant improvement in the fit compared to the original reduced of 1.12., We found no significant improvement in the fit compared to the original reduced of 1.12.
 The limited number ofcounts from the central point source does not permit any further analysis of a possible AGN component., The limited number ofcounts from the central point source does not permit any further analysis of a possible AGN component.
 An upper limit for the X-ray emission from the point source is estimated assuming, An upper limit for the X-ray emission from the point source is estimated assuming
We have written a code which can follow the chemical enrichment in a region as it evolves due to star formation. evolution and its feedback to the ISM.,"We have written a code which can follow the chemical enrichment in a region as it evolves due to star formation, evolution and its feedback to the ISM."
 This code allows up to two changes in the SE ellicieney. ancl gas. inllow and (using luminosity weighted. single stellar populations) predicts spectral features on the Lick/LDS svstem (Worthey et 11994)., This code allows up to two changes in the SF efficiency and gas inflow and (using luminosity weighted single stellar populations) predicts spectral features on the Lick/IDS system (Worthey et 1994).
 These predictions are used in comparison with observations of composite stellar. populations., These predictions are used in comparison with observations of composite stellar populations.
 Using this code we have attempted to fit observations of 10 early-type galaxies., Using this code we have attempted to fit observations of 10 early-type galaxies.
 We find that ellipticals are not made from primordial gas in a closed zone., We find that ellipticals are not made from primordial gas in a closed zone.
 Some pre-cnrichment mechanism. is required at carly times to produce the strong lines seen in earlv-tvpe galaxies., Some pre-enrichment mechanism is required at early times to produce the strong lines seen in early-type galaxies.
 In the current models this enrichment is achieved by varving the SE elliciencv. through increased SE rate after a time delay. and by inllow of gas enriched to the level of the existing ISM.," In the current models this enrichment is achieved by varying the SF efficiency through increased SF rate after a time delay, and by inflow of gas enriched to the level of the existing ISM."
 A delaved starburst) (from. increased SE elliciency and/or enriched gas inflow) can produce the observed strong lines in earlv-tvpe galaxies. without the need for a variation in the IME with time.," A delayed starburst (from increased SF efficiency and/or enriched gas inflow) can produce the observed strong lines in early-type galaxies, without the need for a variation in the IMF with time."
 This simulates the effect of a major merger of two massive. eas-rich galaxies.," This simulates the effect of a major merger of two massive, gas-rich galaxies."
 The strength of the LL? index constrains the age of the delayed burst in these models., The strength of the $\beta$ index constrains the age of the delayed burst in these models.
 We found that 7 of the 10 ealaxies could be fitted: with models in which metal rich stars. which dominate the light now. formed within the first [ow Gar.," We found that 7 of the 10 galaxies could be fitted with models in which metal rich stars, which dominate the light now, formed within the first few Gyr."
 Assuming a constant LAIP with time. this could be achieved either through a slightly delaved. burst. 2 or 3 Gyr) following appreciable star formation. or through an extended: period of enriched. inflow (over the first ~ 2 Cyr).," Assuming a constant IMF with time, this could be achieved either through a slightly delayed burst $*$ = 2 or 3 Gyr) following appreciable star formation, or through an extended period of enriched inflow (over the first $\sim$ 2 Gyr)."
 Phe other 3 galaxies required delaved bursts to explain their linc-streneths., The other 3 galaxies required delayed bursts to explain their line-strengths.
 These were the compact elliptical NGC 22] (te=15 Ger). an elliptical. NCC 5831. (te=10 Cir) and a BCC elliptical. NGC 2329 (1-12 Civr). assuming models which started forming 17 Civr ago.," These were the compact elliptical NGC 221 $*$ =15 Gyr), an elliptical NGC 5831 $*$ =10 Gyr) and a BCG elliptical NGC 2329 $*$ =12 Gyr), assuming models which started forming 17 Gyr ago."
 Pherefore there is evidence for later star formation in some ellipticals in groups and in clusters., Therefore there is evidence for later star formation in some ellipticals in groups and in clusters.
 From our fits to linc-strengths measured in ellipticals we find that in many cases the a-clement sensitive spectral features are systematically underestimated by the moclels. whilst iron sensitive features are overestimated when SSPs with solar ratios are used to model real elliptical galaxies.," From our fits to line-strengths measured in ellipticals we find that in many cases the $\alpha$ -element sensitive spectral features are systematically underestimated by the models, whilst iron sensitive features are overestimated when SSPs with solar ratios are used to model real elliptical galaxies."
 lt is interesting that there are a few exceptions to this general rule (e.g. NGC 221. 5831 and 5846) whose line-strengths appear to be consistent with solar abundance ratios.," It is interesting that there are a few exceptions to this general rule (e.g. NGC 221, 5831 and 5846) whose line-strengths appear to be consistent with solar abundance ratios."
 Allowing for non-solar ratios by using SSPs from Weiss et ((1995). we find that Ales and «Fez features in ellipticals can be well fitted with our mocdels.," Allowing for non-solar ratios by using SSPs from Weiss et (1995), we find that $_2$ and $<$ $>$ features in ellipticals can be well fitted with our models."
 Certain spectral features appear more useful. than others for constraining SE histories in galaxies., Certain spectral features appear more useful than others for constraining SF histories in galaxies.
 These include age indicators such as LL? and higher Balmer series lines (Jones Worthey 1995). plus metallicity sensitive features such as Fe5015 ancl Fe4668. whose ratio changes from ο.I around solar metalliity (\W94).," These include age indicators such as $\beta$ and higher Balmer series lines (Jones Worthey 1995), plus metallicity sensitive features such as Fe5015 and Fe4668 whose ratio changes from $<1$ to $>1$ around solar metallicity (W94)."
 Magnesiuni sensitive features (c.g. Mgb) shou also be included. for the purpose of exploring Mg/Ee dependencies., Magnesium sensitive features (e.g. $b$ ) should also be included for the purpose of exploring Mg/Fe dependencies.
 Such features should be put at a high. priority in observing proerammes when studying SE histories in galaxies., Such features should be put at a high priority in observing programmes when studying SF histories in galaxies.
 Ina future paper we will present spectral line-strengths in à sample of bulges in earlv-tvpe spirals (Proctor et 11997 in preparation) to find out if the current. models can fit the spectra. of spiral bulges and whether they also require delaved starbursts to explain their linc-strengths as do some of the ellipticals we have modelled in this paper., In a future paper we will present spectral line-strengths in a sample of bulges in early-type spirals (Proctor et 1997 – in preparation) to find out if the current models can fit the spectra of spiral bulges and whether they also require delayed starbursts to explain their line-strengths as do some of the ellipticals we have modelled in this paper.
 The authors wish to thanks David Fisher for helpful comments on this paper., The authors wish to thanks David Fisher for helpful comments on this paper.
"R leads to a net inward (outward) force, i.e. a position above (below) the equilibrium line in Fig. 1..","$R$ leads to a net inward (outward) force, i.e. a position above (below) the equilibrium line in Fig. \ref{fig:dec_N}."
 We proceed to calculate the perturbed locations in the N/1—V-diagram., We proceed to calculate the perturbed locations in the $N/l - V$ -diagram.
 We first consider a spherically symmetric cloud initially in a circular orbit at distance R to the black hole., We first consider a spherically symmetric cloud initially in a circular orbit at distance $R$ to the black hole.
" Since the temperature is kept constant by photoionisation, a change in by AR>0 will result in a change of the particle density by n(R+AR)=n(R)|(R+AR)/R] ?."," Since the temperature is kept constant by photoionisation, a change in by $\Delta R>0$ will result in a change of the particle density by $n(R+\Delta R) = n(R)
[(R+\Delta R)/R]^{-s}$ ."
" This results in a change of the column density: The corresponding change in orbital velocity for angular momentum conserving perturbations is: The perturbed cloud will receive a restoring force, if: It may be shown by a few lines of algebra that this condition may be fulfilled only if The analogous derivation for AR<0 leads to the same result."," This results in a change of the column density: The corresponding change in orbital velocity for angular momentum conserving perturbations is: The perturbed cloud will receive a restoring force, if: It may be shown by a few lines of algebra that this condition may be fulfilled only if The analogous derivation for $\Delta R<0$ leads to the same result."
" If the rotational velocity is below this value, positive perturbations to R lead to ejection."," If the rotational velocity is below this value, positive perturbations to $R$ lead to ejection."
 Negative ones lead to highly eccentric orbits., Negative ones lead to highly eccentric orbits.
 We show this by considering the effective potential γε., We show this by considering the effective potential $V_\mathrm{eff}$.
 Inserting eqs. (3)), Inserting eqs. \ref{dN}) )
" and (4)) into eq. (1)),"," and \ref{dV}) ) into eq. \ref{eq:focl}) ),"
 and defining Ro to be a fiducial where eq., and defining $R_0$ to be a fiducial where eq.
" 2 holds, with N(Ro)=No, V(Ro)=Vo and x=R/Ro, we derive: The effective potential is displayed in Fig. 2.."," \ref{Neq} holds, with $N(R_0)=N_0$, $V(R_0)=V_0$ and $x=R/R_0$, we derive: The effective potential is displayed in Fig. \ref{fig:Veff}. ."
" In agreement with the preceding discussion, there is an extremum at x—1, whose character depends on Vo: For Vo>Ve (compare eq. 6)),"," In agreement with the preceding discussion, there is an extremum at $x=1$, whose character depends on $V_0$: For $V_0>V_\mathrm{c}$ (compare eq. \ref{Vc}) ),"
 it is a minimum., it is a minimum.
 Stable bound orbits with small radial motions may be found in this case., Stable bound orbits with small radial motions may be found in this case.
" For Vo« Vc, the extremum at x=1 is a maximum."," For $V_0<V_\mathrm{c}$ , the extremum at $x=1$ is a maximum."
" In this case, their exists a minimum further in."," In this case, their exists a minimum further in."
" The equations can be solved easily analytically for s—0, 3/2 and 3."," The equations can be solved easily analytically for $s=0$, $3/2$ and 3."
" For s—3/2, the second extremum is at: This extremum is a minimum for Vo<V.."," For $s=3/2$, the second extremum is at: This extremum is a minimum for $V_0<V_\mathrm{c}$."
" We have verified numerically that for all 1<s3, ας is very close to the value of the s= 3/2-case, for Vo«Vc."," We have verified numerically that for all $1<s<3$, $x_\mathrm{e2}$ is very close to the value of the $s=3/2$ -case, for $V_0<V_\mathrm{c}$."
" Again, stable bound solutions may be found."," Again, stable bound solutions may be found."
" If a cloud is very deep in that potential well, the orbits are close to circular, and dominated by rotation."," If a cloud is very deep in that potential well, the orbits are close to circular, and dominated by rotation."
 The average column density and rotation velocity for such an orbit are much higher than the values at Ro., The average column density and rotation velocity for such an orbit are much higher than the values at $R_0$.
" Such orbits are therefore effectively on the stable, solid black part of the equilibrium curve in Fig. 1.."," Such orbits are therefore effectively on the stable, solid black part of the equilibrium curve in Fig. \ref{fig:dec_N}."
" Orbits with total energy close to the potential energy at the maximum will follow highly eccentric orbits, with predominantly radial kinematics."," Orbits with total energy close to the potential energy at the maximum will follow highly eccentric orbits, with predominantly radial kinematics."
" For a significant fraction of their orbital period, they have indeed low column densities and rotational velocities, corresponding to the region above the red dashed part of the line in Fig. 1.."," For a significant fraction of their orbital period, they have indeed low column densities and rotational velocities, corresponding to the region above the red dashed part of the line in Fig. \ref{fig:dec_N}."
" The luminosity in Eddington units for a geometrically thin accretion disk is a function of the polar angle and is given by 2l|cos0|, where / is the total luminosity of the source in Eddington units."," The luminosity in Eddington units for a geometrically thin accretion disk is a function of the polar angle and is given by $2 l |\cos \theta |$, where $l$ is the total luminosity of the source in Eddington units."
" The force is still central and therefore, angular momentum is conserved."," The force is still central and therefore, angular momentum is conserved."
 The orbits are planar., The orbits are planar.
 Any particular orbital plane may be characterised by a polar angle 05., Any particular orbital plane may be characterised by a polar angle $\theta_\mathrm{o}$.
 We define the direction of the maximum elevation above the equatorial plane to have an azimuth of @=0., We define the direction of the maximum elevation above the equatorial plane to have an azimuth of $\phi=0$.
" A cloud will then have maximum radiation pressure support at ϕ=0 and $= 7, and no radiation pressure support at ϕ=π/2 and ϕ= 32/2, when passing the equatorial plane."," A cloud will then have maximum radiation pressure support at $\phi=0$ and $\phi=\pi$ , and no radiation pressure support at $\phi=\pi/2$ and $\phi=3\pi/2$ , when passing the equatorial plane."
 Geometrical considerations result in: cos0=$4.," Geometrical considerations result in: $\cos \theta = \cos \phi
\, \cos \theta_\mathrm{o}$."
" Using this, and eqs (1)), (3)) and (4)), results in the following total force equation: where x=R/Ro, R(ó=x/2)Ro, N(ó=x/2)No."," Using this, and eqs \ref{eq:focl}) ), \ref{dN}) ) and \ref{dV}) ), results in the following total force equation: where $x=R/R_0$, $R(\phi=\pi/2)=R_0$, $N(\phi=\pi/2)=N_0$."
 The azimuthal velocity at ϕ=2/2 in Kepler units is denoted by Vo., The azimuthal velocity at $\phi=\pi/2$ in Kepler units is denoted by $V_0$.
" Because of the @-dependence, the force is no longer conservative."," Because of the $\phi$ -dependence, the force is no longer conservative."
" We therefore expect the total energy and hence also the major axis of the orbits to grow or diminish, depending on whether thecloud moves predominantlywith or against the radiation flux."," We therefore expect the total energy and hence also the major axis of the orbits to grow or diminish, depending on whether thecloud moves predominantlywith or against the radiation flux."
" First, we consider the force free locations, which are useful to understand the general structureof theallowed orbits."," First, we consider the force free locations, which are useful to understand the general structureof theallowed orbits."
" In the standard treatment of the Kepler problem, bound orbits require the"," In the standard treatment of the Kepler problem, bound orbits require the"
The spectroscopic observations have been obtained with Director's. Discretionary Time (DDT) program 273.B-5008.,The spectroscopic observations have been obtained with Director's Discretionary Time (DDT) program 273.B-5008.
 The spectra were taken with the VLT/FORS2 mask exchange unit MXU. targeting the five brightest UCD candidates in Abell 1689 from Mieske et al. (2004b))," The spectra were taken with the VLT/FORS2 mask exchange unit MXU, targeting the five brightest UCD candidates in Abell 1689 from Mieske et al. \cite{Mieske04b}) )"
 and several other objects of interest (stellar super cluster candidates. dwarf galaxies. ete.).," and several other objects of interest (stellar super cluster candidates, dwarf galaxies, etc.)."
" The slit-width was 1"". the resolution was about 6A//pixel at a spatial scale of 072 / pixel."," The slit-width was $\arcsec$, the resolution was about /pixel at a spatial scale of $\farcs 2$ / pixel."
" The PSF-FWHM was about 0.8"".", The PSF-FWHM was about $\arcsec$.
 The total on-source integration time was 6500 seconds., The total on-source integration time was 6500 seconds.
 The resulting S/N per pixel in the best exposed wavelength range around 6000 was about 15 for the brightest UCD candidate (;/=22.2 mag) and about 6 for the faintest one (4=23.8 mag)., The resulting S/N per pixel in the best exposed wavelength range around 6000 was about 15 for the brightest UCD candidate $i=22.2$ mag) and about 6 for the faintest one $i=23.8$ mag).
 The UCD candidate spectra are shown in Fig. 3..," The UCD candidate spectra are shown in Fig. \ref{spectra},"
 the derived redshifts in Table 1.., the derived redshifts in Table \ref{resspec}.
 We confirm the cluster membership for UCD candidates No.] and 3., We confirm the cluster membership for UCD candidates No.1 and 3.
 The redshift of candidate No.2 1s consistent with a galactic foreground star., The redshift of candidate No.2 is consistent with a galactic foreground star.
 The spectra of the two faintest UCD candidates No.4 and 5 were partially blended with neighbouring galaxies. inhibiting an unambiguous redshift determination.," The spectra of the two faintest UCD candidates No.4 and 5 were partially blended with neighbouring galaxies, inhibiting an unambiguous redshift determination."
 Both spectra exhibit several weak Cross correlation peaks in a range of redshifts. including peaks corresponding to AT689's distance.," Both spectra exhibit several weak cross correlation peaks in a range of redshifts, including peaks corresponding to A1689's distance."
 Longer spectroscopic integrations and better seeing are needed to definitely confirm or rule out their cluster membership., Longer spectroscopic integrations and better seeing are needed to definitely confirm or rule out their cluster membership.
 Judging from the thumbnail image in Fig. l..," Judging from the thumbnail image in Fig. \ref{mapcmd},"
" candidate No.4 might also be part of a lensed background the absolute luminosities of the two confirmed A1689 UCD candidates are almost identical to M32. such that one might also call them ""compact elliptical"" (cE) or ""compact dwarf elliptical"" (cdE)."," candidate No.4 might also be part of a lensed background the absolute luminosities of the two confirmed A1689 UCD candidates are almost identical to M32, such that one might also call them “compact elliptical” (cE) or “compact dwarf elliptical” (cdE)."
" We choose to assign the two cluster members the more general term ""compact galaxies"" (CGs) or ""M32-like galaxies"".", We choose to assign the two cluster members the more general term “compact galaxies” (CGs) or “M32-like galaxies”.
 In Figs., In Figs.
 2 and 4 the morphological appearance and surface brightness (SB) profiles of the two A1689 compact galaxies (CGs) is compared with a seeing convolved. simulated SB model of M32. as taken from the fit by Graham (2002)).," \ref{m32comp1} and \ref{m32comp2} the morphological appearance and surface brightness (SB) profiles of the two A1689 compact galaxies (CGs) is compared with a seeing convolved, simulated SB model of M32, as taken from the fit by Graham \cite{Graham02}) )."
 This fit consists of a central bulge component described by a Sersic profile with nzl1.5 and nag=105 pe and an underlying exponential disk with exponential scale length h=|NU pe (rag=SOG pe)., This fit consists of a central bulge component described by a Sersic profile with n=1.5 and $r_{\rm eff}=105$ pc and an underlying exponential disk with exponential scale length $h=480$ pc $r_{\rm eff}=806$ pc).
 Ellipticity was adopted as zero., Ellipticity was adopted as zero.
 The PSF profile used for convolution was adopted as à Moffat function with FHWM=0.095” and 3}=2.5., The PSF profile used for convolution was adopted as a Moffat function with $=0.095''$ and $\beta=2.5$.
" For the simulation. an angular distance scale of 3.10 kpe/” was assumed. corresponding to a mean cluster redshift of Do0.1832 (NED database). Πω=το kms ! |. Qa,= 0.3 and O4—0.1. On the ACS image in Fig. 2.."," For the simulation, an angular distance scale of 3.10 $\arcsec$ was assumed, corresponding to a mean cluster redshift of $z=0.1832$ (NED database), $H_{\rm 0}=70$ km $^{-1}$ $^{-1}$, $\Omega_M=$ 0.3 and $\Omega_{\Lambda}=0.7$ On the ACS image in Fig. \ref{m32comp1},"
 the two AI689 CGs and the simulated M32 appear very similar and are notably more extended than a pure stellar source of comparable luminosity., the two A1689 CGs and the simulated M32 appear very similar and are notably more extended than a pure stellar source of comparable luminosity.
 In the surface brightness plots of Fig. 4..," In the surface brightness plots of Fig. \ref{m32comp2},"
 this is confirmed., this is confirmed.
 When fitting an exponential profile to the outer parts of the two CG profiles that are mostly unaffected by the seeing. effective radii in the range 700-800 pc are determined. in agreement with the scale size of M32's exponential disk.," When fitting an exponential profile to the outer parts of the two CG profiles that are mostly unaffected by the seeing, effective radii in the range 700-800 pc are determined, in agreement with the scale size of M32's exponential disk."
 In the lower panel of Fig. 4..," In the lower panel of Fig. \ref{m32comp2},"
 itis shown that also a single Sersic profile with ra= pe and η=1.85 fits quite well the observed CG profiles., it is shown that also a single Sersic profile with $r_{\rm eff}=300$ pc and $n=1.85$ fits quite well the observed CG profiles.
 In order to explore the acceptable range for rq in the case of asingle component. we used the ISHAPE task (Larsen 1999)).," In order to explore the acceptable range for $r_{\rm eff}$ in the case of a single component, we used the ISHAPE task (Larsen \cite{Larsen99}) )."
" As input PSFs we used three different unresolved sources in the / exposure of the ACS image. with PSF-FWHM between 0.09 and 0.10""."," As input PSFs we used three different unresolved sources in the $i$ exposure of the ACS image, with PSF-FWHM between 0.09 and $''$."
 For the intrinsic profile shape we adopted King profiles with concentration parameter between 5 and 30 and a Moffat profile with ./= 2.5., For the intrinsic profile shape we adopted King profiles with concentration parameter between 5 and 30 and a Moffat profile with $\beta= 2.5$ .
" The mean and scatter of the fitted rog Obtained with different models and PSFs is 370 pe + 150 pe for CGaigso., and 225 pe + 75 pe for X169»."," The mean and scatter of the fitted $r_{\rm eff}$ obtained with different models and PSFs is 370 pc $\pm$ 180 pc for $_{\rm A1689,1}$ and 225 pc $\pm$ 75 pc for $_{\rm A1689,2}$."
 All of the profile fits have a comparable 47., All of the profile fits have a comparable $\chi ^2$.
 Doing the same procedure for the simulated. seeing convolved M32 image yields an intrinsic effective radius of 225 pe + 75 pe.," Doing the same procedure for the simulated, seeing convolved M32 image yields an intrinsic effective radius of 225 pc $\pm$ 75 pc."
 Objects like M32 are extremely rare., Objects like M32 are extremely rare.
 In thisletter we have presented the discovery of two M32 analogs in Abell 1689. an," In thisletter we have presented the discovery of two M32 analogs in Abell 1689, an"
Optical microvariability (OM: variations with sanall amplitude in fiue scales from münutes to hours) iu quasars can be a powerful tool to constrain the energv ciuission models of active ealactic nuclei;,Optical microvariability (OM; variations with small amplitude in time scales from minutes to hours) in quasars can be a powerful tool to constrain the energy emission models of active galactic nuclei.
 Some studies indicate that OM does not depend on the radio properties of quasars (sce deDiegoctal.1998.. hereafter PI: Stalinct 200k: Cupta&Joshi 2005: Ramirezetal. 2009.. hereafter PIT).," Some studies indicate that OM does not depend on the radio properties of quasars (see \citealt{de98}, hereafter PI; \citealt{Stalin04}; \citealt{Gupta05}; \citealt{Ram09}, hereafter PII)."
 However. the imechauism responsible for OM. has not been characterized.," However, the mechanism responsible for OM has not been characterized."
 Concerning this topic. variability studies iu tlic optical bands acquired a particular relevance. because the optical/UV excess is identified. witcf the euission from the accretion disk (es... Ixrisha&Wita 199k: Lawrence 2005: Pereyraetal.x 2006: Li&Cao 2008: Wilhiteetal. 2008)).," Concerning this topic, variability studies in the optical bands acquired a particular relevance, because the optical/UV excess is identified with the emission from the accretion disk (e.g., \citealt{Krishan94}; \citealt{Lawrence05}; \citealt{Pereyra06}; \citealt{Li08}; \citealt{Wilh08}) )."
 Unfortunately. all the proposed scenarios are very complex when the enüssion frou additional spectral components is considered.," Unfortunately, all the proposed scenarios are very complex when the emission from additional spectral components is considered."
 Iu photometric studies. the properties of the spectral euergev distribution (SED) of quasars era characterized by the spectral index. which is defined as the slope of a curve ina plane {ουfp)eersuslog(r). and calculated as maguitude differeuces at different bands (e.g.1998: Trevese&Vaenctti 2002)).," In photometric studies, the properties of the spectral energy distribution (SED) of quasars era characterized by the spectral index, which is defined as the slope of a curve ina plane $log(f_\nu)~versus~log(\nu)$, and calculated as magnitude differences at different bands (e.g.,; \citealt{Treve02}) )."
 During variability events. changes in the shape of this SED can give us inportaut information about the enission processes.," During variability events, changes in the shape of this SED can give us important information about the emission processes."
 For instance. evidence from variability studies strouglv indicates that OM has nou-thermal naturein both BL Lac objects aud flat spectrum radio quasars (FSROs: c.¢.. D'muicisetal. 2002: Vaenettictal. 2003: Villataet 2001b:: Villatactal.200 La:: Thietal. 2006a.. Thietal. 2006b.. Duetal. 2007)). but in the case of FSRQs the coutributiou of a thermal component must be considered in the spectrum of these objects (e.g. Cuetal. 2006: Thietal.20060: Ramirezctal.200 1)).," For instance, evidence from variability studies strongly indicates that OM has non-thermal naturein both BL Lac objects and flat spectrum radio quasars (FSRQs; e.g., \citealt{Damicis02}; \citealt{Vagne03}; \citealt{Villata04b}; \citealt{Villata04a}; \citealt{Hu06a}, \citealt{Hu06b}, , \citealt{Hu07}) ), but in the case of FSRQs the contribution of a thermal component must be considered in the spectrum of these objects (e.g., \citealt{Gu06}; ; \citealt{Hu06b}; \citealt{Ram04}) )."
 Concerning quasars. sliort-term variability (variations with amplitudes of tenths of a maenitude aud time scales of weeks to mouths) nüeht have thermal origin (e... Trevese&Vaenetti 2002)).," Concerning quasars, short-term variability (variations with amplitudes of tenths of a magnitude and time scales of weeks to months) might have thermal origin (e.g., \citealt{Treve02}) )."
 Nevertheless. all these stucies usually eimiplov two bands. or oulv the average of color values is used. losing thus information abouttesture of shortest variations (sce 1999: Webb&Alalkan 2000: Tróvese 2002::Vaenettietal. 2003).," Nevertheless, all these studies usually employ two bands, or only the average of color values is used, losing thus information about of shortest variations (see \citealt{Giveon99}; ; \citealt{Webb00}; ; \citealt{Treve02}; ;\citealt{Vagne03}) )."
 Iu order to discern where the OAL in quasars originates. we assunie that ciission from the accretion disk mustbe related to thermal processes.," In order to discern where the OM in quasars originates, we assume that emission from the accretion disk mustbe related to thermal processes,"
All these distrimtons use metallicities estimated by photometric data.,All these distributions use metallicities estimated by photometric data.
 However. they differ in the selection criteria and calibraIOUS 1sed.," However, they differ in the selection criteria and calibrations used."
 In spite of these differences. the agreement of the allicity distributions (Figures 2 and 3) is very good.," In spite of these differences, the agreement of the metallicity distributions (Figures 2 and 3) is very good."
 The fraction of stars with [Fe/T]| < O.L0. for each clistrilnition. is preseuted in Table 2..," The fraction of stars with [Fe/H] $<-0.40$, for each distribution, is presented in Table \ref{fraction}."
 From these data. 1 call © ostimated that around (22 + of the late-tvpe dwarfs in our neighbourlood shold have metallicities lower twn 0.10 dex Note also that al distributions. except that by Ενα Morell (1997). show apronunent single peak around —0.20 dex.," From these data, it can be estimated that around (22 $\pm$ of the late-type dwarfs in our neighbourhood should have metallicities lower than $-0.40$ dex Note also that all distributions, except that by Flynn Morell (1997), show a prominent single peak around $-0.20$ dex."
 As shown bv Rocha-Pinto Maciel (19975)). this feature could be expained by au iuteuse star formation era from 5 to 8 Car ago.," As shown by Rocha-Pinto Maciel \cite{RPM97b}) ), this feature could be explained by an intense star formation era from 5 to 8 Gyr ago."
 Therefore. the main conclusion that can be drawn from the comparisous above is that there is a remarkable cousiscacy amonest the distributions of F. C and Is dwarfs.," Therefore, the main conclusion that can be drawn from the comparisons above is that there is a remarkable consistency amongst the distributions of F, G and K dwarfs."
 This cousisteucyv could only be attained if the chemical euricnuent and star formation lüstorvdwarfs., This consistency could only be attained if the chemical enrichment and star formation history.
 The raw data of tl1ο Inetallicitv. distributions are often subject to a varicty* of corrections due to observational Crrors. Cosmic scatter and scale height effects.," The raw data of the metallicity distributions are often subject to a variety of corrections due to observational errors, cosmic scatter and scale height effects."
 When a sample has stars with lifetimes lower than the disk age.aep corrections due to stellar evolution ist also be applied.," When a sample has stars with lifetimes lower than the disk age, corrections due to stellar evolution must also be applied."
 Such corrections are needed to couvert the imetallieitv distribution iuto tje distribution., Such corrections are needed to convert the metallicity distribution into the distribution.
 For a distribution based ou svectroscopic [Fe/T]. these correctious are eenerallv πιficient.," For a distribution based on spectroscopic [Fe/H], these corrections are generally sufficient."
 Ilowever. for photometric distributions there is an additional correction which has been totally ucelected iji past studies;," However, for photometric distributions there is an additional correction which has been totally neglected in past studies."
 This correction ds needec in order to take iuto account the effects of tje chromosplieric activiv on the photometric iudices., This correction is needed in order to take into account the effects of the chromospheric activity on the photometric indices.
 Dv studving the uetallicity distribution in a sample of 730 late-type dwarts with varving levels of chromospheric activity. Rocha-Pinto Alaciel (1998)) have shown tha for the acive stars. the difference between the spectroscojc aud i6 photometric metallicity. increases systematically as a function of the stellar activity.," By studying the metallicity distribution in a sample of 730 late-type dwarfs with varying levels of chromospheric activity, Rocha-Pinto Maciel \cite{RPM98}) ) have shown that, for the active stars, the difference between the spectroscopic and the photometric metallicity increases systematically as a function of the stellar activity."
 This LOS!It is a consequence of the ny deficiency. which is more VOj0unced in active binaries (Cünuénnez et al 1991)). mt actually seenis also to be present iu normal active stars (Cdampapa e al. 1979:," This result is a consequence of the $m_1$ deficiency, which is more pronounced in active binaries (Giménnez et al \cite{gimenez}) ), but actually seems also to be present in normal active stars (Giampapa et al. \cite{giampapa};"
 Dazii et al. 1989:, Basri et al. \cite{basri};
 Morale ( al. 1996))., Morale et al. \cite{morale}) ).
 A imetallicitv. distribution that does not ake iuto account this effect will be biased towards moetal-voor stars., A metallicity distribution that does not take into account this effect will be biased towards metal-poor stars.
 The elimination of ideuified active stars from 1e sample is no an ideal solution to this problem as. iu single late-type dwarts. the acivitv ds luked o the stellar age (Soderbom ct al. 19913).," The elimination of identified active stars from the sample is not an ideal solution to this problem as, in single late-type dwarfs, the activity is linked to the stellar age (Soderblom et al. \cite{soder91}) )."
 Sunales free of active stars will be also free of young stars. which will iutroduce another bias. in the sense of avoiding the expected inetal-richer dwarfs.," Samples free of active stars will be also free of young stars, which will introduce another bias, in the sense of avoiding the expected metal-richer dwarfs."
 Even if there was no relaion between age aud activity. there wotId always remain some uuideitid activo stars in the photometric surveys. as we do not kuow how to idenitv such stars from their indices.," Even if there was no relation between age and activity, there would always remain some unidentified active stars in the photometric surveys, as we do not know how to identify such stars from their indices."
 The only wav to keep a minim compromise between the achievement of a LObiased sample aud an accurate metallicity. distinition is to make use of approximate corrections for the effects of the chromospheric activity., The only way to keep a minimum compromise between the achievement of a non-biased sample and an accurate metallicity distribution is to make use of approximate corrections for the effects of the chromospheric activity.
" The corrections we are proposing assume that all active stus, for which the chromospheric index logRy>105 "," The corrections we are proposing assume that all active stars, for which the chromospheric index $\log R'_{\rm HK}> -4.75$ "
 section. we will solve equations (51). (9)). (13)) for the Haring region. and equations (17)). (18)). (19)) for the outer region in terms of quantities which can be inferred. eit her fromfits of relativistic Low mocels to radio images (jet and MES ' in the Haring region. together with the radius of the laver in both regions) or from. X-ray observations of the surrounding hot eas (external density. temperature and pressure).," In this section, we will solve equations \ref{eq-mass-flare}) ), \ref{eq-mom-flare}) ), \ref{eq-en-flare}) ) for the flaring region, and equations \ref{eq-mass-out}) ), \ref{eq-mom-out}) ), \ref{eq-en-out}) ) for the outer region in terms of quantities which can be inferred either from fits of relativistic flow models to radio images (jet and layer velocities in the flaring region, together with the radius of the layer in both regions) or from X-ray observations of the surrounding hot gas (external density, temperature and pressure)."
" We can then derive the shape of the laminar jet. rj Gr). thevariation of velocity with distance in the outer region. J,Ge). the values of -# in the various regions. the - entrainment and the velocity of entraüinment."," We can then derive the shape of the laminar jet, $r_{j}(x)$, the variation of velocity with distance in the outer region, $\beta_{s}(x)$, the values of $\mathscr{R}$ in the various regions, the entrainment function and the velocity of entrainment."
 We assume that the laminar jet has a relativistic |- state with D; = 4/3:theenvironment. has DL. = 5/3., We assume that the laminar jet has arelativistic equation of state with $\Gamma_{j}=4/3$; the environment has $\Gamma_{e}=5/3$.
Theshear laver contains mixed material but the energy. density must still be dominated. by relativistic particles (D94). and we therefore take Py = 4/3.~ ]," The shear layer contains mixed material but the energy density must still be dominated by relativistic particles (B94), and we therefore take $\Gamma_{s}=4/3$ ."
 From equations (5)) and (13)). we have:," From equations \ref{eq-mass-flare}) ) and \ref{eq-en-flare}) ), we have:"
ave listed in Table 1. their uncertainties are the 1-0. values estimated. by (throughout this paper. the quoted uncertainties are lear. except when stated otherwise).,"are listed in Table 1, their uncertainties are the $\sigma$ values estimated by (throughout this paper, the quoted uncertainties are $\sigma$, except when stated otherwise)."
 The TOA residuals are clisplaved in the bottom plot o£ FE , The TOA residuals are displayed in the bottom plot of Fig. \ref{fig:frequencies+residuals}.
We have performed a parallel analysis using the software., We have performed a parallel analysis using the software.
package’. The resulting timing parameters. and their uncertainties (as estimated. by TEMPO) are also presented in Table 1., The resulting timing parameters and their uncertainties (as estimated by ) are also presented in Table 1.
 Several timing parameters are clillerent [rom those estimated. byrEMDPO2., Several timing parameters are different from those estimated by.
 ALL time-like quantities (spin [requency. orbital period. projected. semi-major axis. time of passage through periastron and c. which is strongly covariant with the latter) are different because converts the POAs to Barvcentric Dynamic Time (PDB). a time scale that is cillerent [rom PCB.," All time-like quantities (spin frequency, orbital period, projected semi-major axis, time of passage through periastron and $\omega$, which is strongly covariant with the latter) are different because converts the TOAs to Barycentric Dynamic Time (TDB), a time scale that is different from TCB."
 The right ascension and. declination are dillerent because the DE/LE 405 solar system ephemeris ? uses an earlier. celestial reference frame., The right ascension and declination are different because the DE/LE 405 solar system ephemeris \cite{sta98b} uses an earlier celestial reference frame.
 Furthermore. we included in a /mocified version of the DD orbital model that uses the orthometric parameterization of the Shapiro delay. as described. in ?..," Furthermore, we included in a modified version of the DD orbital model that uses the orthometric parameterization of the Shapiro delay, as described in \cite{fw10}."
" The vorthometric amplitude” (5) and the ""orthometric ratio” (4) provide an improved description. of the areas of the (mi.sini) plane where the svstem is likely to be compared to the ""range"" (r) and. “shape” (3) parameters used in the normal parameterization."," The “orthometric amplitude” $h_3$ ) and the “orthometric ratio” $\varsigma$ ) provide an improved description of the areas of the $m_c, \sin i$ ) plane where the system is likely to be compared to the “range” $r$ ) and “shape” $s$ ) parameters used in the normal parameterization."
 ln. cases where we can measure other post-Ixeplerian. (PIN). parameters. the new parameterization vields an improved test of general relativity.," In cases where we can measure other post-Keplerian (PK) parameters, the new parameterization yields an improved test of general relativity."
 Using either or there are no apparent trends in the residuals and the normalised X7 is 1.3. which is similar to what we observe in other ALSPs timed. with these observing svstems.," Using either or there are no apparent trends in the residuals and the normalised $\chi^2$ is 1.3, which is similar to what we observe in other MSPs timed with these observing systems."
 This means that the timing model provides a complete description of the observed TOXs. with no significant. unmocdelled. effects. present.," This means that the timing model provides a complete description of the observed TOAs, with no significant, unmodelled effects present."
 This also means that the 1-0 uncertainties reported. by and are reasonably accurate estimates of the real uncertainties., This also means that the $\sigma$ uncertainties reported by and are reasonably accurate estimates of the real uncertainties.
 The non-unity value of the normalised. S7. likely reflects a slight. underestimation of the TOA uncertainties by the cross-correlation analysis. but it could. also be due to red noise in the pulsar rotation. as suggested by the marginal (almost 3-0) detection of P.," The non-unity value of the normalised $\chi^2$ likely reflects a slight underestimation of the TOA uncertainties by the cross-correlation analysis, but it could also be due to red noise in the pulsar rotation, as suggested by the marginal (almost $\sigma$ ) detection of $\ddot{\nu}$."
 To account for this. the timing parameters in Table E were obtained after the multiplication of the TOA uncertainty estimates by a scale factor: this is calculated separately for each dataset so that its normalised VO is LL," To account for this, the timing parameters in Table 1 were obtained after the multiplication of the TOA uncertainty estimates by a scale factor; this is calculated separately for each dataset so that its normalised $\chi^2$ is 1."
CDhis scale factor is 1.1 for all the Arecibo data and 1.3 for the GBT data., This scale factor is 1.1 for all the Arecibo data and 1.3 for the GBT data.
 Radial velocities and spectral parameters were determined bv comparing the observed. spectra. both individual and averages at cach epoch. with synthetic spectral templates from the spectral library of Munari ct ((2005).," Radial velocities and spectral parameters were determined by comparing the observed spectra, both individual and averages at each epoch, with synthetic spectral templates from the spectral library of Munari et \nocite{mscz05}."
". οσο models range in effective temperature Zi. surface gravity g. rotational velocity Vi, and metallicity M/I] and are sampled at a resolution of A/;NÀ=20000."," These models range in effective temperature $T_\mathrm{eff}$, surface gravity $g$, rotational velocity $V_\mathrm{rot}$ and metallicity $\mathrm{[M/H]}$ and are sampled at a resolution of $\lambda/{\Delta \lambda}=20000$."
 The wavelength range between aancl wavas used., The wavelength range between and was used.
 To remove the effects of the finite resolution of he instrument. the svnthetic spectra were convolved. with a Gaussian profile with a width equal to the seeing. aud runeated at the width of the slit.," To remove the effects of the finite resolution of the instrument, the synthetic spectra were convolved with a Gaussian profile with a width equal to the seeing, and truncated at the width of the slit."
 For each. object. an iterative procedure was used to first determine the racial velocities using a starting template. then use those velocities o shift all spectra to zero velocity and create an average which was then used to obtain a better matching template or determining more accurate radial velocities.," For each object, an iterative procedure was used to first determine the radial velocities using a starting template, then use those velocities to shift all spectra to zero velocity and create an average which was then used to obtain a better matching template for determining more accurate radial velocities."
 The comparison of the average spectrum with synthetic spectra from the library of Ὁ vields a temperature of dup=5825+ POOWW (lo) and sets a 2-0 lower limit on the surface. gravity of loegdems. firmly excluding the possibility of a giant or sub-giant star (που Fig. 4)).," The comparison of the average spectrum with synthetic spectra from the library of \cite{mscz05} yields a temperature of $T_\mathrm{eff}\,=\,5825\pm200$ K $1\sigma$ ) and sets a $\sigma$ lower limit on the surface gravity of $\log g > 4 \rm\,cm\,s^{-2}$, firmly excluding the possibility of a giant or sub-giant star (see Fig. \ref{fig:tefflogg}) )."
" Based on the shape of the absorption lines we estimate a 3-7 upper limit on the rotational broadening of esin,140 where 7, is the inclination of the star."," Based on the shape of the absorption lines we estimate a $\sigma$ upper limit on the rotational broadening of $v_\mathrm{rot} \sin i_* <
140$ $^{-1}$, where $i_*$ is the inclination of the star."
 The observed spectrum does not constrain the metallicity., The observed spectrum does not constrain the metallicity.
 Stellar evolution models ?. show that a main sequence (MS) star with the observed. mass and. effective. temperature of the companion of is consistent with the observed. infrared. colours. for distances between 6 and S kpe. assuming the estimated reddening as a function of distance from 2\LASS stars ?..," Stellar evolution models \cite{gbbc00} show that a main sequence (MS) star with the observed mass and effective temperature of the companion of is consistent with the observed infrared colours for distances between 6 and 8 kpc, assuming the estimated reddening as a function of distance from 2MASS stars \cite{crl+08}."
 This is in agreement with the 6.4 kpe distance estimated from the pulsar’s DM (see Table 1) using the ? model of the Galactic electron distribution., This is in agreement with the 6.4 kpc distance estimated from the pulsar's DM (see Table 1) using the \cite{cl02} model of the Galactic electron distribution.
 The preclicted ages from. these models range from 4 to GGvr., The predicted ages from these models range from 4 to Gyr.
 Table 2 shows the racial velocities of the pulsar counterpart relative to the Solar System. barveentre., Table \ref{tab:radvel} shows the radial velocities of the pulsar counterpart relative to the Solar System barycentre.
 They are 92.4dAkknm ss during the first epoch. and 15= 2kkmss+ curing the second.," They are $V_1\,=\,92.4\pm8.3$ $^{-1}$ during the first epoch, and $V_2\,=\,20.2\pm5.2$ $^{-1}$ during the second."
 The velocity dillerence between the two epochs of VV3=72.2£9 s8kkmss deviates from constant velocity at the 7.460. level.," The velocity difference between the two epochs of $V_1-V_2\,=\,72.2\pm9.8$ $^{-1}$ deviates from constant velocity at the $7.4\sigma$ level."
 Between the two epochs. the averaged radial velocities of stars A and C diller bv Vy15=142+0.3kkm ss.+ for star A," Between the two epochs, the averaged radial velocities of stars A and C differ by $V_1-V_2=-14.2\pm0.3$ $^{-1}$ for star A"
dispersion relation.,dispersion relation.
" This work was partially supported by PRIN-MIUR 2006, by ASI through contract ASI-INAF 1/088/06/0 and (for PB) by the US DOE and by NASA grant NAG5-10842."," This work was partially supported by PRIN-MIUR 2006, by ASI through contract ASI-INAF I/088/06/0 and (for PB) by the US DOE and by NASA grant NAG5-10842."
" Fermilab is operated by Fermi Research Alliance, LLC under Contract No."," Fermilab is operated by Fermi Research Alliance, LLC under Contract No."
 DE-ACO2-07CHI1359 with the United States DOE., DE-AC02-07CH11359 with the United States DOE.
componnents of planetary systems. such as comets in our solar system.,"nents of planetary systems, such as comets in our solar system."
non-dusty galaxies.,non-dusty galaxies.
 The MIDS catalogue of in sources. though. is selected. on cust emission to high. redshift. and there is a known correlation between the lux densities at mic-infrared: ancl far-infrarecl wavelengths (Charv 2001).," The MIPS catalogue of $\,$$\micron$ sources, though, is selected on dust emission to high redshift, and there is a known correlation between the flux densities at mid-infrared and far-infrared wavelengths \citep{chary01}."
. The stacking equation we use is similar to Iq. 1::, The stacking equation we use is similar to Eq. \ref{eqn:coadd}:
" where Shin is the stacked [lux density of Nig, sources in à bin of ini Hux density or redshift. and;46 and Ait. are the measured. 1.16 mim Dux density ancl noise at the position of the ;-th 24jai source."," where $S_{\mathrm{bin}}$ is the stacked flux density of $N_{\mathrm{bin}}$ sources in a bin of $\,$$\micron$ flux density or redshift, and $S_{\mathrm{i},1.16}$ and $\sigma_{\mathrm{i},1.16}$ are the measured $\,$ mm flux density and noise at the position of the $i$ -th $\,$$\micron$ source."
 Vhis equation does not include any constant terms (es) because the goal of stacking is to ect an average [ux density for all sources from a map at one wavelength., This equation does not include any constant terms $w$ 's) because the goal of stacking is to get an average flux density for all sources from a map at one wavelength.
 The noise decreases with the inclusion of DUOFC SOULCES: In a mathematical sense. this equation is only valid when all of the σιτο are independent: because there are many pm sources in the area of one PSE. this requirement is strictlv not met.," The noise decreases with the inclusion of more sources: In a mathematical sense, this equation is only valid when all of the $\sigma_{\mathrm{i},1.16}$ are independent; because there are many $\,$$\micron$ sources in the area of one PSF, this requirement is strictly not met."
 We fit a Gaussian to the distribution of stacked. [lux densities for. 2484 random: pixels. and the σ ," We fit a Gaussian to the distribution of stacked flux densities for 2484 random pixels, and the $\sigma$ "
This is a noteworthy result.,This is a noteworthy result.
 It means (that. independent of the strength: of radiative cooling and hence of plasma compression. a finite fraction of the incoming magnetic energy is not dissipated into heat. (which is then promptly radiated away). but instead is converted directly into the mechanical kinetic energy of the outflow.," It means that, independent of the strength of radiative cooling and hence of plasma compression, a finite fraction of the incoming magnetic energy is not dissipated into heat (which is then promptly radiated away), but instead is converted directly into the mechanical kinetic energy of the outflow."
 This implies that. strictly speaking. one cannot use the total energy. conservation to estimate the compression ratio ;1 bv simply equating the Poynting flux with the radiative losses!," This implies that, strictly speaking, one cannot use the total energy conservation to estimate the compression ratio $A$ by simply equating the Poynting flux with the radiative losses!"
 What one should do instead. is to look al the [ate of plasmaentropy. i.e. al the balance between the plasma heating due to the dissipation of magnetic enerev (the ohmic heating) and the plasma cooling both via the advection of heat out of the laver and via the radiation losses.," What one should do instead, is to look at the fate of plasma, i.e., at the balance between the plasma heating due to the dissipation of magnetic energy (the ohmic heating) and the plasma cooling both via the advection of heat out of the layer and via the radiation losses."
 In other words. one needs to focus on the thermal content of the plasma. as opposed to its total energv.," In other words, one needs to focus on the thermal content of the plasma, as opposed to its total energy."
 Then one can sav that the svstem is in the strong radiative cooling regime if the olmic heating rate is primarily balanced by radiative losses. whereas the heat losses by advection are small.," Then one can say that the system is in the strong radiative cooling regime if the ohmic heating rate is primarily balanced by radiative losses, whereas the heat losses by advection are small."
 This is in contrast with the classical (non-radiative) SweetParker model. where it is forward to show that ohmic heating is balanced by the adyection of thermal energy out of the laver.," This is in contrast with the classical (non-radiative) Sweet–Parker model, where it is straight-forward to show that ohmic heating is balanced by the advection of thermal energy out of the layer."
 The ohmic heatinge rate in our moclel can be estimated as NEN— qED a)\\2 0— SVG where ij=4x/c is the resistivity. and where we made use of equation (2.2)).," The ohmic heating rate in our model can be estimated as = j^2 )^2 = , where $\eta' = 4\pi\,\eta/c^2$ is the resistivity, and where we made use of equation \ref{eq-Ohm}) )."
 Combining this with the mass conservation condition (2.2)). eL~udel. and with (2.2)). we find Cu. where we defined acharacteristic power per unit volume Qy=Bz /(dar4u).," Combining this with the mass conservation condition \ref{eq-mass-conserv}) ), $v_{\rm rec} L \sim u\delta A$, and with \ref{eq-u}) ), we find Q_0, where we defined acharacteristic power per unit volume $Q_0\equiv B_0^2/(4\pi\tau_{A0})$ ."
for this data set is aand the Strehl ratio is 0.34.,for this data set is and the Strehl ratio is 0.34.
 The initial image of rrevealed an interesting faint nearby source which. based on a comparison with observations made by Kouwenhovenetal.(2005. 2007).. was very likely to be à true co- companion.," The initial image of revealed an interesting faint nearby source which, based on a comparison with observations made by \citet{kouwenhoven05, kouwenhoven07}, was very likely to be a true co-moving companion."
 To confirm the common proper motion of this candidate and. get additional photometry measurements. follow-up imaging observations in J. H and K’ were obtained on 2009 July 2 using the same instrument as before.," To confirm the common proper motion of this candidate and get additional photometry measurements, follow-up imaging observations in $J$, $H$ and $K^\prime$ were obtained on 2009 July 2 using the same instrument as before."
 To avoid saturation. m these broad filters. we used a 512«$12 subarray. and kept the strategy of a short. unsaturated image immediately followed by a longer. saturated image at each of 5 dither positions.," To avoid saturation in these broad filters, we used a $512\times512$ subarray, and kept the strategy of a short, unsaturated image immediately followed by a longer, saturated image at each of 5 dither positions."
" For each wavelength. three sky frames were obtained at offsets of >15""."," For each wavelength, three sky frames were obtained at offsets of $>15\arcsec$."
 The short exposure times were 0.055 s with 40 co-additions. while the long exposures were 6 s. 5 s. and 4 s with one co-addition in J. H and K'. respectively.," The short exposure times were 0.055 s with 40 co-additions, while the long exposures were 6 s, 5 s, and 4 s with one co-addition in $J$, $H$ and $K^\prime$, respectively."
 To increase the significance of our common proper motion confirmation further. another follow-up observation was made on 2010 August 30 using the filter. full frame. and one co-addition of twelve 0.6 s Απintegrations followed by a single 10 s integration. at each of five dither positions.," To increase the significance of our common proper motion confirmation further, another follow-up observation was made on 2010 August 30 using the $K_{\rm cont}^{2.09}$ filter, full frame, and one co-addition of twelve 0.6 s integrations followed by a single 10 s integration at each of five dither positions."
 Other follow-up data were acquired in spring and summer 2010 but owing to an overseen difference in the instrumental setup. these data suffer from large systematic astrometric errors and are not used here.," Other follow-up data were acquired in spring and summer 2010 but owing to an overseen difference in the instrumental setup, these data suffer from large systematic astrometric errors and are not used here."
 The imaging data were reduced using custom routines., The imaging data were reduced using custom routines.
 For the 2008 and 2010 imaging data. a sky frame was constructed by taking the median of the images at all dither positions after masking out the regions dominated by the target’s signal. while for the 2009 imaging data. the median of the three sky frames was obtained.," For the 2008 and 2010 imaging data, a sky frame was constructed by taking the median of the images at all dither positions after masking out the regions dominated by the target's signal, while for the 2009 imaging data, the median of the three sky frames was obtained."
 After subtraction of this sky frame. the images were divided by a normalized flat-field.," After subtraction of this sky frame, the images were divided by a normalized flat-field."
 Then isolated bad pixels were replaced by the interpolated value of a third-order polynomial surface fit to the good pixels in à 7« pixels box while clustered bad pixels were simply masked out., Then isolated bad pixels were replaced by the interpolated value of a third-order polynomial surface fit to the good pixels in a $7\times7$ pixels box while clustered bad pixels were simply masked out.
 Next the images were distortion corrected using the distortion solution provided by the Gemini Finally. the long-exposure images were properly scaled in intensity and merged with their corresponding short-exposure images.," Next the images were distortion corrected using the distortion solution provided by the Gemini Finally, the long-exposure images were properly scaled in intensity and merged with their corresponding short-exposure images."
 The color composite /HK image from the 2009 follow-up observations is shown imn figure I.., The color composite $JHK$ image from the 2009 follow-up observations is shown in figure \ref{fig:im}.
 The likely companion detected in the imaging data was very faint and. assuming if was a true companion. its magnitude suggested a substellar mass.," The likely companion detected in the imaging data was very faint and, assuming it was a true companion, its magnitude suggested a substellar mass."
 Thus. to verify its late-type nature. we obtained spectroscopic observations of it in A. K and J on 2009 July 2. 2009 July 3. and 2009 August 8. respectively. using the integral field spectrograph NIFS (McGregoretal.2003). with the ALTAIR adaptive optics system at the Gemini North telescope.," Thus, to verify its late-type nature, we obtained spectroscopic observations of it in $H$, $K$ and $J$ on 2009 July 2, 2009 July 3, and 2009 August 8, respectively, using the integral field spectrograph NIFS \citep{mcgregor03} with the ALTAIR adaptive optics system at the Gemini North telescope."
 The / grating was used with the ZJ blocking filter. the H grating with the /H blocking filter and the K grating with the ΗΝ blocking filter.," The $J$ grating was used with the $ZJ$ blocking filter, the $H$ grating with the $JH$ blocking filter and the $K$ grating with the $HK$ blocking filter."
 The spectral resolving power ts ~6000 in J and ~5300 in both H and K., The spectral resolving power is $\sim$ 6000 in $J$ and $\sim$ 5300 in both $H$ and $K$.
" Given the large angular separation of the compantor (~4.5"")). the primary star is not visible in the 3""«3” NIFS field-of-view."," Given the large angular separation of the companion $\sim$ ), the primary star is not visible in the $3\arcsec\times 3\arcsec$ NIFS field-of-view."
 In each band. the companion was positionec near the center of the field-of-view and was dithered by aalong a line between each of five exposures to enable gooc sky subtraction.," In each band, the companion was positioned near the center of the field-of-view and was dithered by along a line between each of five exposures to enable good sky subtraction."
 The individual exposure times were 480 s. 300 s and 240 s in J. H and K. respectively. in low read noise mode.," The individual exposure times were 480 s, 300 s and 240 s in $J$, $H$ and $K$, respectively, in low read noise mode."
 After the J and H sequences. the AOV star HIP 79229 was observed at a similar airmass for telluric and instrumental transmission correction. while the AOV star HIP 73820 was observed just before the K band sequence for the same purpose.," After the $J$ and $H$ sequences, the A0V star HIP 79229 was observed at a similar airmass for telluric and instrumental transmission correction, while the A0V star HIP 73820 was observed just before the $K$ band sequence for the same purpose."
 The reduction of the NIFS data. up to the reconstructior of the data cube. was made using the Gemini IRAF pipeline.," The reduction of the NIFS data, up to the reconstruction of the data cube, was made using the Gemini IRAF pipeline."
 The steps covered in this pipeline are sky background subtraction. flat-field and bad pixel correction. spatial and spectral calibration. and data cube reconstruction.," The steps covered in this pipeline are sky background subtraction, flat-field and bad pixel correction, spatial and spectral calibration, and data cube reconstruction."
 The spatial sampling of the reconstructed data cube is ppixel!., The spatial sampling of the reconstructed data cube is $^{-1}$.
 The instrumental/telluric transmission. correctior and spectrum extraction. which. could be done using the Gemini pipeline. were instead done using custom IDL.," The instrumental/telluric transmission correction and spectrum extraction, which could be done using the Gemini pipeline, were instead done using custom IDL."
 The center of each PSF was first registered to à commor position in all spectral slices and all cubes of the sequence: the center positions were calculated by fitting a 2D Gaussiar function., The center of each PSF was first registered to a common position in all spectral slices and all cubes of the sequence; the center positions were calculated by fitting a 2D Gaussian function.
 The spectrum of the source was extracted by summing the flux in a circular aperture of diameter 5 pixels., The spectrum of the source was extracted by summing the flux in a circular aperture of diameter 5 pixels.
 The spectrum of the telluric standard was corrected for its spectral slope using a 9520 K blackbody curve. and any hydrogen line absorption was removed by dividing out a Voigt profile fit.," The spectrum of the telluric standard was corrected for its spectral slope using a 9520 K blackbody curve, and any hydrogen line absorption was removed by dividing out a Voigt profile fit."
 The companion spectrum was divided by the standard spectrum to correct for telluric and instrumental transmission., The companion spectrum was divided by the standard spectrum to correct for telluric and instrumental transmission.
 Finally. the median of the 5 spectra was obtained.," Finally, the median of the 5 spectra was obtained."
 A summary of the properties of 78530A and its companion. taken from the literature or derived in this section. is presented in Table 2..," A summary of the properties of A and its companion, taken from the literature or derived in this section, is presented in Table \ref{tbl:properties}."
 The relative position of the companion and primary was determined by fitting a 2D Gaussian function to each component., The relative position of the companion and primary was determined by fitting a 2D Gaussian function to each component.
 The pixel positions were converted to areseconds using the 21.4 mas pixel scale of NIRI with ALTAIR and the field lens as indicated on the instrument webpage?:: the position angle of the image was taken from the FITS header., The pixel positions were converted to arcseconds using the 21.4 mas pixel scale of NIRI with ALTAIR and the field lens as indicated on the instrument web; the position angle of the image was taken from the FITS header.
 The astrometric errors. 6 mas in separation and ün position angle. were estimated in Lafreniéreetal.(2010) using the position of background stars for a similar sequence of observations: they are mostly dominated by residual distortion errors.," The astrometric errors, 6 mas in separation and in position angle, were estimated in \citet{lafreniere10} using the position of background stars for a similar sequence of observations; they are mostly dominated by residual distortion errors."
 We also note that the pixel scale of ALTAIR/NIRI with the field lens has not been extensively calibrated and that our values may be systematically different from measurements made using other instruments or telescopes. but they should be internally consistent.," We also note that the pixel scale of ALTAIR/NIRI with the field lens has not been extensively calibrated and that our values may be systematically different from measurements made using other instruments or telescopes, but they should be internally consistent."
 The, The
stars. but the linear resolution implied by the accretion radii (0.05pc) is rather large.,"stars, but the linear resolution implied by the accretion radii (0.05pc) is rather large."
 Our sink particles should strictly be regarded: as binaries or small multiple systems so that our mass functions. in the terminologv of. 2.. are Alultiple Star Mass. Functions CMSMES). not Single Star Alass Funetions (SSMES).," Our sink particles should strictly be regarded as binaries or small multiple systems so that our mass functions, in the terminology of \cite{2008A&A...477..823G}, are Multiple Star Mass Functions (MSMFs), not Single Star Mass Functions (SSMFs)."
 This would be true to an even &reater degree if our results were rescaled to higher masses., This would be true to an even greater degree if our results were rescaled to higher masses.
 7? lind. not surprisingly. that the SSME contains fewer ligh-mass objects anc more lowmass objects relative to he MSME. and that the SSME peaks at lower masses.," \cite{2008A&A...477..823G} find, not surprisingly, that the SSMF contains fewer high-mass objects and more low–mass objects relative to the MSMF, and that the SSMF peaks at lower masses."
 The SSME from our simulations would. probably therefore. be slightly less deficient in lowmass objects and contain fewer ighmass objects than the mass functions we present. but he cllects of multiplicity will not quantitatively change our inding that the mass functions become very topheavy due o oligarchic 7. observed the mass function of CO clumps in the Carina Flare supershell anc observed that. their mass spectrum follows a powerlaw.," The SSMF from our simulations would probably therefore be slightly less deficient in low–mass objects and contain fewer high–mass objects than the mass functions we present, but the effects of multiplicity will not quantitatively change our finding that the mass functions become very top–heavy due to oligarchic \cite{2008PASJ...60.1297D} observed the mass function of CO clumps in the Carina Flare supershell and observed that their mass spectrum follows a power–law."
 Although they caution against overinterpretation of their data. their clump mass measurements imply that the clumps are not eravitationallv xound (in the sense that their inferred masses are. below heir virial masses).," Although they caution against over–interpretation of their data, their clump mass measurements imply that the clumps are not gravitationally bound (in the sense that their inferred masses are below their virial masses)."
" This result is consistent with our indings that the asvetunbound objects produced in the inear phase of shell fragmentation have a power law mass ""unction (although note that the power law we derive is considerably steeper than that inferred by ?)).", This result is consistent with our findings that the as–yet–unbound objects produced in the linear phase of shell fragmentation have a power law mass function (although note that the power law we derive is considerably steeper than that inferred by \citet{2008PASJ...60.1297D}) ).
" Our findings suggest. however. that the mass function of these unbound ragments cannot necessarily be used to predict. the mass ""unction of stars or star clusters that will eventually form."," Our findings suggest, however, that the mass function of these unbound fragments cannot necessarily be used to predict the mass function of stars or star clusters that will eventually form."
 Vhere are as vet no observations of the stellar or cluster mass functions generated by the fragmentation of expanding shells. owing to the extreme cilliculty of making such a The initial conditions used in our simulations. that of a shell which is pressureconfined anc vet does not sweep up any mass are somewhat unusual and the peculiar mass function we obtain implies that most stars or clusters clo not form in shells of this tvpe.," There are as yet no observations of the stellar or cluster mass functions generated by the fragmentation of expanding shells, owing to the extreme difficulty of making such a The initial conditions used in our simulations, that of a shell which is pressure–confined and yet does not sweep up any mass are somewhat unusual and the peculiar mass function we obtain implies that most stars or clusters do not form in shells of this type."
 However. such shells must surely occur sometimes. since a shell may expand. into the hot rarefied LSAT and be pressureconfined. while accreting negligible mass.," However, such shells must surely occur sometimes, since a shell may expand into the hot rarefied ISM and be pressure–confined while accreting negligible mass."
 Our results demonstrate that shells of this nature will produce anomalously high. numbers of massive stars or clusters and may therefore lead to the propagation of star formation by triggering., Our results demonstrate that shells of this nature will produce anomalously high numbers of massive stars or clusters and may therefore lead to the propagation of star formation by triggering.
 We find that. in the case of a momentumdriven shell pressureconfined. internally anc externally. the PACL moclel faithfully reproduces the mass function of fragments produced. by the gravitational instability during the time for which perturbations in the shell surface. density are relatively small and the behaviour is In the case studied here of a shell which does. not sweep up mass. we find that a form of accretion which we term oligarchic accretion operates within the shell wherchy the first objects to form become the most massive because they aecrete from the gas. reservoir. for longer.," We find that, in the case of a momentum–driven shell pressure–confined internally and externally, the PAGI model faithfully reproduces the mass function of fragments produced by the gravitational instability during the time for which perturbations in the shell surface density are relatively small and the behaviour is In the case studied here of a shell which does not sweep up mass, we find that a form of accretion which we term oligarchic accretion operates within the shell whereby the first objects to form become the most massive because they accrete from the gas reservoir for longer."
 Coupled with the finite number of p.cores formed during the linear fragmentation of the shell. this leads to a very topheavy mass function in which the most unstable wavelength ab the time when perturbations first. become bound. is overrepresented.," Coupled with the finite number of p–cores formed during the linear fragmentation of the shell, this leads to a very top–heavy mass function in which the most unstable wavelength at the time when perturbations first become bound is over–represented."
 A shell whose gas reservoir is replenished may form [fragments continuously. resulting in a mass function more closely. resembling a power Since the mass function generated by our simulations is so unusual. we infer that shells of the type we have studied cannot make a strong contribution to the global stellar population. although no direct. measurements of the mass functions produced by expanding shells exist against which we could test our conclusions.," A shell whose gas reservoir is replenished may form fragments continuously, resulting in a mass function more closely resembling a power Since the mass function generated by our simulations is so unusual, we infer that shells of the type we have studied cannot make a strong contribution to the global stellar population, although no direct measurements of the mass functions produced by expanding shells exist against which we could test our conclusions."
 Vhe authors thank the. referee. sven Van Loo. for incisive and helpful comments which considerably improvec the structure of the paper.," The authors thank the referee, Sven Van Loo, for incisive and helpful comments which considerably improved the structure of the paper."
 JED acknowledges: suppor from a Marie. Curie. fellowship as part of the European Commission FPG Rescarch Training Network ‘Constellation’ under contract NICEN.CP2006035890., JED acknowledges support from a Marie Curie fellowship as part of the European Commission FP6 Research Training Network `Constellation' under contract MRTN–CT–2006–035890.
 JIZD.. RAV anc JP acknowledge support from the Institutional. Research Plan AVOZLO0030501 of the Academy of Seicnces of the Czech Republic and project LCO6014.Centre for Theoretica Astrophysics of the Ministry of Education. Youth and Sports of the Czech Republic.," JED, RW and JP acknowledge support from the Institutional Research Plan AV0Z10030501 of the Academy of Sciences of the Czech Republic and project LC06014–Centre for Theoretical Astrophysics of the Ministry of Education, Youth and Sports of the Czech Republic."
 AW gratefully acknowledges the support of erant. ST/11001530/1 from the Ulx Science anc Technology. Facilities Council., AW gratefully acknowledges the support of grant ST/H001530/1 from the UK Science and Technology Facilities Council.
 The thin shell dispersion relation may be recast in a dimensionless form with the help of an expansion Law a prescription for the variations in the shell racius 2. expansion velocity V. and surface density X with time.," The thin shell dispersion relation may be recast in a dimensionless form with the help of an expansion law – a prescription for the variations in the shell radius $R$, expansion velocity $V$ and surface density $\Sigma$ with time."
" The shells described in this work are ballistic and of fixed mass M. expanding due to their inertia into a vacuum and decelerated by self.gravity. We celine a dimensionless radius £ as ∖∖⋎↓↕⋖⋅↓⋅∢⊾∫∖⋡∶⇀∪⇖⋡⇉⇉⋜⋯∠⇂⊓⊽∶↿∖∩∆∪⊒∃↿∖∐∩⋜⋯⋅⇂↓⊔⊾ kinetic and. potential energy respectively. and is the maximum radius to which the shell expands (the initial total cnerev Au|Wy at radius Z2, is assumed to be negative so that the shell is bound)."," The shells described in this work are ballistic and of fixed mass $M$, expanding due to their inertia into a vacuum and decelerated by self–gravity We define a dimensionless radius $\xi$ as where $K = MV^2/2$ and $W = (GM^2)/(2R)$ are the kinetic and potential energy respectively, and is the maximum radius to which the shell expands (the initial total energy $K_0 + W_0$ at radius $R_0$ is assumed to be negative so that the shell is bound)."
 The free fall time obtained from Ixepler's third law is, The free fall time obtained from Kepler's third law is
Quasars were discovered in the 1960s by virtue of their radio emission.,Quasars were discovered in the 1960s by virtue of their radio emission.
 While this emission later proved to be exception rather than the rule. the radio sources associated with quasi-stellar objects have found great interest over the past decades.," While this emission later proved to be exception rather than the rule, the radio sources associated with quasi-stellar objects have found great interest over the past decades."
 With ever increasing angular resolution (including. VLBI). the study of extragalactic radio sources — quasars and radio galaxies — led to our understanding of the relevant processes. whereby the ultimate energy source in these objects must be located in their nuclei. on the subparsee scale. and whereby jets transport some part of this energy to radio-emitting lobes in intra- or intergalactic space.," With ever increasing angular resolution (including VLBI), the study of extragalactic radio sources – quasars and radio galaxies – led to our understanding of the relevant processes, whereby the ultimate energy source in these objects must be located in their nuclei, on the subparsec scale, and whereby jets transport some part of this energy to radio-emitting lobes in intra- or intergalactic space."
 In fact. the study of extragalactic radio sources and the subsequent need for increasing angular resolution provided the stimulus to develop VLBI techniques.," In fact, the study of extragalactic radio sources and the subsequent need for increasing angular resolution provided the stimulus to develop VLBI techniques."
 Fine reviews of the morphologies and other properties of extragalactic radio sources can be found in e.g. ?.. ?.. ?.. 2.. and ?..," Fine reviews of the morphologies and other properties of extragalactic radio sources can be found in e.g., \citet{Miley1980}, , \citet{Bridle1984}, , \citet{Saikia1988}, \citet{Carilli1996}, and \citet{Zensus1997}."
 The radio morphologies of quasars can be broadly classified into two categories: core-dominated or lobe-dominated., The radio morphologies of quasars can be broadly classified into two categories: core-dominated or lobe-dominated.
 Examination of the detailed properties of these two classes. incorporating properties in other wavelength regimes. has led to the notion that the orientation of the radio jets with respect to the observer is the prime discriminant.," Examination of the detailed properties of these two classes, incorporating properties in other wavelength regimes, has led to the notion that the orientation of the radio jets with respect to the observer is the prime discriminant."
 Core-dominated objects typically display one short (curved) jet and relatively weak halo-emission besides the dominant core (coincident with the actual QSO). whereas the morphologies of lobe-dominated (1.e.. Fanaroff Riley Class 2. FR2) objects comprise extended double lobes straddling the QSO radio core.," Core-dominated objects typically display one short (curved) jet and relatively weak halo-emission besides the dominant core (coincident with the actual QSO), whereas the morphologies of lobe-dominated (i.e., Fanaroff Riley Class 2, FR2) objects comprise extended double lobes straddling the QSO radio core."
 Only single-sided jets are observed. in either classes.," Only single-sided jets are observed, in either classes."
" This jet asymmetry together with the occurrence and the magnitude of the nilliaresee scale proper motion measured | the jet components led to the picture whereby radio-loud quasars are preferably oriented objects. the lobe-dominated objects being at substantial angles to the line-of-sight. the core-""Sominated objects being at small angles to that line."," This jet asymmetry together with the occurrence and the magnitude of the milliarcsec scale proper motion measured in the jet components led to the picture whereby radio-loud quasars are preferably oriented objects, the lobe-dominated objects being at substantial angles to the line-of-sight, the core-dominated objects being at small angles to that line."
 Projection. foreshortening. and relativistic effects are naturally invoked to explain the observed radio properties of the compact subclass. icluding the frequently observed superluminal motion (e.g.. ?)..," Projection, foreshortening, and relativistic effects are naturally invoked to explain the observed radio properties of the compact subclass, including the frequently observed superluminal motion \citep[e.g.,][]{Orr1982}."
 The combined properties of the radio-loud QSO class together. however. imply that quasars 1n or close to the plane of the sky simply do not exist: such objects must masquerade as another class of object. namely extended double-lobed FR2 radio galaxies.," The combined properties of the radio-loud QSO class together, however, imply that quasars in or close to the plane of the sky simply do not exist: such objects must masquerade as another class of object, namely extended double-lobed FR2 radio galaxies."
 Following the spectropolarimetric detection of an obscured quasar in radio galaxy 22344 (?).. this unified picture for powerful extragalacticradio sources was initially explored by ? and ?.. and subsequently put on à firm basis by ?..," Following the spectropolarimetric detection of an obscured quasar in radio galaxy 234 \citep{Ski1984}, this unified picture for powerful extragalacticradio sources was initially explored by \citet{Scheuer1987} and \citet{Peacock1987}, and subsequently put on a firm basis by \citet{Barthel1989b}."
 An optically opaque ‘torus’ surrounding the central accretion disk in its equatorial plane. blocking the view towards the accretion disk and its associated broad line region. is an essential ingredient of these unification models.," An optically opaque `torus' surrounding the central accretion disk in its equatorial plane, blocking the view towards the accretion disk and its associated broad line region, is an essential ingredient of these unification models."
 Relevant reviews of the quasar — radio galaxy as well as complementary unificatior theories can be found in ?.. ?.. and ?..," Relevant reviews of the quasar – radio galaxy as well as complementary unification theories can be found in \citet{Antonucci1993}, \citet{Urry1995}, and \citet{Maiolino2002}."
 We stress that a complete explanation of extragalactic radio sources should obviously also take into account the effects of radio source evolution (youth-adulthood-sentority). radio source environment. and activity duty cycle. besides the aspect-dependent effects.," We stress that a complete explanation of extragalactic radio sources should obviously also take into account the effects of radio source evolution (youth-adulthood-seniority), radio source environment, and activity duty cycle, besides the aspect-dependent effects."
" Their asymmetric radio depolarization (""depolarizatioi asymmetry"") provided independent support for the preferred orientation of the quasar class.", Their asymmetric radio depolarization (“depolarization asymmetry”) provided independent support for the preferred orientation of the quasar class.
 Following up on radio galaxy depolarization studies by e.g.. 2.. ? and ?. drew first attention to the fact that the “jetted” lobes in quasars display less depolarization (with increasing wavelength) than the unjettec lobes.," Following up on radio galaxy depolarization studies by e.g., \citet{Strom1988}, \citet{Laing1988} and \citet{Garrington1988} drew first attention to the fact that the “jetted” lobes in quasars display less depolarization (with increasing wavelength) than the unjetted lobes."
 They suggested that this effect was due to natural line-of-sight depolarization effects towards a near lobe with approaching jet and a distant lobe with receding Jet. through αἱ extended. Faraday-thick. magneto-ionie halo hosting the radio source.," They suggested that this effect was due to natural line-of-sight depolarization effects towards a near lobe with approaching jet and a distant lobe with receding jet, through an extended, Faraday-thick, magneto-ionic halo hosting the radio source."
The effect appeared to be strong in small radio sources. but virtually absent in large ones (?)..,"The effect appeared to be strong in small radio sources, but virtually absent in large ones \citep{Garrington1991b}."
 Extensive radio imaging studies by e.g.. 2.. 2.. 2.. 22. ?? and most recently ? have yielded broad agreement with the orientation-dependent unification scheme. but the abovementioned studies also found evidence for radio source environmental effects.," Extensive radio imaging studies by e.g., \citet{Hough1989}, \citet{Hough1999}, \citet{Bridle1994}, \citet{Fernini1993, Fernini1997}, \citet{Dennett1997, Dennett1999} and most recently \citet{Mullin2008}
 have yielded broad agreement with the orientation-dependent unification scheme, but the abovementioned studies also found evidence for radio source environmental effects."
 Consistency with aspect-dependent unification was also found in optical and infrared studies of the AGN and their host galaxies (e.g..222?)..," Consistency with aspect-dependent unification was also found in optical and infrared studies of the AGN and their host galaxies \citep[e.g.,][]{Ogle1997, Dunlop2003, Haas2004,
Haas2005}."
 Open issues obviously remain. such as thetorusopening angle (e.g.. ?).. the nature of the class of Broad Line Radio Galaxies (e.g.. 22).. and the nature of Low Excitation Radio Galaxies. LERGs as well as mid-infrared weak radio galaxies (e.g..2?) ," Open issues obviously remain, such as thetorusopening angle \citep[e.g.,][]{Grimes2004}, , the nature of the class of Broad Line Radio Galaxies \citep[e.g.,][]{Dennett2000, Bemmel2001}, , and the nature of Low Excitation Radio Galaxies, LERGs \citep[e.g.,][]{Hardcastle2004}
 as well as mid-infrared weak radio galaxies \citep[e.g.,][]{Ogle2006,
Wolk2010}
 "
the identification of the X-ray sources in paper 2 (the CERS 3 and 14-h fields).,the identification of the X-ray sources in paper 2 (the CFRS 3 and 14-h fields).
 The CEDE data were taken with the using the ULSI mosaic camera in D. V and £2 with C data supplied by either the οΓιο (3-h field) or the IKPNO (14-h field).," The CFDF data were taken with the using the UH8K mosaic camera in $B$ , $V$ and $I$, with $U$ data supplied by either the CTIO (3-h field) or the KPNO (14-h field)."
 Total exposure times were typically ~5 hours for 2. V and { and ~10 hours for C.," Total exposure times were typically $\sim5$ hours for $B$ , $V$ and $I$, and $\sim10$ hours for $U$."
 The lengthy data reduction process is described in detail in (?).., The lengthy data reduction process is described in detail in \citep{2001A&A...376..756M}.
 The selection of the AGN sample requires careful consideration. in order to avoid uncertainties leading [rom degeneracies in redshift ancl X-ray luminosity. for example.," The selection of the AGN sample requires careful consideration, in order to avoid uncertainties leading from degeneracies in redshift and X-ray luminosity, for example."
" In any flux Πίος, survey. such as our surveys. an inevitable correlation arises between redshift and. X-ray uminositv."," In any flux limited survey, such as our surveys, an inevitable correlation arises between redshift and X-ray luminosity."
 TFherefore. if we wish to study. for example. xotential correlations between clustering amplitude ancl X-ray Iuminositv. or between clustering amplitude anc redshift. hen we must select a sample accordingly so that the trends associated with one ellect are not confused with those caused ον the other.," Therefore, if we wish to study, for example, potential correlations between clustering amplitude and X-ray luminosity, or between clustering amplitude and redshift, then we must select a sample accordingly so that the trends associated with one effect are not confused with those caused by the other."
 ‘To be able to reduce the error introduced by uncertain xickeround: and foreground. number counts. as well as reducing uncertainties in X-ray luminosity. we require the yest possible. photometric redshift estimates for both the AGN and the surrounding Ποιά galaxies.," To be able to reduce the error introduced by uncertain background and foreground number counts, as well as reducing uncertainties in X-ray luminosity, we require the best possible photometric redshift estimates for both the AGN and the surrounding field galaxies."
 Therefore. we restrict this analysis to only the 1d4-h field. which has more accurate photometric redshifts than the 3-h. field.," Therefore, we restrict this analysis to only the 14-h field, which has more accurate photometric redshifts than the 3-h field."
 In this work we use photometric redshifts derived. from a slightly improved. version of DPZ (?).. compared to that used in paper 2: this version allows the photometric zero-points of the galaxy catalogue to be adjusted. leading to an improvement in the accuracy of the photometric redshifts.," In this work we use photometric redshifts derived from a slightly improved version of BPZ \citep{2000ApJ...536..571B}, compared to that used in paper 2; this version allows the photometric zero-points of the galaxy catalogue to be adjusted, leading to an improvement in the accuracy of the photometric redshifts."
 After fitting a Saggalaxy template1 and redshift to cach input1 galaxy. the code then compares the input colours with those of the fitted. templates. for cach galaxy.," After fitting a galaxy template and redshift to each input galaxy, the code then compares the input colours with those of the fitted templates for each galaxy."
 Any systematic dillerence between the input and template colours. for the whole catalogue. may indicate a systematic. photometry error in the input catalogue. which can be accounted. for before re-running the code by applving a global photometric zero-point olfset.," Any systematic difference between the input and template colours, for the whole catalogue, may indicate a systematic photometry error in the input catalogue, which can be accounted for before re-running the code by applying a global photometric zero-point offset."
 After several iterations the number of ealaxies with reliable redshifts is increased. with far lower catastrophic redshift errors.," After several iterations the number of galaxies with reliable redshifts is increased, with far fewer catastrophic redshift errors."
 Adjusting. photometric zero-points ereathy improves the aecuracy ancl reliability. of the photometric redshifts but we can eo one step further., Adjusting photometric zero-points greatly improves the accuracy and reliability of the photometric redshifts but we can go one step further.
 To ensure that we only consider galaxies with good redshifts we construc a reduced CEDE catalogue containing only those galaxies with a high reliability. measure. 24.0.9 (2)...," To ensure that we only consider galaxies with good redshifts we construct a reduced CFDF catalogue containing only those galaxies with a high reliability measure, $P_{\Delta z}>0.9$ \citep{2000ApJ...536..571B}."
 Driellv. his quantity represents. the peakedness. of the. redshif »obability function for a particular galaxy template fit.," Briefly, this quantity represents the peakedness of the redshift probability function for a particular galaxy template fit."
 Lt is he integration of the probability function around the best fi redshift out to limits defined in the code as As=0.2.(1]2).," It is the integration of the probability function around the best fit redshift out to limits defined in the code as $\Delta z =
0.2\times(1+z)$."
 If the probability. function has a well defined. peak. tha ies entirely within the integration range. then Z2.=1.0. whereas a function that is very broad. or multi-mocdoed. wil iive a low Py.," If the probability function has a well defined peak, that lies entirely within the integration range, then $P_{\Delta
z}=1.0$, whereas a function that is very broad, or multi-moded, will have a low $P_{\Delta z}$."
 Therefore. Py. gives a measure of the reliability of a particular photometric redshift estimate.," Therefore, $P_{\Delta z}$ gives a measure of the reliability of a particular photometric redshift estimate."
 Using the above eriterion further reduces the number of catastrophic redshift errors in the optical catalogue. by retaining only those galaxies with a sharp. sinele-moced redshift probability functions in the catalogue.," Using the above criterion further reduces the number of catastrophic redshift errors in the optical catalogue, by retaining only those galaxies with a sharp, single-moded redshift probability functions in the catalogue."
 This criterion is also used. for the AGN selection. along with the other criteria described below.," This criterion is also used for the AGN selection, along with the other criteria described below."
 The photometric redshifts are most accurate for z< (determined from the whole CEDE. catalogue). so to maximise the number of sources in our GN sample (because the redshift distribution peaks at z0.7 (paper 2)) while maintaining a narrow enough range to minimise redshift/Luminositv correlations. we select. sources in the range 0.4xzx;0.6.," The photometric redshifts are most accurate for $z<0.6$ (determined from the whole CFDF catalogue), so to maximise the number of sources in our AGN sample (because the redshift distribution peaks at $z\sim
0.7$ (paper 2)), while maintaining a narrow enough range to minimise redshift/luminosity correlations, we select sources in the range $0.4\leq z \leq0.6$."
 For all galaxies with both photometric and spectroscopic redshifts the rms dilference is ὃς=0.1. over this redshift range.," For all galaxies with both photometric and spectroscopic redshifts the rms difference is $\delta z = 0.1$, over this redshift range."
 Figure 1. shows the photometric vs. spectroscopic redshifts for all the non-stellar AGN in the lth CERS field hat have a CERS measured. redshift. and that also have shotometric redshifts with a high reliability measure (P.> 49).," Figure \ref{AGN_z} shows the photometric vs. spectroscopic redshifts for all the non-stellar AGN in the 14-h CFRS field that have a CFRS measured redshift, and that also have photometric redshifts with a high reliability measure $P_{\Delta z}>0.9$ )."
 The 95percent confidence limits on the photometric redshifts are shown by the error. bars., The $95~per~cent$ confidence limits on the photometric redshifts are shown by the error bars.
 By selecting only AGN with non-stellar light profiles we ensure that the yhotometric redshift estimates are not overly allected by contamination [from nuclear light., By selecting only AGN with non-stellar light profiles we ensure that the photometric redshift estimates are not overly affected by contamination from nuclear light.
 ? clemonstrate that for X-ray selected AGN DPZ is à reliable wav of obtaining hotometric redshifts. as long as the galaxies themselves are not quasar dominated. so we are confident that our method is robust.," \citet{2002ApJ...581..155G}
 demonstrate that for X-ray selected AGN BPZ is a reliable way of obtaining photometric redshifts, as long as the galaxies themselves are not quasar dominated, so we are confident that our method is robust."
 Figure 2 shows the final sample of 31 ACN plotted with hard X-ray luminosity (calculated assuming a photon index of 1.7 for the Ix-correction) versus redshift., Figure \ref{sample_selection} shows the final sample of 31 AGN plotted with hard X-ray luminosity (calculated assuming a photon index of 1.7 for the K-correction) versus redshift.
 We split the full sample into two sub-samples. based on X-ray Iuminosity. to test for any recishift/luminosity correlation: 16 sources are in the low luminosity sample. 15are in the high luminosity sample.," We split the full sample into two sub-samples, based on X-ray luminosity, to test for any redshift/luminosity correlation; 16 sources are in the low luminosity sample, 15are in the high luminosity sample."
 Both sampleshave a median redshift of 0.51. and a Ixolmogorov-Smirnov test shows no evidence for a dillerence in the redshift. distributions.," Both sampleshave a median redshift of 0.51, and a Kolmogorov-Smirnov test shows no evidence for a difference in the redshift distributions."
 Figure 11 of ? , Figure 11 of \citet{2003ApJ...598..886U} 
motion lor in the plane of the Galaxy. compared to out of the plane (see 833.2).,motion for in the plane of the Galaxy compared to out of the plane (see 3.2).
 Table 3 summarizes the apparent motion of in Galactic coordinates with various known sources of motion removed., 	 Table 3 summarizes the apparent motion of in Galactic coordinates with various known sources of motion removed.
 Clearly. these results are of great interest in regard (o the structure and kinematics of the Galaxy. and will be the subject of a later paper.," Clearly, these results are of great interest in regard to the structure and kinematics of the Galaxy and will be the subject of a later paper."
 Whereas (he orbital motion of the Sun (around the Galactic Center) coniplicates estimates of the “in-plane” component of the peculiar motion ofÀ*.. motions out of the plane are simpler to interpret.," 	Whereas the orbital motion of the Sun (around the Galactic Center) complicates estimates of the “in-plane"" component of the peculiar motion of, motions out of the plane are simpler to interpret."
 One needs only to subtract the small Z-component of the Solar Motion [rom the observed motion of to estimate the out-of-plane component of (he peculiar motion ofA*., One needs only to subtract the small Z-component of the Solar Motion from the observed motion of to estimate the out-of-plane component of the peculiar motion of.
. Using stus within about 0.1 kpe of the Sun and measured bv IHipparcos. Dehnen Binney (1998) find (he Z-component lor the Solar Motion to be T.1720.38kms+.," Using stars within about 0.1 kpc of the Sun and measured by Hipparcos, Dehnen Binney (1998) find the Z-component for the Solar Motion to be $7.17 \pm 0.38~\kms$."
 Other determinations of the Z-component of the Solar Motion. 7.61£0.64kms! for stars with distances out to a few kpe by Feast Whitelock (1997). are generally in agreement with the result of Dehnen&Binney(1993).. but have ereater uncertainties.," Other determinations of the Z-component of the Solar Motion, $7.61 \pm 0.64~\kms$ for stars with distances out to a few kpc by Feast Whitelock (1997), are generally in agreement with the result of \citet{DB98}, but have greater uncertainties."
 Thus. we adopt the Dehnen&Binney(19938) value in this paper.," Thus, we adopt the \citet{DB98} value in this paper."
 The motion of in Galactic latitude. shown in Fig.," The motion of in Galactic latitude, shown in Fig."
 4. is re-plotted with an expanded scale in Fig.," 4, is re-plotted with an expanded scale in Fig."
 5., 5.
 The dashed line in (he figure is a weighted least-squares fit to the data., The dashed line in the figure is a weighted least-squares fit to the data.
 The fitted slope of —0.202220.019 masFl. or —7.62:0.7 for Ry—8.0 kpc. agrees almost exactly with the reflex of the Solar Motion out of the Galactic plane.," The fitted slope of $-0.202\pm0.019$ mas, or $-7.6\pm0.7$ for $\rnot=8.0$ kpc, agrees almost exactly with the reflex of the Solar Motion out of the Galactic plane."
 Allowing for a 0.1 degree uncertaintyin the (lt of the Galactic pole (Dlaauwetal.1900) increases slightly the uncertainty in the out of plane motion from 0.7 to 0.8.., Allowing for a 0.1 degree uncertaintyin the tilt of the Galactic pole \citep{Blaauw60} increases slightly the uncertainty in the out of plane motion from $0.7$ to $0.8$.
 Subtracting —TAT+O038 from our measured apparent motion of out of the plane of the Galaxy. we estimate the peculiar motion of to be —0.4d:0.9 (lowarel the north Galactic pole (see Table 3).," Subtracting $-7.17\pm0.38$ from our measured apparent motion of out of the plane of the Galaxy, we estimate the peculiar motion of to be $-0.4\pm0.9$ toward the north Galactic pole (see Table 3)."
" The expected acceleration of (he Sun in its orbit about the Galactic center currently is undetectably small: OG,2/A,z(220kms!?/8.0kpe&6x105 km ! ! ior in angular units OQ42/R7~107 mas vy>7. Thus. unlike velocity measurements. which require precise knowledge and subtraction of the Sun's orbital contribution. any measurement of acceleration can be directly. attributed toA*."," 	The expected acceleration of the Sun in its orbit about the Galactic center currently is undetectably small: $\tnot^2/\rnot\approx(220~\kms)^2/8.0~{\rm kpc}
 \approx6\times10^{-6}$ km $^{-1}$ $^{-1}~,$ or in angular units $\tnot^2/\rnot^2\sim10^{-7}$ mas $^{-2}.$ Thus, unlike velocity measurements, which require precise knowledge and subtraction of the Sun's orbital contribution, any measurement of acceleration can be directly attributed to."
. Since the motion of on the sky appears rectilinear. our data can be used {ο set an upper limit on any apparent acceleration ol A*..," Since the motion of on the sky appears rectilinear, our data can be used to set an upper limit on any apparent acceleration of ."
large masses.,large masses.
 Ελπίς has been confirmed by several studies of numerical scattering experiments (ee.Hut&Baheall1983:Fregeauetal. 2004).," This has been confirmed by several studies of numerical scattering experiments \citep[e.g.][]{hut83, fregeau04}."
. Jased on this. we expect that Ni; should be very small ancl so. as a first approximation. we take δρ=0.," Based on this, we expect that $_{out}$ should be very small and so, as a first approximation, we take $_{out} = 0$."
 However. we also explore the effects of a No; bv assigning a kick velocity to all BSs at birth.," However, we also explore the effects of a non-zero $_{out}$ by assigning a kick velocity to all BSs at birth."
 Lf cdnamical interactions play a role in BS formation. we might naturally expect DSs to be imparted a recoil velocity at birth (or shortly before) due to momentum conservation.," If dynamical interactions play a role in BS formation, we might naturally expect BSs to be imparted a recoil velocity at birth (or shortly before) due to momentum conservation."
 We will return to this assumption in Section 4.., We will return to this assumption in Section \ref{results}.
 Our model contains 4 free parameters. which correspond to the fraction of outcomes that. produce a blue strageler for each formation mechanism (111 collisions. 112 collisions. 2|2 collisions. and binary star evolution).," Our model contains 4 free parameters, which correspond to the fraction of outcomes that produce a blue straggler for each formation mechanism (1+1 collisions, 1+2 collisions, 2+2 collisions, and binary star evolution)."
 These are the f values described in the previous section: fiafi2-fe2.fine," These are the $f$ values described in the previous section: $f_{1+1}, f_{1+2}, f_{2+2},
f_{mt}$."
 We assume that these values are constant throughout each cluster. ancl are also constant between clusters.," We assume that these values are constant throughout each cluster, and are also constant between clusters."
 By fitting the predictions of our model to the observational data. we can determine best-fit values for each of these f£. parameters. and therefore make predictions about whieh blue strageler formation processes are more important.," By fitting the predictions of our model to the observational data, we can determine best-fit values for each of these $f$ parameters, and therefore make predictions about which blue straggler formation processes are more important."
 In order to determine the best values for these f parameters. we need an appropriate statistical test.," In order to determine the best values for these $f$ parameters, we need an appropriate statistical test."
 For this. we follow the approach of Verbunt.Pooley&Bassa(2008).," For this, we follow the approach of \citet{verbunt08}."
. The number of BSs seen in the core of a globular cluster can be cleseribecl by Poisson statistics., The number of BSs seen in the core of a globular cluster can be described by Poisson statistics.
 In. particular. the probability of observing NW sources when (p are expected is: We can calculate a probability for cach cluster. ane then caleulate an overall probability 2 for the model by multiplying the individual 2? values.," In particular, the probability of observing $N$ sources when $\mu$ are expected is: We can calculate a probability for each cluster, and then calculate an overall probability $P^{\prime}$ for the model by multiplying the individual $P$ values."
 We can then vary the f values to maximize this value., We can then vary the $f$ values to maximize this value.
 In practice. these Z values are typically of order ten percent per cluster. and with 35 clusters. the value of £7 quickly becomes extremely small.," In practice, these $P$ values are typically of order ten percent per cluster, and with 35 clusters, the value of $P^{\prime}$ quickly becomes extremely small."
 “Pherefore we chose to work with a moclified version of this value: the deviance of our model to the saturated mocel, Therefore we chose to work with a modified version of this value: the deviance of our model to the saturated model.
 A saturated model is one in which the observed number of sources is exactly equal to the expected. number in cach cluster., A saturated model is one in which the observed number of sources is exactly equal to the expected number in each cluster.
 In other words. this is the best that we can possibly do.," In other words, this is the best that we can possibly do."
 However. because of the nature of Poisson statistics. the probability 2 of such a model (calculated by setting No=ye in Equation 12)) is not equal to 1. but in fact has some smaller value.," However, because of the nature of Poisson statistics, the probability $P$ of such a model (calculated by setting $N=\mu$ in Equation \ref{eqn:stats}) ) is not equal to 1, but in fact has some smaller value."
 For the numbers of blue stragelers in our clusters. the 2? values for the saturated model run from 0.044 to 0.149. and the value of P is 2.08.10.H.," For the numbers of blue stragglers in our clusters, the $P$ values for the saturated model run from 0.044 to 0.149, and the value of $P^{\prime}$ is $2.08 \times 10^{-41}$."
 The deviance of any model from the saturated mocel is given by The model which minimizes this quantity will be our best-fit’ model., The deviance of any model from the saturated model is given by The model which minimizes this quantity will be our best-fit model.
 Ideally. the sealed deviance (2/ Nporometers) should be equal to 1 fora best fit.," Ideally, the scaled deviance $D/(N_{data}-N_{parameters}$ ) should be equal to 1 for a best fit."
 Given the simplicity of our model we expect that our values will not provide this kind of agreement. ancl we simply look for the nmocel which provides the minimum of the scaled deviance.," Given the simplicity of our model, we expect that our values will not provide this kind of agreement, and we simply look for the model which provides the minimum of the scaled deviance."
 In this section. we present the results of comparing our model predictions to the observations.," In this section, we present the results of comparing our model predictions to the observations."
 After presenting the results for a constant core binary fraction for all clusters. we explore the implications of adopting a core binary fraction that depends on the cluster. luminosity. as reported. in Sollimaetal.(2007) and Miloneetal.(2008).," After presenting the results for a constant core binary fraction for all clusters, we explore the implications of adopting a core binary fraction that depends on the cluster luminosity, as reported in \citet{sollima07} and \citet{milone08}."
. The predictions of our model for our initial choice of assumptions are shown in Figure L.., The predictions of our model for our initial choice of assumptions are shown in Figure \ref{fig:model_best}.
 These numbers are plotted against the total stellar mass in the core along with he number of BSs observed in the core (filled triangles)., These numbers are plotted against the total stellar mass in the core along with the number of BSs observed in the core (filled triangles).
" We xot both the total number of BSs predicted to have formed within rye, in the last res vers (Nes in Equation 1: open circles). as well as the total number formed. only in he core (Nea | Nos: small filled circles)."," We plot both the total number of BSs predicted to have formed within $r_{max}$ in the last $\tau_{BS}$ years $_{BS}$ in Equation \ref{eqn:number-bss}; open circles), as well as the total number formed only in the core $_{coll}$ + $_{bin}$; small filled circles)."
 Upon comparing gs do the observed. number of BSs in the core. the best-ting model parameters predict that most DSs are formed rom binary star evolution. with a small contribution from 2|2 collisions being needed. in order to obtain the best xossible match to the observations.," Upon comparing $_{BS}$ to the observed number of BSs in the core, the best-fitting model parameters predict that most BSs are formed from binary star evolution, with a small contribution from 2+2 collisions being needed in order to obtain the best possible match to the observations."
 Ehe ideal contribution rom 2| collisions constitutes at most a few percent of the wedieted total for most of the clusters in our sample., The ideal contribution from 2+2 collisions constitutes at most a few percent of the predicted total for most of the clusters in our sample.
 I5ven or our best-fitting mocel parameters. our initial choice of assumptions prediets too few BSs in clusters with small core κ...," Even for our best-fitting model parameters, our initial choice of assumptions predicts too few BSs in clusters with small core masses."
" We tried changing our assumption of a constant D, for all clusters to one for which the core binary fraction varies with the total cluster magnitude.", We tried changing our assumption of a constant $_b$ for all clusters to one for which the core binary fraction varies with the total cluster magnitude.
 First. we adopted a dependence of the form: where My is the total cluster V. magnitude.," First, we adopted a dependence of the form: where $_V$ is the total cluster V magnitude."
 “Phis relation comes from [fitting a line of best-fit το the observations of Sollimaetal.(2007)... who studied the binary fractions in a sample of 13 low-density GC's (we caleulated an average of columns 3 and 4 in their Table 3 and. used these binary fractions to obtain Equation 14)).," This relation comes from fitting a line of best-fit to the observations of \citet{sollima07}, who studied the binary fractions in a sample of 13 low-density GCs (we calculated an average of columns 3 and 4 in their Table 3 and used these binary fractions to obtain Equation \ref{eqn:sollima07}) )."
" In order to prevent negative binary fractions. we impose a minimum binary [fraction of 7""D=0.01."," In order to prevent negative binary fractions, we impose a minimum binary fraction of $_b^{min} = 0.01$."
" In other words. we set f,= OP"" if Equation 14 gives a binary fraction less than [7."," In other words, we set $_b =$ $_b^{min}$ if Equation \ref{eqn:sollima07}
 gives a binary fraction less than $_b^{min}$."
 As res," As before, we adopt an average semi-major axis of 2 AU."
ults of this comparison are presented in Figure 2.., The results of this comparison are presented in Figure \ref{fig:model_rmax_sollima_best}.
 As in Figure 1.. both the numbers of observed. (filled: triangles) ancl predicted. (open circles) BSs in the core are. plotted versus the total stellar mass in the core.," As in Figure \ref{fig:model_best}, both the numbers of observed (filled triangles) and predicted (open circles) BSs in the core are plotted versus the total stellar mass in the core."
 Once again. the sreclictecl numbers inelude all BSs formed. within tye. in he last rgs vears.," Once again, the predicted numbers include all BSs formed within $_{max}$ in the last $\tau_{BS}$ years."
 The best-itting. model parameters. for his comparison suggest that both sinele-single collisions and binary star evolution are significant contributors to BS ormation., The best-fitting model parameters for this comparison suggest that both single-single collisions and binary star evolution are significant contributors to BS formation.
 Single-single collisions contribute up to several ens of a percent of the predicted total in several clusters., Single-single collisions contribute up to several tens of a percent of the predicted total in several clusters.
" We obtain a deviance in this case that is significantly larger han that obtained for the best-Litting model parameters assuming a constant b, of 104. for all clusters.", We obtain a deviance in this case that is significantly larger than that obtained for the best-fitting model parameters assuming a constant $_b$ of $10\%$ for all clusters.
 Equation 14 eives higher binary fractions in low-mass clusters relative, Equation \ref{eqn:sollima07} gives higher binary fractions in low-mass clusters relative
"with 'The other four components8) 1(¢)""can be derived as c conjugates.",with The other four components can be derived as complex conjugates.
". The results are given. as a function ο.of Z3 and therefore when the coordinates (in partigul shear coordinates) are expressed in the (r1,r2) bas"," The results are given as a function of $z_3$ and $\zb_3$ and therefore when the coordinates (in particular the shear coordinates) are expressed in the $(\vr_1,\vr_2)$ basis."
 The amplitude of the shear correlation fungtig then be compared to the amplitude of the conv correlation function., The amplitude of the shear correlation function can then be compared to the amplitude of the convergence correlation function.
 To start with one can consideP£$(8)? defined as It is to be noted that there are some symmetry properties in this quantity due to the fact it should be invariant under z>1/(1--2) and z—11/z when properly rescaled (these changes correspond to transform that leaves the triangle shape unchanged; the amplitude of the shear functions are then only changed by a scale factor effect)., To start with one can consider $\xi_{\gamma}(z)$ defined as It is to be noted that there are some symmetry properties in this quantity due to the fact it should be invariant under $z\to 1/(1-z)$ and $z\to 1-1/z$ when properly rescaled (these changes correspond to transform that leaves the triangle shape unchanged; the amplitude of the shear functions are then only changed by a scale factor effect).
 Its amplitude in shown on Fig., Its amplitude in shown on Fig.
 1 where the spatial dependence of the amplitude of the shear correlation function is compared to that of the convergence correlation function., \ref{xi3norm} where the spatial dependence of the amplitude of the shear correlation function is compared to that of the convergence correlation function.
" Whereas the convergence correlation function has a very simple spatial dependence, the spatial dependence of the shear correlation is much more complex exhibiting minima for certain configurations."," Whereas the convergence correlation function has a very simple spatial dependence, the spatial dependence of the shear correlation is much more complex exhibiting minima for certain configurations."
 It is to be noted however that these patterns are quite dependent of the power spectrum index (right panel of Fig., It is to be noted however that these patterns are quite dependent of the power spectrum index (right panel of Fig.
 1 and Fig., 1 and Fig.
 2)., 2).
 In particular the existence of a local maximum of the signal for equilateral triangles vanishes for steep spectra., In particular the existence of a local maximum of the signal for equilateral triangles vanishes for steep spectra.
 A fair fraction of the signal however can be extracted from a subpart of the correlation functions., A fair fraction of the signal however can be extracted from a subpart of the correlation functions.
 Following Bernardeau et al. (, Following Bernardeau et al. (
2003) it is possible to map at least,2003) it is possible to map at least
0 mmaps.,$\sigma$ maps.
 Lhe observed velocity field exhibits some departures from what is expected from axisvmmetrey., The observed velocity field exhibits some departures from what is expected from axisymmetry.
 We measure a tilt of the zero-velocity curve of953., We measure a tilt of the zero-velocity curve of $9\deg \pm3\deg$.
 This might partly be due to dust. extinction. perturbing the line-of-sight. kincmiatics. although extinction is significant. only on the South-East side.," This might partly be due to dust extinction perturbing the line-of-sight kinematics, although extinction is significant only on the South-East side."
 This should be confirmed by spectroscopy with higher signal-to-noise ratio., This should be confirmed by spectroscopy with higher signal-to-noise ratio.
 The central velocity gradient is about 110 + for ~1094 kkpe 1)., The central velocity gradient is about 110 $^{-1}$ (or $\sim 1094$ $^{-1}$ $^{-1}$ ).
 The velocity dispersion field is nearly [lat in the central two areseconds with a value around ~215 and slowly clecreases outwards.," The velocity dispersion field is nearly flat in the central two arcseconds with a value around $\sim 215$, and slowly decreases outwards."
 We derived line strength maps from the sspectra. namelyMeb.. FeS270. Fe5335 and. Fe5406.," We derived line strength maps from the spectra, namely, Fe5270, Fe5335 and Fe5406."
. Here we present the nunap. (classically defined as (ΡΟ| be5335)/2) as well as the nunap and. correction for the contribution of he da55200 emission line (Figs.," Here we present the map, (classically defined as $(\mbox{Fe5270} + \mbox{Fe5335}) / 2$ ) as well as the map and correction for the contribution of the 5200 emission line (Figs."
 T and 8)). to illustrate its elfect on theline strength (see CGoudfrooij Enisellem 1996).," \ref{fig:Mgmap} and \ref{fig:Femap}) ), to illustrate its effect on theline strength (see Goudfrooij Emsellem 1996)."
 The highest vvalue of 5.1c0.2 is reached about 0.9 aresee [rom he centre along the major-axis. with its svmnmetric point wing ool 4.40.2 (the central value is 4.3 20.2).," The highest value of $5.1 \pm 0.2$ is reached about $-0.9$ arcsec from the centre along the major-axis, with its symmetric point having of $4.4 \pm 0.2$ (the central value is $4.3 \pm 0.2$ )."
 There is a slight decrease of outwards. with values at the edge of the field being around 4.," There is a slight decrease of outwards, with values at the edge of the field being around 4."
 Phe aand be5406 maps are consistent with being Hat., The and Fe5406 maps are consistent with being flat.
 Higher signal-to-noise spectra would help confirming the absence of significant line strength gradients in the central areseconcds., Higher signal-to-noise spectra would help confirming the absence of significant line strength gradients in the central arcseconds.
 We now present the emission-line maps derived from fits of the emission lines present in the ddata cube., We now present the emission-line maps derived from fits of the emission lines present in the data cube.
 Phese maps were built using a pixel size of 0718. half of the final sky sampling.," These maps were built using a pixel size of $0\farcs18$ , half of the final sky sampling."
S00 models: both models suggest a full-width half maximum spread of factors of a few to >10 and therefore an overall spread in absolute ages that is larger than the central age of the distribution and the median age of the ONC found from the H-R diagram.,S00 models: both models suggest a full-width half maximum spread of factors of a few to $>10$ and therefore an overall spread in absolute ages that is larger than the central age of the distribution and the median age of the ONC found from the H-R diagram.
" The exponential distribution has a favoured starting age of essentially zero for the DAM97 models and a decay time scale, λα of MMyr."," The exponential distribution has a favoured starting age of essentially zero for the DAM97 models and a decay time scale, $\lambda_{a}$ of Myr."
 Again this indicates a spread of ages that is larger than the median age of the sample from the H-R diagram., Again this indicates a spread of ages that is larger than the median age of the sample from the H-R diagram.
" The S00 models favour a non-zero maximum in the age distribution, which is also seen in the H-R diagram-based age distribution shown in Fig. 3,,"," The S00 models favour a non-zero maximum in the age distribution, which is also seen in the H-R diagram-based age distribution shown in Fig. \ref{age2},"
 although a very young starting age cannot be ruled out., although a very young starting age cannot be ruled out.
 The decay timescale is about 2MMyr and again this implies a spread of ages that is larger than the median age of the sample., The decay timescale is about Myr and again this implies a spread of ages that is larger than the median age of the sample.
 I ran some more simulations using a smaller value of ith=15°.," I ran some more simulations using a smaller value of $i_{\rm
th}=15^{\circ}$."
 The results for these are also given in Table 2.., The results for these are also given in Table \ref{results2}. .
" As expected a smaller it, slightly reduces the age dispersion required to explain the data, but does not affect the broad conclusions."," As expected a smaller $i_{\rm th}$ slightly reduces the age dispersion required to explain the data, but does not affect the broad conclusions."
" Conversely, a larger itn would increase the required age dispersion."," Conversely, a larger $i_{\rm th}$ would increase the required age dispersion."
" As final test I constructed age distributions based upon the aages determined from the H-R diagram for the rotation sample, which are shown in the right hand panels of Figs."," As a final test I constructed age distributions based upon the ages determined from the H-R diagram for the rotation sample, which are shown in the right hand panels of Figs."
 2 and 3.., \ref{age} and \ref{age2}. .
 These age distributions were used to construct model Rsini/Rawy: distributions which in turn were compared with the observed Rsini/R3myr distributions via K-S tests.," These age distributions were used to construct model $R\sin i/R_{\rm 3Myr}$ distributions which in turn were compared with the observed $R\sin
i/R_{\rm 3Myr}$ distributions via K-S tests."
" Initially, 4:4 was fixed at 30°."," Initially, $i_{\rm th}$ was fixed at $30^{\circ}$."
 A further set of models was generated with itn=15? to investigate the influence of this parameter.," A further set of models was generated with $i_{\rm
th}=15^{\circ}$ to investigate the influence of this parameter."
 The comparison between the model and observed distributions of Rsini/Hawy. are shown in Fig. 9..," The comparison between the model and observed distributions of $R\sin
i/R_{\rm 3Myr}$ are shown in Fig. \ref{hragedist}."
 Numerical results are listed in Table 3.., Numerical results are listed in Table \ref{results3}.
" All of the model distributions are formally compatible with the data and cannot be rejected at the 90 per cent confidence level, although the DAM97 iin=30° model barely passes this test."," All of the model distributions are formally compatible with the data and cannot be rejected at the 90 per cent confidence level, although the DAM97 $i_{\rm th}=30^{\circ}$ model barely passes this test."
 Hence the age distributions determined from the Hertzsprung-Russell diagram cannot be discounted as a valid description of the age distribution required by the projected radii., Hence the age distributions determined from the Hertzsprung-Russell diagram cannot be discounted as a valid description of the age distribution required by the projected radii.
" Given the uncertainties in determining ages from the H-R diagram - binarity, accretion, reddening, variability (see Hartmann2001) --that do not afflict determinations of the projected radii, this might seem"," Given the uncertainties in determining ages from the H-R diagram – binarity, accretion, reddening, variability (see Hartmann2001) –that do not afflict determinations of the projected radii, this might seem"
a triplet in the Fourier spectrum were seen in the Llipparcos measurements analvzed by Ixoen (2001). (,a triplet in the Fourier spectrum were seen in the Hipparcos measurements analyzed by Koen (2001). (
ii) Another example is provided by the sdB stars: significant amplitude. variations are detected with a time scale of about a vear (Ixilkenny et al.,ii) Another example is provided by the sdB stars: significant amplitude variations are detected with a time scale of about a year (Kilkenny et al.
 1999) or even shorter (Pereira Lopes 2004)., 1999) or even shorter (Pereira Lopes 2004).
 O7Foole et al. (, O'Toole et al. (
2002) have argued that the variability may be caused by the beating of close frequencies with a frequency separation of 0.1 evele 3 or loss.,2002) have argued that the variability may be caused by the beating of close frequencies with a frequency separation of 0.1 cycle $^{-1}$ or less.
 However. more detailed investigations may be needed. (," However, more detailed investigations may be needed. ("
iit) Most. 2 Cephei variables also. show amplitude variations with a variety of time-scales of under à vear (e.g. in EN Lac. Lehmann et al.,"iii) Most $\beta$ Cephei variables also show amplitude variations with a variety of time-scales of under a year (e.g., in EN Lac, Lehmann et al."
 2001) up to decades or even several centuries (see Chapellier 1986). (, 2001) up to decades or even several centuries (see Chapellier 1986). (
iv) White Dwarfs: Lancler et al. (,iv) White Dwarfs: Handler et al. (
2003) analvzed new time-series photometry of two pulsating DD white chwarls anc concluded that. the observed. amplitude. anc period variability cannot be explained by the beating of multiple pulsation miocles. (,2003) analyzed new time-series photometry of two pulsating DB white dwarfs and concluded that the observed amplitude and period variability cannot be explained by the beating of multiple pulsation modes. (
v) Many 9 Seuti variables show amplitude variability associated with the multimode pulsation.,v) Many $\delta$ Scuti variables show amplitude variability associated with the multimode pulsation.
 Similar to the situation seen in white dwarls. the same star may appear to change its pulsation spectrum to appear as if they were different stars at diferent times.," Similar to the situation seen in white dwarfs, the same star may appear to change its pulsation spectrum to appear as if they were different stars at different times."
 However. detailed analyses of 4 CVn showed that the modes do not disappear. but ave still present at small amplitudes.," However, detailed analyses of 4 CVn showed that the modes do not disappear, but are still present at small amplitudes."
 For the majority of the well-studicd stars. frequeney pairs with frequency separations less than 0.06 evcle + can be seen in the power spectra.," For the majority of the well-studied stars, frequency pairs with frequency separations less than 0.06 cycle $^{-1}$ can be seen in the power spectra."
 These pairs may be a rellection of close frequencies. amplitude variability of a single mode or observational errors.," These pairs may be a reflection of close frequencies, amplitude variability of a single mode or observational errors."
 For the star DI CALL. Dreger Bischof (2002) showed from. detailed. phasc-amplitude tests that the amplitude variability was the result of beating of separate modes with close frequencies.," For the star BI CMi, Breger Bischof (2002) showed from detailed phase-amplitude tests that the amplitude variability was the result of beating of separate modes with close frequencies."
 This paper also summarizes the situation in these stars., This paper also summarizes the situation in these stars.
 The Blazhko Elfect. therefore. exists in many (vpes of stellar pulsators.," The Blazhko Effect, therefore, exists in many types of stellar pulsators."
 LO is presently uncertain whether a single physical mechanism is responsible for the phenomenon in all these stars., It is presently uncertain whether a single physical mechanism is responsible for the phenomenon in all these stars.
 We wish to distinguish between two explanations for the observed amplitude ancl phase variability observed in so many types of pulsating stars: (i) The amplitude variability is. caused: by beating between two (or more) close frequencies., We wish to distinguish between two explanations for the observed amplitude and phase variability observed in so many types of pulsating stars: (i) The amplitude variability is caused by beating between two (or more) close frequencies.
 Since amplitude variability is common. the occurrence of close [requencies cannot be an accidental coincidence between randomly distributed. frequencies.," Since amplitude variability is common, the occurrence of close frequencies cannot be an accidental coincidence between randomly distributed frequencies."
 In. this. hypothesis. a physical mechanism must exist to excite the close eigenfrequencies to observable levels. (," In this hypothesis, a physical mechanism must exist to excite the close eigenfrequencies to observable levels. ("
ii) “Truc amplitude ancl phase variations not caused by beating.,ii) True amplitude and phase variations not caused by beating.
 These include periodic behavior. possibly caused by rotational modulation or magnetic fields. or evele-to-cvcle variations as observed in. Miras.," These include periodic behavior, possibly caused by rotational modulation or magnetic fields, or cycle-to-cycle variations as observed in Miras."
 The observational tests to. distinguish. between the different physical explanations require a very large amount of data covering many months over several vears. so that detailed. analyses are so far available only for Bl CAL.," The observational tests to distinguish between the different physical explanations require a very large amount of data covering many months over several years, so that detailed analyses are so far available only for BI CMi."
 In this paper. we want to examine the ὁ Seuti star PC: Vir. for which extensive data ancl miultifrequeney analyses are available (Breeer et al.," In this paper, we want to examine the $\delta$ Scuti star FG Vir, for which extensive data and multifrequency analyses are available (Breger et al."
 2004. 2005: note that the table of frequencies given in the latter reference already refers to the analysis carried out in this paper. and lists some preliminary results).," 2004, 2005: note that the table of frequencies given in the latter reference already refers to the analysis carried out in this paper, and lists some preliminary results)."
 The previous work has shown that in PG Vir three modes sect to possess strongly variable amplitudes., The previous work has shown that in FG Vir three modes seem to possess strongly variable amplitudes.
 The present paper also examines the svstematies of the occurrence of close frequencies. defined here as multiple [requencies separated by 0.1 evele 1 or Less.," The present paper also examines the systematics of the occurrence of close frequencies, defined here as multiple frequencies separated by 0.1 cycle $^{-1}$ or less."
 In the Fourier spectrum. of short cata sets. (wo. close frequencies heating with each other appear as a single frequeney with variable amplitude.," In the Fourier spectrum of short data sets, two close frequencies beating with each other appear as a single frequency with variable amplitude."
 For long data sets with sullicient frequency. resolution. two peaks will be shown.," For long data sets with sufficient frequency resolution, two peaks will be shown."
 In the case of a single frequency with variable amplitude and/or phasing. also two (or more) peaks occur.," In the case of a single frequency with variable amplitude and/or phasing, also two (or more) peaks occur."
 Consequently.modes. while the absence of a double peak does no exclude the possibility.," Consequently, while the absence of a double peak does not exclude the possibility."
 This is a fact well recognized in the iterature., This is a fact well recognized in the literature.
 Lt is important to apply an additional method bevon recognizing two peaks in the power spectrum to separate he two hypotheses., It is important to apply an additional method beyond recognizing two peaks in the power spectrum to separate the two hypotheses.
 Εις is especially the case whenever he detection of the two peaks could be due to imporfec data coverage such as data length being of the order of the reat [requeney., This is especially the case whenever the detection of the two peaks could be due to imperfect data coverage such as data length being of the order of the beat frequency.
 Fortunately. beating between two (or more) requencies produces amplitude ancl phase variations which are mathematically related.," Fortunately, beating between two (or more) frequencies produces amplitude and phase variations which are mathematically related."
 Phe most extreme ancl easily recognizable situation occurs when two close frequencies jwe the same amplitude: this leads to a half à cycle phase shift at the time of minimum amplitude ofthe visible (single) requenev., The most extreme and easily recognizable situation occurs when two close frequencies have the same amplitude: this leads to a half a cycle phase shift at the time of minimum amplitude of the visible (single) frequency.
 Even when the amplitudes. differ. beating has a specific signature in the amplitude of phase shifts of an assumed single frequeney.," Even when the amplitudes differ, beating has a specific signature in the amplitude of phase shifts of an assumed single frequency."
 In. particular. the amplitude anc hase vary svachronously with a phase shift close to 90 (Psesevich 1975).," In particular, the amplitude and phase vary synchronously with a phase shift close to $\degr$ (Tsesevich 1975)."
 Therefore. it is possible to separate beating from true amplitude variability by studying the relationship between he amplitude and phase variations of an assumed single requenev.," Therefore, it is possible to separate beating from true amplitude variability by studying the relationship between the amplitude and phase variations of an assumed single frequency."
 “Lhe method. consists of obtaining an extensive whotometric data base. subcdividing the data into short time ins and calealating the amplitude and phase of an assumed optimum single frequency for each time bin.," The method consists of obtaining an extensive photometric data base, subdividing the data into short time bins and calculating the amplitude and phase of an assumed optimum single frequency for each time bin."
 Furthermore. if other pulsation frequencies are also present. they need to be corrected for.," Furthermore, if other pulsation frequencies are also present, they need to be corrected for."
 In practice. large cilferences in the amplitudes of the frequencies beating with each other. observational noise. as well as short (relative to the beat period) data sets may make the recognition cillicult.," In practice, large differences in the amplitudes of the frequencies beating with each other, observational noise, as well as short (relative to the beat period) data sets may make the recognition difficult."
 The recent multifrequeney analysis of FG Vir (Breger ct al., The recent multifrequency analysis of FG Vir (Breger et al.
 2005) led to the detection of more than 75 frequencies. of," 2005) led to the detection of more than 75 frequencies, of"
low-metallicity ISED obviously presents stronger flix at shorter wavelength when compared wilh high-metallicily ISED (Schulz et al.,low-metallicity ISED obviously presents stronger flux at shorter wavelength when compared with high-metallicity ISED (Schulz et al.
 2002)., 2002).
 Discussions about these problems in some saniple clusters in (his case are present in (he following., Discussions about these problems in some sample clusters in this case are present in the following.
 Thirty stars were cited as BSSs in Nine 2 in AL95. based on the Johnson-Consins UBVR CCD photometry of Aparicio et al. (," Thirty stars were cited as BSSs in King 2 in AL95, based on the Johnson-Cousins UBVR CCD photometry of Aparicio et al. ("
1990). in which only BSS candidates within the core radius are considered. in order to avoid the probable contamination by field stars.,"1990), in which only BSS candidates within the core radius are considered, in order to avoid the probable contamination by field stars."
 As a cluster of large Nyy number (as shown in Figure 7). the effects of BSSs on the ISED of this cluster are obviously due to the large number of bright DSSs.," As a cluster of large $_{BS}$ number (as shown in Figure 7), the effects of BSSs on the ISED of this cluster are obviously due to the large number of bright BSSs."
 As shown in Figure 8. four stars were defined as BSSs in NGC 6939 in Cannon Llovd (19690).," As shown in Figure 8, four stars were defined as BSSs in NGC 6939 in Cannon Lloyd (1969)."
 Star 59 (59° is the klenüfication number of the star marked in (he finding chart in the photometric work of Geisler (1933)., Star 59 ('59' is the identification number of the star marked in the finding chart in the photometric work of Geisler (1988).
 All the munbers in the following are the same identifications) was mentioned as à BSS in Geisler (1983). although it was unlikely a member.," All the numbers in the following are the same identifications) was mentioned as a BSS in Geisler (1988), although it was unlikely a member."
 In AL95. NGC 6939 has No number of eighty which is smaller than that of Ning 2 and bigger than that of NCC3680. thus NGC 6939 has an intermediate richness in our working sample.," In AL95, NGC 6939 has $_2$ number of eighty which is smaller than that of King 2 and bigger than that of NGC3680, thus NGC 6939 has an intermediate richness in our working sample."
 In (his situation. the effects of BSSs on the ISED should be better attributed to the high luminosities of the DSSs. since the Nyy of this eluster is much smaller than that of Ning 2.," In this situation, the effects of BSSs on the ISED should be better attributed to the high luminosities of the BSSs, since the $_{BS}$ of this cluster is much smaller than that of King 2."
 These (wo clusters possess small numbers of both Nyy and Ne., These two clusters possess small numbers of both $_{BS}$ and $_2$.
 The single BSS in Berkeley 42 was selected from the CCD photometry of Aparicio et al. (, The single BSS in Berkeley 42 was selected from the CCD photometry of Aparicio et al. (
1991).,1991).
 Four stars (29. 33. 43 and 56) were identified as BSSs in NGC 3680 in AL95. according to the photo-electric observations of Eggen (1969). who later (Eggen 1933) identified stars 3. 29. 33 ancl 56 as BSSs and stars 20 and 34 as red stragglers’ wilh ubvy. photometry.," Four stars (29, 33, 43 and 56) were identified as BSSs in NGC 3680 in AL95, according to the photo-electric observations of Eggen (1969), who later (Eggen 1983) identified stars 3, 29, 33 and 56 as BSSs and stars 20 and 34 as 'red stragglers' with ubvy photometry."
 In the CCD photometry of Anthony-Twarog et al. (, In the CCD photometry of Anthony-Twarog et al. (
1991). stars 29 and 33 were considered as probable cluster members. but star 56 is unlikely a member.,"1991), stars 29 and 33 were considered as probable cluster members, but star 56 is unlikely a member."
 Meanwhile. an important. presence of binaries arouncl (he turnolf point of NGC 3680 was observed by Anthony-Twarog et al. (," Meanwhile, an important presence of binaries around the turnoff point of NGC 3680 was observed by Anthony-Twarog et al. ("
1991).,1991).
 As shown in Figure 9. (here is a BSS in both (wo clusters whose positions in the CMDs highly exceed the cluster turnoff that contributes most of the lieht in the blue part of the ISED.," As shown in Figure 9, there is a BSS in both two clusters whose positions in the CMDs highly exceed the cluster turnoff that contributes most of the light in the blue part of the ISED."
estimate the seusitivity and accuracy of SZ observatious.,estimate the sensitivity and accuracy of SZ observations.
 Given a sky scan rate. how accurately can the SZ power spectrum be measured?," Given a sky scan rate, how accurately can the SZ power spectrum be measured?"
 How mauy cluster cau be found?, How many cluster can be found?
 How accurately cau he SZ decrement be determined for iudividual clusters?, How accurately can the SZ decrement be determined for individual clusters?
 What are optimal strategies for these ueasurements?, What are optimal strategies for these measurements?
 La this paper. we will take AMIBA as our target to address these questions.," In this paper, we will take AMIBA as our target to address these questions."
 This paper is organized as follows: in 82.. we describe the SZ ellect. its statistics ancl our uethod to analyze these statistics.," This paper is organized as follows: in \ref{sec:application}, we describe the SZ effect, its statistics and our method to analyze these statistics."
 We theu present our simulation results of these statistics aud »ossible coustrain [rom observatious bu., We then present our simulation results of these statistics and possible constrain from observations \\ref{sec:simulation}) ).
 ‘efsec:AMIBA we simulate AMIBA cdriftscau observations to estimate the accuracy of AMIBA ueasurement of these statistics., In \\ref{sec:AMIBA} we simulate AMIBA driftscan observations to estimate the accuracy of AMIBA measurement of these statistics.
 We conclude in 85.., We conclude in \ref{sec:conclusion}.
 The thermal SZ effect causes CMB temperature and intensity fluctuations in the sky., The thermal SZ effect causes CMB temperature and intensity fluctuations in the sky.
" At a eivenposition given by a unit vector 5 pointiug [rom the earth to some point ou the sky. the fractional change in the CMB intensity 6/,/2, depends on both the observing frequency v anc the iutegral of the gas pressure along the liue of sight (Zel'dovichaudSuuyvaev1969):: For the CMB iutensitv. the spectral dependeuce lor inverse compton scattering off non-relativistic electronis is described by SyGr)=avay(2—v/[2tauhtce/2)]) where: =hv/(kpToun)7/97 Ghz."," At a givenposition given by a unit vector $\hat{n}$ pointing from the earth to some point on the sky, the fractional change in the CMB intensity $\delta
I_{\nu}/I_{\nu}$ depends on both the observing frequency $\nu$ and the integral of the gas pressure along the line of sight \citep{Zeldovich69}: For the CMB intensity, the spectral dependence for inverse compton scattering off non-relativistic electrons is described by $S_I(x)=\frac{xe^x}{e^{x}-1}\left(2-x/[2\tanh(x/2)]\right)$ where $x\equiv h\nu/(k_B T_{\rm CMB})=\nu/57$ Ghz."
 The correspoudiug fractional chauge in CMB temperature O(n)=Sau) is then: Iu the Ravleieh-Jeaus limit Ge« 1). Spr)=1.," The corresponding fractional change in CMB temperature $\Theta(\hat{n}) \equiv\frac{\Delta
T_{\rm CMB}(\hat{n})}{T_{\rm CMB}}$ is then: In the Rayleigh-Jeans limit $x\ll 1$ ), $S_I(x)=S_T(x)=1$."
 δε) reaches the highest response. Str)~1.6 atv~100 Ciz. whieh is well matched to the AMIBA [requency range of 80-100 Gli.," $S_I(x)$ reaches the highest response, $S(x)\simeq 1.6$ at $\nu \simeq 100$ Ghz, which is well matched to the AMIBA frequency range of 80-100 Ghz."
 Spor) mouotonously decreases with [requeucy., $S_T(x)$ monotonously decreases with frequency.
 We show ο) aud ο) in fig. (1)), We show $S_I(x)$ and $S_T(x)$ in fig. \ref{fig:sx}) )
 where the[requeney responses of various experiments which already have the CMB observation data at small angular scales {ο 2000) such as ATCA (Subralimauyanetal.2000).. ΒΙΛΑ (Dawsonetal.2001)..CBE. SUZIE (Churchetal.L997) aud VLA (Partridgeetal.L997) are shown.," where thefrequency responses of various experiments which already have the CMB observation data at small angular scales $l>2000$ ) such as ATCA \citep{Subrahmanyan00}, BIMA \citep{Dawson01}, SUZIE \citep{Church97} and VLA \citep{Partridge97} are shown."
" The gas pressure depeudence is described by the Comptou y parameter Ti. n, aud 2. are the temperature. uumber density aud pressure of free electrous. respectively."," The gas pressure dependence is described by the Compton $y$ parameter $T_e$ , $n_e$ and $P_e$ are the temperature, number density and pressure of free electrons, respectively."
 dl is the proper distauce interval along the path of CMB photous., $dl$ is the proper distance interval along the path of CMB photons.
 oy. m. aud ce have their usual," $\sigma_T$ , $m_e$ and $c$ have their usual"
the literature incorporates such mass and energy loss in (he shock jump conditions.,the literature incorporates such mass and energy loss in the shock jump conditions.
 In the framework of our model presented here. mass loss is coupled to energy loss via the shock jump conditions. and therefore it must be considered simultaneously.," In the framework of our model presented here, mass loss is coupled to energy loss via the shock jump conditions, and therefore it must be considered simultaneously."
 It is (his point that motivates us to explore. for the first (nme. the lormation of outflowing particles as a consequence of 1e formation of dissipative standing shocks in accretion in a fully relativistic treatment.," It is this point that motivates us to explore, for the first time, the formation of outflowing particles as a consequence of the formation of dissipative standing shocks in accretion in a fully relativistic treatment."
 The formalism of our current model is partly based on (he previous works (Yang&Ixalfatos1993:Fukumura&Tsuruta 2004).," The formalism of our current model is partly based on the previous works \citep[][]{Yang95,Lu98,FT04}."
. Our main objectivelows in this study is jerefore (o explore in details a possibility that the formation of oni (jets/winds) can scour ab a dissipative shock front in (ransonic accreting flows: i.e.. a connection. between ponjockede-accreting gas and the outflowineg particles.," Our main objective in this study is therefore to explore in details a possibility that the formation of outflows (jets/winds) can occur at a dissipative shock front in transonic accreting flows: i.e., a connection between shocked-accreting gas and the outflowing particles."
 The structure of (his paper is as follows., The structure of this paper is as follows.
 and.In 82 we review and explain our simplified moclel jab simultaneouslv considers both shocks outflows with jjump conditions., In 2 we review and explain our simplified model that simultaneously considers both shocks and outflows with appropriate jump conditions.
 cLlain parameterdependence of the shock-outllow solutions appropriateis explored in 83. where we V.row the nature of the shock-outflow solutions and the corresponding global accretion solutions.," Main parameter dependence of the shock-outflow solutions is explored in 3, where we show the nature of the shock-outflow solutions and the corresponding global accretion solutions."
 Our primary eoal in these analysis is to examine the coupling between the solutions., Our primary goal in these analysis is to examine the coupling between the shock-outflow solutions.
 In 84 we cliscuss our results and make some observation implications., In 4 we discuss our results and make some observation implications.
 Briel simnmarvy and concluding remarks are given there as well., Brief summary and concluding remarks are given there as well.
 In black hole accretion. accreting gas must be (ransonic.," In black hole accretion, accreting gas must be transonic."
 After passing through a first sonic radius. the gas is slowed down. ancl a shock max develop.," After passing through a first sonic radius, the gas is slowed down, and a shock may develop."
 For causality. the shocked eas must become supersonic again belore crossing the event horizon.," For causality, the shocked gas must become supersonic again before crossing the event horizon."
 Below. we will explain the details of our simplified mocel.," Below, we will explain the details of our simplified model."
 We consider a steady-state. axisvmmetric accreting flows in Ixerr geometry.," We consider a steady-state, axisymmetric accreting flows in Kerr geometry."
 The spacetime metric is expressed bv (he Bover-Lindqtist coordinates as ⊄⋯↲∐↓≺↕∖∖≺↕∐≼⇂∖↕≻≼↲≺↕∐≺≺↕∐⋃∏∟∐(or ," The spacetime metric is expressed by the Boyer-Lindquist coordinates as where $\Delta \equiv r^2-2mr+a^2$, $\Sigma \equiv r^2+a^2 \cos^2
\theta$, $A \equiv (r^2+a^2)^2-a^2 \Delta \sin^2 \theta$."
momentum err parameter) a black hole., $m$ and $a$ are mass and specific angular momentum (or Kerr parameter) of a black hole.
 Following the standard eeometrized units. we have taken Go =ο1in equation ol(1)).," Following the standard geometrized units, we have taken $G=c=1$ in equation \ref{eq:BL}) )."
 Thus. the length (distance) r and @ are measured in units of m," Thus, the length (distance) $r$ and $a$ are measured in units of $m$."
 Since we are interested in (he equatorial flows. we set 9=7/2 throughout this paper.," Since we are interested in the equatorial flows, we set $\theta=\pi/2$ throughout this paper."
 The black hole horizon is then expressed by ryaamvnm?—«0.," The black hole horizon is then expressed by $r_{\rm{H}} \equiv
m+\sqrt{m^2-a^2}$."
 The sell-gravitv of the flow is ignored. ancl we do not include the effects of magnetic [ields for simplicity.," The self-gravity of the flow is ignored, and we do not include the effects of magnetic fields for simplicity."
 We follow the earlier works on relativistic shock formation as follows: accretion time is asstuned (o be shorter than that of energy diffusion. thus treating the flows as acliabatic," We follow the earlier works on relativistic shock formation as follows; accretion time is assumed to be shorter than that of energy diffusion, thus treating the flows as adiabatic"
Rich clusters of. galaxies are used to study the large-scale structure of the universe on scales larger than about 105.Mpe th=Ty100Ins“Mpc *).,Rich clusters of galaxies are used to study the large-scale structure of the universe on scales larger than about $10 \hmpc$ $h = {\rm H_0\;/\;100\;km\;s^{-1}\;Mpc^{-1}}$ ).
" To measure the clustering streugth of clusters the spatial two-point correlation function Abell clusters (Abell1958: Abell. Corwin. (CF)Olowin ofthe1959) has been estimated by many authors (Baleall Sonecira 1983: IIuchraetal. 1990:: Postman. Huchlra. Celler 1992: Peacock&WestL992: Millerotal, 1999: sce Bahcall1988 fora review). aud hasbeen fouud tobe cousisteut with the power law with the correlation leugth ry~20—26h|Mpe and with the iudex 5~2."," To measure the clustering strength of clusters the spatial two-point correlation function (CF) of the Abell clusters \cite{abe58}; Abell, Corwin, Olowin 1989) has been estimated by many authors (Bahcall Soneira 1983; \cite{huc90}; ; Postman, Huchra, Geller 1992; \cite{pea92}; \cite{mil99}; see \cite{bah88} for a review), and has been found to be consistent with the power law with the correlation length $r_0 \simeq 20 \sim 26 \hmpc$ and with the index $\gamma \simeq 2$."
 Receuthy new catalogs of clusters have been obtained by automated selection from the Ediuburel-Dirham Southern Galaxy Catalog and the Automatic Plate Measming(APM) Galaxy19973). Survev (Liuusdenetal1992: M, Recently new catalogs of clusters have been obtained by automated selection from the Edinburgh-Durham Southern Galaxy Catalog and the Automatic Plate Measuring (APM) Galaxy Survey \cite{lum92}; \cite{dal97}) ).
Motal. AM data of these clusters d(Collinsetal. 1995: Daltonctal.199 153) havebeen to estimate the CF (Nicholetal. 1992:; Daltonetal.1992:; Daltonetal.L99 lay)., Redshift data of these clusters \cite{col95}; ; \cite{dal94b}) ) have been used to estimate the CF \cite{nic92}; ; \cite{dal92}; \cite{dal94a}) ).
 The correlation lengths ry measured for these new cluster saaples have becu reported to be 1Lxry16h !Mpe. which is πιο smaller than that of the Abell clusters.," The correlation lengths $r_0$ measured for these new cluster samples have been reported to be $14 \lesssim r_0 \lesssim 16 \hmpc$ , which is much smaller than that of the Abell clusters."
 There are also more recent new catalogs such as Northern Sky Optical Cluster Survey based ou Digitized Second. Palomar Sky Survey (DPOSS) etal. 2000)) and Sloan Digital Sky Survey. (8DSS-CE) galaxy. cluster catalog (Cotoct 2002))., There are also more recent new catalogs such as Northern Sky Optical Cluster Survey based on Digitized Second Palomar Sky Survey (DPOSS) \cite{gal00}) ) and Sloan Digital Sky Survey Cut-and-Enhance (SDSS-CE) galaxy cluster catalog \cite{got02}) ).
 They need. however. either muore follow- confinuuiiions or spectroscopic observatious before being able to be used to study the spatial clusteriug.," They need, however, either more follow-up confirmations or spectroscopic observations before being able to be used to study the spatial clustering."
 It hasbeen argued. that the countiug luecradius n.the=1.5h.pe for the Abell clusters is so that catalog contains serious projection effects. which cause artificial Hne-of-sieht correlations (Sutherland1988: Efstathiouetal. 1992)).," It has been argued that the counting radius $r_c = 1.5 \hmpc$ for the Abell clusters is so large that the catalog contains serious projection effects, which cause artificial line-of-sight correlations \cite{sut88}; \cite{efs92}) )."
 It has also been poiuted out that the intrinsically subjective nature of the Abell catalog can cause problems in homogeneity aud statistical conipleteuness (Nicholetal. 1992))., It has also been pointed out that the intrinsically subjective nature of the Abell catalog can cause problems in homogeneity and statistical completeness \cite{nic92}) ).
" Mowever. Balcall West (1992) have claimed that the discrepancy between the Abell aud the APM. cluster CFs can be explüued by the richucss dependence of cluster correlation. amplitudes rq=0.14. where u="" is the mean iuterclister separation aud i. is the urea space deusity of clusters."," However, Bahcall West (1992) have claimed that the discrepancy between the Abell and the APM cluster CFs can be explained by the richness dependence of cluster correlation amplitudes $r_0 = 0.4 d_c$, where $d_c = n_c ^{-1/3}$ is the mean intercluster separation and $n_c$ is the mean space density of clusters."
 But Croft et al. (, But Croft et al. (
1997) have analyzed vichucss subsamples of the APA clusters aud argued that there is only aweak depeudence ofcorrelation amplitudes on the cluster richness.,1997) have analyzed richness subsamples of the APM clusters and argued that there is only aweak dependence of correlation amplitudes on the cluster richness.
" Itis. therefore, necessary to understand why there have been disagreements in the cluster correlation length ryand on the relation between ry and d, "," Itis, therefore, necessary to understand why there have been disagreements in the cluster correlation length $r_0$and on the relation between $r_0$ and $d_c$ ."
Tn the, In the
πιοΊο of height over several diffusion timescales.,function of height over several diffusion timescales.
 The first and the second term together on the right hand side ensures that the tracer particles will be distributed according to the backeround eas deusitv field., The first and the second term together on the right hand side ensures that the tracer particles will be distributed according to the background gas density field.
 The third ena describes the settling of the particles., The third term describes the settling of the particles.
 Equation 5 is valid if the motion of the dust does not influence the notion of the eas (the back-reaction from the dust to the eas is negligible)., Equation \ref{eq:diff} is valid if the motion of the dust does not influence the motion of the gas (the back-reaction from the dust to the gas is negligible).
 This condition is iet if the dust to gas ratio ds <1., This condition is met if the dust to gas ratio is $\ll 1$.
 We note that we do not solve for the dust deusitv directly., We note that we do not solve for the dust density directly.
 We follow the motion of dustparticles.. cach of which represents a portion of the total dust mass inside the column.," We follow the motion of dust, each of which represents a portion of the total dust mass inside the column."
 Thus we derive the corresponding velocities (or fluxes) for the first. second. aud third terms of Eq.," Thus we derive the corresponding velocities (or fluxes) for the first, second, and third terms of Eq."
 58 to calculate the position update of the particles., \ref{eq:diff} to calculate the position update of the particles.
 We calculate Dy. the diffusion coefficient of the dust. aud clefine the diffusion velocity. ep.," We calculate $D_d$, the diffusion coefficient of the dust, and define the diffusion velocity, $v_{D1}$."
 The diffusion cocticicnut of the gas can be defined as (Shakura&Sun-vaev.1973) Based on Youdin&Lithwick(2007).. the diffusion coefficieut of the dust cau be calculated as where St is the Stokes umber The average displacement of à particle in 1D during the time step of Af then is The real displacement of the particle (A+) is drawn from a Caussian distribution which has zero mean aud a half width of L.," The diffusion coefficient of the gas can be defined as \citep{Shakura1973}
 Based on \cite{Youdin:2007p576}, the diffusion coefficient of the dust can be calculated as where $St$ is the Stokes number The average displacement of a particle in 1D during the time step of $\Delta t$ then is The real displacement of the particle $\Delta z$ ) is drawn from a Gaussian distribution which has zero mean and a half width of $L$."
 The “diffusion velocity” can then be calculated as This velocity component fries to smear out dust concentrations., The “diffusion velocity” can then be calculated as This velocity component tries to smear out dust concentrations.
 It is iuportaut to note that the real. physical velocity of the particle during turbulent diffusion changes raudonlv every time the ageregate interacts with a turbulent eddy.," It is important to note that the real, physical velocity of the particle during turbulent diffusion changes randomly every time the aggregate interacts with a turbulent eddy."
 The diffusion velocity defined above Is a nunicerical construct to calculate the time-averaged velocity of the particle during a time-interval Af., The diffusion velocity defined above is a numerical construct to calculate the time-averaged velocity of the particle during a time-interval $\Delta t$.
 The secoud terii on the right side of Eq., The second term on the right side of Eq.
 8. results ina systematic velocity terii which pushes particles towards the density maxima of the gas., \ref{eq:diff} results in a systematic velocity term which pushes particles towards the density maxima of the gas.
" This velocity can be determinedl by Usine these two velocity componcuts (ep, and 052). the particles will be distributed according to the eas density profile in a diffusion timescale."," This velocity can be determined by Using these two velocity components $v_{D1}$ and $v_{D2}$ ), the particles will be distributed according to the gas density profile in a diffusion timescale."
" The fact that particles with f;>0 settle towards the midplane aud have a scale height less than JZ, is the result of the third term of Eq. &.."," The fact that particles with $t_s>0$ settle towards the midplane and have a scale height less than $H_g$ is the result of the third term of Eq. \ref{eq:diff},"
 the setthue velocity., the settling velocity.
 The settling velocity of a particle is given by The new position of the particles cau then be determined by using these three velocity teris: This description os identical to the one used in (2010)., The settling velocity of a particle is given by The new position of the particles can then be determined by using these three velocity terms: This description os identical to the one used in \cite{Ciesla2010}.
 The collision model used in this work is similar to the one used in Paper IL., The collision model used in this work is similar to the one used in Paper II.
 There are however two differences., There are however two differences.
 The first cüfference is the additional source of relative velocity due to differential vertical settling (see Eq. 18)):, The first difference is the additional source of relative velocity due to differential vertical settling (see Eq. \ref{eq:set}) ):
 The second difference concerns the calculation of the acrodvuamucal cross section Which is used to calculate the stopping time., The second difference concerns the calculation of the aerodynamical cross section which is used to calculate the stopping time.
 Iu Paper II πο used the ecometrical cross section of the particles (Ormeletal..2007) where r2 is the compact radius of the ageregateCoco (assunüng that the mass of the particle is contaiued iu a colupact sphere of radius (à). and V is the eulareemeut xuuneter.," In Paper II we used the geometrical cross section of the particles \citep{Ormel:2007p93}
 where $r_c$ is the compact radius of the aggregate (assuming that the mass of the particle is contained in a compact sphere of radius $r_c$ ), and $\Psi$ is the enlargement parameter."
 V-- for compact particles aud the value is hieher for fluff& particles COriueletal..2007).., $\Psi=1$ for compact particles and the value is higher for fluffy particles \citep{Ormel:2007p93}.
 This ornmla works well for particles with fracta dinieusionu above 2., This formula works well for particles with fractal dimension above 2.
 However. we use he porosity model of Okuzuiietal. (2009).. aud their model produces ageregates with yactal dimension low 2.," However, we use the porosity model of \cite{Okuzumi2009a}, and their model produces aggregates with fractal dimension below 2."
 If one calculates the stopping ine of such an ageregate iu the Epstein regime (Eq. 103), If one calculates the stopping time of such an aggregate in the Epstein regime (Eq. \ref{eq:ts1}) )
 using the formula in Eq. 21.. ," using the formula in Eq. \ref{eq:cross1}, ,"
one ects that the stopping tue is less than the stopping time of a monomer., one gets that the stopping time is less than the stopping time of a monomer.
 This is clearly unphysical., This is clearly unphysical.
 The reason for this low stopping time (large area) is that Eq., The reason for this low stopping time (large area) is that Eq.
 21 docs not take into account he οί space between the fracta branches’., \ref{eq:cross1} does not take into account the empty space between the `fractal branches'.
 To avoid such unplysical results for agerceates with low fractal dinensious. we use the aerodyvnaniücal cross section as defined in Eq.," To avoid such unphysical results for aggregates with low fractal dimensions, we use the aerodynamical cross section as defined in Eq."
 17 in Okuzuiietal.(2009)., 47 in \cite{Okuzumi2009a}.
 It is eoncrally assumed. that dust particles in the xotopl:rotoplanctaryi disk are ageregates. 1.¢1.0. thev consist of mucron-sized spheres (monomers) Which are counected via surface forces Tenming&Stoguieuko(1996).," It is generally assumed that dust particles in the protoplanetary disk are aggregates, i.e. they consist of micron-sized spheres (monomers) which are connected via surface forces \cite{Henning1996}."
". The material of these monomers can be iron. silicate, organics (tar). different ices aud a nüxture of these (c.e.. silicates with organic or icv mantels)."," The material of these monomers can be iron, silicate, organics (tar), different ices and a mixture of these (e.g., silicates with organic or icy mantels)."
 The collision modeldescribed iu Paper Lis based on laboratory experiments performed ou, The collision modeldescribed in Paper I is based on laboratory experiments performed on
2006b).,.
. No variable optical counterpart was found in deep image dillereneing of r' GMOS/Gemini 8-m data taken at 0.7 and 1.7 davs after the GRD (Priceetal.2006)., No variable optical counterpart was found in deep image differencing of $^\prime$ GMOS/Gemini 8-m data taken at 0.7 and 1.7 days after the GRB \citep{pbf+06}.
. Three sources in the refined Swill NRT error circle were identified. one of which was shown through spectroscopy to be a Galactic star (Berger&Mirabal20068:Raunvantsevetal. 2006).," Three sources in the refined Swift XRT error circle were identified, one of which was shown through spectroscopy to be a Galactic star \citep{bcr06,hm06a,rks+06}."
. Three additional sources are located in or near our modified NRT error circle., Three additional sources are located in or near our modified XRT error circle.
 On 2006 May. 30UTC. using the Low-Resolution Imaging Spectrograph 1995) on the Neck 1 10 m telescope. we imaged the field of OGO502B in the A and ο fillers for 300 and 330 s. respectively.," On 2006 May 30, using the Low-Resolution Imaging Spectrograph \citep{occ+95} on the Keck I 10 m telescope, we imaged the field of 060502B in the $R$ and $g'$ filters for 300 and 330 s, respectively."
 The images were processed in (he usual manner., The images were processed in the usual manner.
 We also observed the fieldl Lrom 2006 May 3 7:48:22 00 9:47:05 with the 1.3m Peters Automated Infrared Imaging Telescope (PAIRITEL) (Bloometal. 2006c).., We also observed the field from 2006 May 3 7:48:22 to 9:47:05 with the 1.3m Peters Automated Infrared Imaging Telescope (PAIRITEL) \citep{bsb+06}. .
" We reduced and stacked the images in J. £7. and A, band using the standard pipeline."," We reduced and stacked the images in $J$, $H$, and $K_s$ band using the standard pipeline."
 IXeck and PAIRITEL fincling chart of the field is presented in Figure 1.., A Keck and PAIRITEL finding chart of the field is presented in Figure \ref{fig:finder}. .
 Astrometry was performed on all images relative to the USNO D1.0 catalog (Monetetal.2003) with (vpical lo rms relative to that catalog of 250 mas im each coordinate., Astrometry was performed on all images relative to the USNO B1.0 catalog \citep{mlc+03} with typical 1 $\sigma$ rms relative to that catalog of 250 mas in each coordinate.
" From the PAIRITEL imaging. we took note of an extended red source (67) to the south of the XRT position. at a position of a = 4455.76. 0 = 336""..7 (J2000)."," From the PAIRITEL imaging, we took note of an extended red source $G^*$ ) to the south of the XRT position, at a position of $\alpha$ = .76, $\delta$ = .7 (J2000)."
 Motivated by the inference of old galaxies at low redshilt ἐς~ 0.2) associated with some SIIDs αἱ large projected offsets (Bloom&Prochaska2006) we investigated the native of C., Motivated by the inference of old galaxies at low redshift $z \sim 0.2$ ) associated with some SHBs at large projected offsets \citep{bp06} we investigated the nature of $G^*$.
 Photometry trom the Keck data were performed using observations of the standard star field PG 2213 (Lanclolt1992)., Photometry from the Keck data were performed using observations of the standard star field PG 2213 \citep{lan92}.
. For G. we use a ((vaclius) aperture. while a smaller aperture of wwas used lor photometry of several fainter objects in and around the ART error circle.," For $G^*$, we use a (radius) aperture, while a smaller aperture of was used for photometry of several fainter objects in and around the XRT error circle."
 PAIRITEL data were photometered relative to the 2\TASS (Skrutskieetal.2006) using a rradius aperture. to capture most of the flux of G7.," PAIRITEL data were photometered relative to the 2MASS \citep{scs+06} using a radius aperture, to capture most of the flux of $G^*$."
 A summary of the photometry of this object is lound in Table 1.., A summary of the photometry of this object is found in Table \ref{tab:gphot}.
 We further investigate (he nature of G* bv fitting different profiles to our Neck imaging ol the galaxy using the software package GALFIT (Pengetal.2002)., We further investigate the nature of $G^*$ by fitting different profiles to our Keck imaging of the galaxy using the software package GALFIT \citep{phi+02}.
. Initially we fit a bulge-disk model. modeling the galaxy as a sum of an exponential profile and a cle Vaucouleurs prolile.," Initially we fit a bulge+disk model, modeling the galaxy as a sum of an exponential profile and a de Vaucouleurs profile."
 The de Vaucouleurs component was reduced to à point-source by the fit. and the residuals were very large.," The de Vaucouleurs component was reduced to a point-source by the fit, and the residuals were very large."
 The residualsfor a fit with a single. general Sérrsic profile were also unacceptable.," The residualsfor a fit with a single, general Sérrsic profile were also unacceptable."
 A signilicantly betterfit was obtained with a model of the, A significantly betterfit was obtained with a model of the
comparing the template image containing the variable aud (he corresponding CCD image of the CMD dataset.,comparing the template image containing the variable and the corresponding CCD image of the CMD dataset.
 The V and V-I information in Tables 3.4 and 5 are associated with uncertainties of 0.03 and 0.04 respectively. and are (he non-randonm errors in the zero-point determination and incorporate the errors in the OGLE calibration.," The V and V-I information in Tables 3,4 and 5 are associated with uncertainties of 0.03 and 0.04 respectively, and are the non-random errors in the zero-point determination and incorporate the errors in the OGLE calibration."
 A Lomb-Searele Periodogram (LSP) method was chosen to search for periodic variables in the final lighteurve database (Bretthorst2001)., A Lomb-Scargle Periodogram (LSP) method was chosen to search for periodic variables in the final lightcurve database \citep{Brett01}.
.. This method was produced specifically to overcome (he problem of unevenly spaced data. which in our case is caused by daylight and cloudy. nights.," This method was produced specifically to overcome the problem of unevenly spaced data, which in our case is caused by daylight and cloudy nights."
 A Fourier power spectrum is produced. with the same statistical properties as a standard power spectrum.," A Fourier power spectrum is produced, with the same statistical properties as a standard power spectrum."
 If. a periocicityv (P) is detected in the data. a spike is produced in the spectrum at a [frequency of 27/P. As implemented. any spike above 2xrms of (he power spectrum is identified as a candidate.," If a periodicity (P) is detected in the data, a spike is produced in the spectrum at a frequency of $\pi$ /P. As implemented, any spike above $\times$ rms of the power spectrum is identified as a candidate."
 Using this condition. 29.314 liehtcurves which contained such an apparent periodicity (~23% total database. ~300 times the final variable number) were found. which were then examined in detail.," Using this condition, 29,314 lightcurves which contained such an apparent periodicity $\sim$ $\%$ total database, $\sim$ 300 times the final variable number) were found, which were then examined in detail."
 Ho was noticed that the vast majoritv of this list was composed of common systematic effects inherent in the data. lighteurves [rom stars close to saturation. stars verv close to our magnitude limits and the 226 variables.," It was noticed that the vast majority of this list was composed of common systematic effects inherent in the data, lightcurves from stars close to saturation, stars very close to our magnitude limits and the $>$ $\sigma$ variables."
 All (hose candidates which exhibited a clear periodicity by eve were catalogued. leaving a sample of 100 for the variable list.," All those candidates which exhibited a clear periodicity by eye were catalogued, leaving a sample of 100 for the variable list."
 Fie.G shows our detection limits., \ref{detlim} shows our detection limits.
 The log of the total amplitude (in magnitudes) of the detected variables have been plotted against V. It is clear (hat there is a sharp cutolE to the detections. marked with a dotted line.," The log of the total amplitude (in magnitudes) of the detected variables have been plotted against V. It is clear that there is a sharp cutoff to the detections, marked with a dotted line."
 Any amplitudes which are >3% are detectable (ο V=17.0. and onlv those with an amplitude 71 mag can be found for 222.," Any amplitudes which are $>$ $\%$ are detectable to V=17.0, and only those with an amplitude $>$ 1 mag can be found for $>$ 22."
 Hence those variables which lie underneath: (his limit are missed in our dataset., Hence those variables which lie underneath this limit are missed in our dataset.
 By phase-wrapping the candidate lighteurves at the detected period. the nature of the," By phase-wrapping the candidate lightcurves at the detected period, the nature of the"
Note that we cid not always have a phase reference source within 2° of Cassini. as desired.,"Note that we did not always have a phase reference source within $2^{\circ}$ of Cassini, as desired."
 There is often a tradeoff between augular separation and reference source [lux density. aud we cid uot consider sources with f[Iux deusities less than 100 1uJy to ensure an adequate signal/noise ratio.," There is often a tradeoff between angular separation and reference source flux density, and we did not consider sources with flux densities less than 100 mJy to ensure an adequate signal/noise ratio."
 Iu retrospect this Εαν deusity cutoff μιαν have been too «'Oliservative., In retrospect this flux density cutoff may have been too conservative.
 Duriο each observing epoch we alteruated scaus |jetween Cassini and the phase refereuce source every 2-3 minutes. aud included occasional scans OL a stroug source lor instrumental delay and bandpass calibration.," During each observing epoch we alternated scans between Cassini and the phase reference source every 2-3 minutes, and included occasional scans of a strong source for instrumental delay and bandpass calibration."
 In addition. LO strong sou‘ces covering as wide an elevation range as possible were observe during a 30-10) minute period o allow a zenith troposphere delay to be estimated for each of the VLBA antenna sites.," In addition, $\sim 10$ strong sources covering as wide an elevation range as possible were observed during a 30-40 minute period to allow a zenith troposphere delay to be estimated for each of the VLBA antenna sites."
 The data [rom each epoch were stored prior to correatiou until a reconstructed orbit solution for Cassini was available from JPL., The data from each epoch were stored prior to correlation until a reconstructed orbit solution for Cassini was available from JPL.
 The Cassini orbit is determined. from Doppler tracking by the Deep Space Network. aud provides the poslious used for Cassini during correlation.," The Cassini orbit is determined from Doppler tracking by the Deep Space Network, and provides the positions used for Cassini during correlation."
 Because Cassini transmits in right circular polarization ou.y. we used ouly the R-R corralator output in our analysis.," Because Cassini transmits in right circular polarization only, we used only the R-R corralator output in our analysis."
 To avoikl decorrelation over time or [recueucy. aud to accurately measure plase across the relatively narrow bandwidth of the Cassini sigial. we used I-second integrations aud 256 spectral elianuels across each 8-MHz IE. band.," To avoid decorrelation over time or frequency, and to accurately measure phase across the relatively narrow bandwidth of the Cassini signal, we used 1-second integrations and 256 spectral channels across each 8-MHz IF band."
 The four available IF bands. were separated iu [reqiency by up to 162 MHz to improve the uiulti-baud delay response., The four available IF bands were separated in frequency by up to 462 MHz to improve the multi-band delay response.
 Sinall changes iu the relerenee frequency were necessary at each epoch to keep the Cassini signal away from the edges of IF baid 1. (, Small changes in the reference frequency were necessary at each epoch to keep the Cassini signal away from the edges of IF band 1. (
The frequency of the Cassini signal Changes by a lew MHz when switching between a oue-wa'v downlink aud a two-way coherent link.),The frequency of the Cassini signal changes by a few MHz when switching between a one-way downlink and a two-way coherent link.)
 The center lrequencies for each LF baud were approximaely 5.128. 8.500. 8.790. and 8.890 GHz.," The center frequencies for each IF band were approximately 8.428, 8.500, 8.790, and 8.890 GHz."
 Initial delays were calculated from the geometric moclel used during correlation. 1999)., Initial delays were calculated from the geometric model used during correlation \citep{Romney99}.
. The spacecraft position aud proper motjon at the start of our observations were corrected. if necessary. from the values used duriug correlation. (," The spacecraft position and proper motion at the start of our observations were corrected, if necessary, from the values used during correlation. ("
The most accurate Cassini orbit solutions were not available until a few weeks alter a particular date.),The most accurate Cassini orbit solutions were not available until a few weeks after a particular date.)
 A subtle but important part ol the geometric model is the differeuce in general relativity corrections for signals propagating through gravitational fields in the solar system. particularly those of the Sun aud Saturu.," A subtle but important part of the geometric model is the difference in general relativity corrections for signals propagating through gravitational fields in the solar system, particularly those of the Sun and Saturn."
 Unlike the phase reference sources. Cassini cau not be assume Lio be at infinity: only. part of the solar system gravitational field applies to signals [rom Cassini.," Unlike the phase reference sources, Cassini can not be assumed to be at infinity; only part of the solar system gravitational field applies to signals from Cassini."
 Amplitude calibration was based on recorded. system temperatures aloug with previously determined antenna gain curves., Amplitude calibration was based on recorded system temperatures along with previously determined antenna gain curves.
 Additional amplituce correctious were applied to remove the ellect of using two-bit quantization when recording signals at the VLBA antennas., Additional amplitude corrections were applied to remove the effect of using two-bit quantization when recording signals at the VLBA antennas.
 Phases were corrected for parallactic angle and for improvements in tle values of Earth orientation parameters, Phases were corrected for parallactic angle and for improvements in the values of Earth orientation parameters
Reeent observations have confirmed our basic understanding of cosmology ancl showed. an. impressive consistency. with the predictions of the standard. ACDAL model (e.g..Astierοἱal.2006:Percivalet2007:Ixo-matsuetal. 2009).,"Recent observations have confirmed our basic understanding of cosmology and showed an impressive consistency with the predictions of the standard $\Lambda$ CDM model \citep[e.g.,][]{SNV06,BAO07,wmap}."
. In this model. a cosmological constant dominates the cosmic mass budget today. but ealaxies and other structures were assembled earlier. primarily out of cold dark matter.," In this model, a cosmological constant dominates the cosmic mass budget today, but galaxies and other structures were assembled earlier, primarily out of cold dark matter."
 HnThis medium of non-interacting. low. velocity-clispersion paricles. started out with small Gaussian density perturbations that were subsequentIv enhanced by gravity.," This medium of non-interacting, low velocity-dispersion particles, started out with small Gaussian density perturbations that were subsequently enhanced by gravity."
 While he model successfully matches observations of the large-scale anisotropies of the cosmic microwave background and the large-scale structure in galaxy survevs. it ds also important to test its validity on smaller scales.," While the model successfully matches observations of the large-scale anisotropies of the cosmic microwave background and the large-scale structure in galaxy surveys, it is also important to test its validity on smaller scales."
 Phe abundance ane structure of non-linear objects are potentially sensiive probes of the properties of dark matter (e.g... whether it is cold) and of the density —uctuations (e.g. whether they are Gaussian).," The abundance and structure of non-linear objects are potentially sensitive probes of the properties of dark matter (e.g., whether it is cold) and of the density fluctuations (e.g., whether they are Gaussian)."
" However. o""as cooling and astrophysical feedback. complicate the interpretation of observations regarine the cark matter istribution in galaxies."," However, gas cooling and astrophysical feedback complicate the interpretation of observations regarding the dark matter distribution in galaxies."
 Thus. it is most attractive to study 1re largest virialized. objects. namely X-ray. clusters. in which most of the gas is too hot and rarefied to cool and is jus expected to trace the gravitational potential.," Thus, it is most attractive to study the largest virialized objects, namely X-ray clusters, in which most of the gas is too hot and rarefied to cool and is thus expected to trace the gravitational potential."
 The mass profiles of galaxy clusters can be measured irectly through gravitational lensing., The mass profiles of galaxy clusters can be measured directly through gravitational lensing.
 Observations of the most massive clusters now [find dozens of multiplv-imaged background. sources. allowing a precise measurement of the central 2-D mass distribution in each cluster as projected on the sky.," Observations of the most massive clusters now find dozens of multiply-imaged background sources, allowing a precise measurement of the central 2-D mass distribution in each cluster as projected on the sky."
 Also crucial for characterizing each cluster is its total virial mass. which can be measured. precisely. by supplementing the central strong lensing signal with weal: lensing distortions measured. out to the cluster edge.," Also crucial for characterizing each cluster is its total virial mass, which can be measured precisely by supplementing the central strong lensing signal with weak lensing distortions measured out to the cluster edge."
" Lt is useful to characterize the total. projected. profile with one scale. the elfective Einstein radius rg (or angle de) defined so that a circle of that radius around the cluster center contains a mean enclosed surface mass density X equal to the critical density for lensing. X—€(4xCO]DosGo, Dis). where D denotes various angular diameter distances. (Observer- Observer-Lens. ancl Lens-Source).This definition is"," It is useful to characterize the total, projected profile with one scale, the effective Einstein radius $\rE$ (or angle $\tE$ ) defined so that a circle of that radius around the cluster center contains a mean enclosed surface mass density $\bar{\Sigma}$ equal to the critical density for lensing, $\Sigma_{\rm cr}=[c^2/(4\pi G)]D_{\rm OS}/(D_{\rm OL}
D_{\rm LS})$ , where $D$ denotes various angular diameter distances (Observer-Source, Observer-Lens, and Lens-Source).This definition is"
10 AAP function were found: using an identical jackknile sampling method to that used. for the ACDAL result.,the AAP function were found using an identical jackknife sampling method to that used for the $\Lambda$ CDM result.
 1t is clear from bie., It is clear from Fig.
 7 that the measured. mean for the iree simulations agree with the respective AAP function. provided the correct expansion. history is known and that le parameter à can be determined. at cach z.," \ref{lcdminvsugra} that the measured mean for the three simulations agree with the respective AAP function, provided the correct expansion history is known and that the parameter $\alpha$ can be determined at each $z$."
 As these results show. the measurements in a AC'DAL or a dynamical lark energy. model agree very well with the predictions. if 10 correct. Cosmology is used to analyse the data.," As these results show, the measurements in a $\Lambda$ CDM or a dynamical dark energy model agree very well with the predictions, if the correct cosmology is used to analyse the data."
 For the two quintessence mocels the deviations from .XC'DM are due to the different expansion histories (see Fig. 2))., For the two quintessence models the deviations from $\Lambda$ CDM are due to the different expansion histories (see Fig. \ref{H}) ).
 This is a consistency check which confirms that the method works., This is a consistency check which confirms that the method works.
 In reality. in à galaxy survey. it is not possible to measure the parameter a accurately at high recshilts xecause of the dillieulties associated with measuring galaxy »eculiar. velocities. ο sullicient precision.," In reality, in a galaxy survey, it is not possible to measure the parameter $\alpha$ accurately at high redshifts because of the difficulties associated with measuring galaxy peculiar velocities to sufficient precision."
 We shall now degrade the status of the icdealisecl observer mentioned above and consider a more realistic observer who still knows the correct. cosmological model but who is unable Oo measure à cdirectiv at any redshift other than z=0., We shall now degrade the status of the idealised observer mentioned above and consider a more realistic observer who still knows the correct cosmological model but who is unable to measure $\alpha$ directly at any redshift other than $z=0$.
 Using the measured distribution of pairs at z=0. we fit the distribution given in Eq.," Using the measured distribution of pairs at $z=0$, we fit the distribution given in Eq."
 6. to set a and test the accuracy of he AAP function using this a(z=0) value at cach recdshilt. as suggested by Alarinoni Buzzi.," \ref{dist} to set $\alpha$ and test the accuracy of the AAP function using this $\alpha(z=0)$ value at each redshift, as suggested by Marinoni Buzzi."
 La does not evolve with redshift we would expect this approach to result. in accurate agreement between the measured. mean anc the AAP function., If $\alpha$ does not evolve with redshift we would expect this approach to result in accurate agreement between the measured mean and the AAP function.
 In Fig., In Fig.
 SN. the measured clistribution of the angle 7. in radians. for . CDM is shown as a red hashed region with error bars.," \ref{fit_alpha} the measured distribution of the angle $\tau$, in radians, for $\Lambda$ CDM is shown as a red hashed region with error bars."
 Note the y-axis shows the fraction of the total number of pairs per bin., Note the y-axis shows the fraction of the total number of pairs per bin.
 The distribution (eq. 6)), The distribution (Eq. \ref{dist}) )
 with best fitting value lor αι=5.67£0.) (with Lo errors) is plotted. as a purple dashed line., with best fitting value for $\alpha_{\mbox{\tiny FIT}} = 5.67 \pm 0.1$ (with $\sigma$ errors) is plotted as a purple dashed line.
 Phe grey dotted. lines show the distribution adopting a|le and ala., The grey dotted lines show the distribution adopting $\alpha + 1\sigma$ and $\alpha - 1\sigma$.
 Note the error we obtain for a. 0.1. is much smaller than tha obtained by Marinoni Buzzi (0.3) due to the cillerence in sample size and the dillerent methods used to estimate the errors.," Note the error we obtain for $\alpha$, 0.1, is much smaller than that obtained by Marinoni Buzzi (0.3) due to the difference in sample size and the different methods used to estimate the errors."
 This value for @ agrees with the measure value from the simulations of a=5.56., This value for $\alpha$ agrees with the measured value from the simulations of $\alpha = 5.56$.
 In Fig., In Fig.
 9. the AAP function assuming a . CDM cosmology and using this Oppp(2=O) value at each redshift is shown as a black clashec line with error bars., \ref{plot_alpha} the AAP function assuming a $\Lambda$ CDM cosmology and using this $\alpha_{\mbox{\tiny FIT}}(z=0)$ value at each redshift is shown as a black dashed line with error bars.
 Phe mismatch between this curve aux the simulation results clearly indicates that à evolves with redshift. invalidating one of the main assumptions mace in the analysis of Marinoni Buzzi.," The mismatch between this curve and the simulation results clearly indicates that $\alpha$ evolves with redshift, invalidating one of the main assumptions made in the analysis of Marinoni Buzzi."
 Note that this black dashed line in Fig., Note that this black dashed line in Fig.
 9 is much smoother than the shaded ereen. band for the AAP function in ACDAL where the value of à is measured directly from the simulation at cach redshift., \ref{plot_alpha} is much smoother than the shaded green band for the AAP function in $\Lambda$ CDM where the value of $\alpha$ is measured directly from the simulation at each redshift.
 Using the z=0 value for a produces an AAP function which is systematically and significantly below the measured results for a ACDAL cosmology for 2>0., Using the $z=0$ value for $\alpha$ produces an AAP function which is systematically and significantly below the measured results for a $\Lambda$ CDM cosmology for $z>0$.
 Applying this measure of à as proposed by Marinoni Buzzi could lead to a spurious detection of deviations from ACDAL., Applying this measure of $\alpha$ as proposed by Marinoni Buzzi could lead to a spurious detection of deviations from $\Lambda$ CDM.
 lt is clear from Fig., It is clear from Fig.
 9. that the method. proposed. by Marinoni Buzzi contains a serious svstematic error which is apparent when applied to a Large sample of pairs., \ref{plot_alpha} that the method proposed by Marinoni Buzzi contains a serious systematic error which is apparent when applied to a large sample of pairs.
 Marinoni Buzzi considered a smaller sample than in the simulations where the statistical errors dominated this svstematic., Marinoni Buzzi considered a smaller sample than in the simulations where the statistical errors dominated this systematic.
 Lt is clear from Fig., It is clear from Fig.
 9. that à does evolve with recshift and that we can improve on the estimates of this parameter bv fitting Eq., \ref{plot_alpha} that $\alpha$ does evolve with redshift and that we can improve on the estimates of this parameter by fitting Eq.
 ο to the measured. distribution. at each redshift., \ref{dist} to the measured distribution at each redshift.
 In Fig., In Fig.
 9 the AAP function ina ACDAL cosmology using the best Gt values lor à measured at cach redshift is shown as a red dot-dashed line with error bars., \ref{plot_alpha} the AAP function in a $\Lambda$ CDM cosmology using the best fit values for $\alpha$ measured at each redshift is shown as a red dot-dashed line with error bars.
 The jackknife errors on o are estimated using 100. subsamples for the distributions at z=01 and 50 subsamples for z=1.5 and 2 as there are fewer pairs at these higher redshifts., The jackknife errors on $\alpha$ are estimated using 100 subsamples for the distributions at $z=0 - 1$ and 50 subsamples for $z=1.5$ and 2 as there are fewer pairs at these higher redshifts.
 This approach to measuring à gives much better agreement with the mean measured. from the simulations. shown as yellow data points in Fi; 9.," This approach to measuring $\alpha$ gives much better agreement with the mean measured from the simulations, shown as blue-yellow data points in Fig. \ref{plot_alpha}."
 Note this method of extracting o assumes that the correct cosmology is (CDM., Note this method of extracting $\alpha$ assumes that the correct cosmology is $\Lambda$ CDM.
 In this section we are no longer idealised. observers. who know the correct. cosmology. so the only possible choice is to assume the same cosmology in the data fitting and in the theoretical prediction. of the galaxy distribution.," In this section we are no longer idealised observers who know the correct cosmology, so the only possible choice is to assume the same cosmology in the data fitting and in the theoretical prediction of the galaxy distribution."
 Specifically we will assume CDM. for simplicity. in order to set the expansion history Z/(2) in Eq.," Specifically we will assume $\Lambda$ CDM, for simplicity, in order to set the expansion history $H(z)$ in Eq."
 S and to compute the comoving distances in Eq. 1..," 8 and to compute the comoving distances in Eq. \ref{sin2t},"
 as well as to extract the parameter o., as well as to extract the parameter $\alpha$.
 In order to find the parameter à we must fit to the observed. distribution of the orientations of pairs which has been found also by assuming a ACDAL cosmology., In order to find the parameter $\alpha$ we must fit to the observed distribution of the orientations of pairs which has been found also by assuming a $\Lambda$ CDM cosmology.
 Assuming that the true cosmological model chosen by nature is à cvnamical clark energy model. for instance the INV or SUGRA cosmology. we will check if the wrong cosmology. ACDALin our case. can be excluded or not. and consequentIv if the method. is applicable to future galaxy surveys.," Assuming that the true cosmological model chosen by nature is a dynamical dark energy model, for instance the INV or SUGRA cosmology, we will check if the wrong cosmology, $\Lambda$ CDM in our case, can be excluded or not, and consequently if the method is applicable to future galaxy surveys."
 For this analysis we take subhalo pairs in the INV and οἱολ simulations and at cach redshift we rescale the comoving, For this analysis we take subhalo pairs in the INV and SUGRA simulations and at each redshift we rescale the comoving
wavelength given for each sample member in Table 2.,wavelength given for each sample member in Table 2.
" As with the FUSE data, the instrumental spectral resolution depends on the angular size of the far-UV continuum within the COS aperture."," As with the FUSE data, the instrumental spectral resolution depends on the angular size of the far-UV continuum within the COS aperture."
" Our HST UV images (O09) yield sizes of 0.1 to 0.5 arcsec (FWHM), and we estimate that the resulting spectral resolution is R~ 3000 to 13,000, corresponding to a velocity dispersion of o~ 10 to 43 km s! (similar to the FUSE data)."," Our HST UV images (O09) yield sizes of 0.1 to 0.5 arcsec (FWHM), and we estimate that the resulting spectral resolution is $R \sim$ 3000 to 13,000, corresponding to a velocity dispersion of $\sigma \sim$ 10 to 43 km $^{-1}$ (similar to the FUSE data)."
 For the new COS spectra we have fit each merged G130M and G160M spectrum with a Starburst99 et al., For the new COS spectra we have fit each merged G130M and G160M spectrum with a Starburst99 (Leitherer et al.
 1999) model spectrum for a young stellar (Leithererpopulation with a metallcity similar to that of the corresponding LBA (Table 2)., 1999) model spectrum for a young stellar population with a metallcity similar to that of the corresponding LBA (Table 2).
" We did not attempt to obtain a rigorous best-fit, but simply adopted a standard Chabrier (2003) IMF and assumed a constant rate of star-formation."," We did not attempt to obtain a rigorous best-fit, but simply adopted a standard Chabrier (2003) IMF and assumed a constant rate of star-formation."
 We varied the duration of the star-formation until we obtained (by eye) a satisfactory fit to the strong stellar wind features., We varied the duration of the star-formation until we obtained (by eye) a satisfactory fit to the strong stellar wind features.
 We then took the ratio of the model and the data and fit the result with a low- polynomial (avoiding spectral regions with strong emission or absorption features)., We then took the ratio of the model and the data and fit the result with a low-order polynomial (avoiding spectral regions with strong emission or absorption features).
 We then divided the data by the polynomial to obtain a normalized spectrum., We then divided the data by the polynomial to obtain a normalized spectrum.
 An example of this is shown in Figure 2., An example of this is shown in Figure 2.
 The opacity of the interstellar medium to Lyman continuum photons is due to two primary sources: the photoelectric opacity of the neutral hydrogen and the opacity of dust., The opacity of the interstellar medium to Lyman continuum photons is due to two primary sources: the photoelectric opacity of the neutral hydrogen and the opacity of dust.
 The latter significantly affects the entire UV continuum in typical local starbursts and high-z galaxies., The latter significantly affects the entire UV continuum in typical local starbursts and high-z galaxies.
 We will focus first on measuring constraints on the photoelectric opacity., We will focus first on measuring constraints on the photoelectric opacity.
 Later we will comment on the importance of dust., Later we will comment on the importance of dust.
 'The only direct probe of the photoelectric opacity is the direct measurement of the relative intensity in the continuum shortward of the Lyman edge., The only direct probe of the photoelectric opacity is the direct measurement of the relative intensity in the continuum shortward of the Lyman edge.
 The FUSE data provide clean access to this spectral region and adequate signal-to-noise in only 5 targets., The FUSE data provide clean access to this spectral region and adequate signal-to-noise in only 5 targets.
 G09 have used these data to derive upper limits to the escape of ionizing radiation for these galaxies (see also Deharveng et al., G09 have used these data to derive upper limits to the escape of ionizing radiation for these galaxies (see also Deharveng et al.
 2001; Bergvall et al., 2001; Bergvall et al.
 2006)., 2006).
 The COS data do not sample this spectral region in any galaxy., The COS data do not sample this spectral region in any galaxy.
" Thus, to gain insight into the escape of ionizing radiation for our two samples we will follow H01 and G09 and use as probes the residual intensity in the core of the strongest interstellar absorption lines that primarily arise in the neutral interstellar medium."," Thus, to gain insight into the escape of ionizing radiation for our two samples we will follow H01 and G09 and use as probes the residual intensity in the core of the strongest interstellar absorption lines that primarily arise in the neutral interstellar medium."
 For the FUSE data the strongest such line is CIIA1036.3 and for the HST-COS data the strongest such line is CIIA1334.5., For the FUSE data the strongest such line is $\lambda$ 1036.3 and for the HST-COS data the strongest such line is $\lambda$ 1334.5.
 Thus in both samples we will use the same ionic species as a tracer of the neutral ISM., Thus in both samples we will use the same ionic species as a tracer of the neutral ISM.
" We have measured the relative residual intensity in the core of the CII absorption line, following the same methodology as in H01 and G09."," We have measured the relative residual intensity in the core of the CII absorption line, following the same methodology as in H01 and G09."
 The results are listed in Tables 1 (FUSE) and 2 (COS)., The results are listed in Tables 1 (FUSE) and 2 (COS).
" Before considering constraints on the escape of ionizing radiation, it is important to establish that an ionizing population of massive stars is in fact present in our targets."," Before considering constraints on the escape of ionizing radiation, it is important to establish that an ionizing population of massive stars (O stars) is in fact present in our targets."
 The strongest (Ospectroscopicstars) signature of these stars is provided by the broad P-Cygni profiles in the resonance transitions of highly ionized species that arise in the stellar winds (e.g. Robert et al., The strongest spectroscopic signature of these stars is provided by the broad P-Cygni profiles in the resonance transitions of highly ionized species that arise in the stellar winds (e.g. Robert et al.
 1993; 2003)., 1993; 2003).
 'The strongest such feature that is present in all our COS spectra is NVA1240., The strongest such feature that is present in all our COS spectra is $\lambda$ 1240.
 In Figure 3 we show the NV, In Figure 3 we show the NV
pixel and a spatial scale of iucesec/pixel eave a spectral resolutiou ofAA.,/pixel and a spatial scale of arcsec/pixel gave a spectral resolution of.
 Initial data reduction used standard techniques for bias removal. flat fielding. aud optimal extraction of the spectra.," Initial data reduction used standard techniques for bias removal, flat fielding, and optimal extraction of the spectra."
 The slit was rotated to include another star on the slit. aud this was used for both relative flux calibration and relative wvavelensth calibration.," The slit was rotated to include another star on the slit, and this was used for both relative flux calibration and relative wavelength calibration."
 Owing to techuical difüculties. we were unable to iutersperse the target observations with are calibrations and so wavelength calibration was done only using HeNeÀr arcs from the beeiuuiug of the first wieght. with subsequent offsets derived using absorption lines in the spectrum of the comparison star.," Owing to technical difficulties, we were unable to intersperse the target observations with arc calibrations and so wavelength calibration was done only using HeNeAr arcs from the beginning of the first night, with subsequent offsets derived using absorption lines in the spectrum of the comparison star."
 The comparison star was calibrated relative to the spectrophotometric standard EC21 (Ibunuxetal.1992.1991) using au observation taken at low airuass with a wide-slit.," The comparison star was calibrated relative to the spectrophotometric standard EG21 \citep{Hamuy:1992a,Hamuy:1994a} using an observation taken at low airmass with a wide-slit."
 All object spectra were then calibrated relative to this., All object spectra were then calibrated relative to this.
 The mean spectrum frou both uights of observation is shown in Figures 1 and 2.., The mean spectrum from both nights of observation is shown in Figures \ref{fig:compcalspec} and \ref{fig:comnorspec}.
 Optical spectroscopy was also obtained on 2003 February 7 aud 2003 February 28 with the Verv Large Telescope (VET) at Cerro Paranal. Chile.," Optical spectroscopy was also obtained on 2003 February 7 and 2003 February 28 with the Very Large Telescope (VLT) at Cerro Paranal, Chile."
 The first night of observation used the FORS2 spectrograph and 1100V erating with waveleneth coverage aand mean dispersion of /pixel: the second nieht the 600 RI erating with ccoverage and mean dispersion of 1.62 AA//pixcl., The first night of observation used the FORS2 spectrograph and 1400V grating with wavelength coverage and mean dispersion of /pixel; the second night the 600 RI grating with coverage and mean dispersion of $1.62$ /pixel.
 Both sets of observations used a slit width of 0.1 arcsec. παπάς slit losses sieuificant aud. absolute Hux calibration unreliable.," Both sets of observations used a slit width of 0.4 arcsec, making slit losses significant and absolute flux calibration unreliable."
 With ao spatial scale. of 4.125 arcsec/pixel this gave a spectral resolution of for night 1 aud for nieht 2., With a spatial scale of 0.125 arcsec/pixel this gave a spectral resolution of for night 1 and for night 2.
 Tuitial data reduction was again carried out usingLRAL., Initial data reduction was again carried out using.
 Ie | Ne aud IIgCd arc lines were obtained during the dav vefore and after the observations aud interpolated through he night., He + Ne and HgCd arc lines were obtained during the day before and after the observations and interpolated through the night.
 The mean spectrum frou both uights of observation is shown in Figure 2.., The mean spectrum from both nights of observation is shown in Figure \ref{fig:comnorspec}.
 Two spectra were obtained with the 6.5 im Walter Baade. Magellan telescope at Las Campanas Observatory. Cerro. Aanqui. Chile on 2003 December 11/15.," Two spectra were obtained with the 6.5 m Walter Baade, Magellan telescope at Las Campanas Observatory, Cerro Manqui, Chile on 2003 December 14/15."
 The IMACS instrmucut was used in f/l1 configuration with a 0.7 aresec slit aud an skxak CCD mosaic detector operating iu 242 binning mode., The IMACS instrument was used in f/4 configuration with a 0.7 arcsec slit and an $\times$ 8k CCD mosaic detector operating in $\times$ 2 binning mode.
 The low-resolution spectrin was obtained using the 300 erating with a rauge of aand a AA/pixel mage scale giving iresolution., The low-resolution spectrum was obtained using the 300 grating with a range of and a A/pixel image scale giving resolution.
 The ligher-resolution spectrum was obtained using the 600 erating with a waveleneth range ofAA... a AA//pixcl image scale and iresolution.," The higher-resolution spectrum was obtained using the 600 grating with a wavelength range of, a /pixel image scale and resolution."
 Data reduction was carried out usine PAMELA software., Data reduction was carried out using software.
 Both spectra are plotted in Figure 2.., Both spectra are plotted in Figure \ref{fig:comnorspec}.
 Three spectra of wwere obtained using the low-resolutiou short-wavelcueth prine camera (SWDP) in 19901991., Three spectra of were obtained using the low-resolution short-wavelength prime camera (SWP) in 1990–1991.
 As these have not been published previously. we extracted. the reprocessed spectra (Nichols&Linsky1906) from the archive for comparison with our new data.," As these have not been published previously, we extracted the reprocessed spectra \citep{Nichols:1996a} from the archive for comparison with our new data."
 We constructed. an exposure time weighted average., We constructed an exposure time weighted average.
 This is based on ~5 binary orbits aud so should adequately represcut the mean spectrmn in 19901901., This is based on $\sim5$ binary orbits and so should adequately represent the mean spectrum in 1990–1991.
 The extracted spectra are plotted alongside our sspectra for comparison in Figure 3..Pa, The extracted spectra are plotted alongside our spectra for comparison in Figure \ref{fig:meaniue}.
 A broad range of cinission lines are seen in the data. saluplue a rauge of ionization stages aud excitations.," A broad range of emission lines are seen in the data, sampling a range of ionization stages and excitations."
 These are παλκος in Table 2.., These are summarised in Table \ref{tab:lines}.
 Iu the optical. Daliner lues of IIT are prominent alongside both Wel aud ΠΟΠ.," In the optical, Balmer lines of HI are prominent alongside both HeI and HeII."
 The Bowen bleud of NITE and/or CTI is clearly seeu as iu most LAINB spectra., The Bowen blend of NIII and/or CIII is clearly seen as in most LMXB spectra.
 The UV spectra is also typical for an LAINB. being dominated by the stroug lines of CIV aud NY. alongside CIIT. NIV. OIV. aud OV.," The UV spectrum is also typical for an LMXB, being dominated by the strong lines of CIV and NV, alongside CIII, NIV, OIV, and OV."
 The line streugtlis seen in the £u-UV show no obvious anomalies and the ratio of CIV to NV is typical: there is no indication of substantial CNO processing as seen in XTE J1118|£80. for example (aswellctal.2002).," The line strengths seen in the far-UV show no obvious anomalies and the ratio of CIV to NV is typical; there is no indication of substantial CNO processing as seen in XTE J1118+480, for example \citep{haswell02}."
. Iu addition. to ciission. there are also absorption features.," In addition to emission, there are also absorption features."
 Broad Lya absorption is common in all Ελ». and may arise from the absorption iu the disk atinospliere. and/or the interstellar medimm.," Broad $\alpha$ absorption is common in all LMXBs, and may arise from the absorption in the disk atmosphere, and/or the interstellar medium."
 Mauy metallic resonance absorption lines iu the £u-UV are also ofinterstellar origin., Many metallic resonance absorption lines in the far-UV are also of interstellar origin.
 The Na D lines ave ouly partially resolved iu the CTIO data and also blended with ΠΟ limiting the precision with which their individual strengths can be measured., The Na D lines are only partially resolved in the CTIO data and also blended with HeI limiting the precision with which their individual strengths can be measured.
 The total equivalent width is ~0.6AA.. with the D» line stronger than D4.," The total equivalent width is $\sim0.6$, with the $_2$ line stronger than $_1$."
 Assmuing a ratio between Lil and 1:2. for D4:D». implies a Do EW ofAA.. and reddening E(BV) likely in the rauge 0.07 Amari&Zwitter 1997).," Assuming a ratio between 1:1 and 1:2, for $_1$ $_2$, implies a $_2$ EW of, and reddening $E(B-V)$ likely in the range 0.07--0.24 \citep{munari97}."
. This is comparable to. but somewhat larger than. that deduced from the interstellar line iu the forthcoming Paper ITI (Ives et al..," This is comparable to, but somewhat larger than, that deduced from the interstellar line in the forthcoming Paper III (Hynes et al.,"
 in prop. LEBWV)20.06£0.03.," in prep.), $E(B-V)=0.06\pm0.03$."
 The latter is probably a more reliable indicator., The latter is probably a more reliable indicator.
 Both are markedly lower than the frequently quoted value of E(BV)=0.12£0.03 (Schoenibs&Zoeschiuger1990).. but this was based on au reliable method for an LAND. so our lower values are more credible.," Both are markedly lower than the frequently quoted value of $E(B-V)=0.42\pm0.03$ \citep{schoembs90}, but this was based on an unreliable method for an LMXB, so our lower values are more credible."
The final absorption components worthyof mention are moving features. present in the Balmer Hines.,"The final absorption components worthyof mention are moving features, present in the Balmer lines,"
"Theoretically, it is only now becomiue possible to reliably iuvestigate the dependence of fyas Ol temperature with the iuclusiou of radiative cooling. star formation. feedback. and other relevant processes i numerical siauulatious (e.g... Lewis et al.","Theoretically, it is only now becoming possible to reliably investigate the dependence of $f_{gas}$ on temperature with the inclusion of radiative cooling, star formation, feedback, and other relevant processes in numerical simulations (e.g., Lewis et al."
 2000: Pearce et al., 2000; Pearce et al.
 2000: Abluunvous et al., 2000; Muanwong et al.
 2001: Davé et al., 2001; Davé et al.
 2001)., 2001).
 As of vet. however. there is little aegreciment iu the approaches adopted to model these processes aud prevent the so-called cooling crisis (compare. for example. the findings of Lewis et al.," As of yet, however, there is little agreement in the approaches adopted to model these processes and prevent the so-called cooling crisis (compare, for example, the findings of Lewis et al."
 2000 with those of Pearce et al., 2000 with those of Pearce et al.
 2000)., 2000).
 This is not surprising., This is not surprising.
 As discussed in detail by Balogh ct al. (, As discussed in detail by Balogh et al. (
2001). attempting to model the effects of cooling across the wide range of halo masses found in clusters is inhereutly very diffieult.,"2001), attempting to model the effects of cooling across the wide range of halo masses found in clusters is inherently very difficult."
 The addition of “sub-erid™ processes. such as star formation aud feedback. further complicates matters.," The addition of “sub-grid” processes, such as star formation and feedback, further complicates matters."
 Thus. the effects that these additional pliysical processes have ou the eas mass fraction of clusters will not be fully realized until such issues are resolved.," Thus, the effects that these additional physical processes have on the gas mass fraction of clusters will not be fully realized until such issues are resolved."
 Iu this paper. however. we show that the observed variation of the VyasTy velatiou(s) arises quite naturally within the class of models that invoke preheating of the intracluster medium durius the carly stages of cluster formation.," In this paper, however, we show that the observed variation of the $M_{gas} 
- T_X$ relation(s) arises quite naturally within the class of models that invoke preheating of the intracluster medium during the early stages of cluster formation."
" Iu these iodels. fj,; IS coustaut ou cluster scales (Fy23 keV). and the self-similarity is instead broken bv an cutropy floor generated bv carly non-eravitational heating eveuts."," In these models, $f_{gas}$ is constant on cluster scales $T_X \gtrsim 3$ keV), and the self-similarity is instead broken by an entropy floor generated by early non-gravitational heating events."
 Preheating has previously been shown to bring consistency between a number of other observed and predicted scaling relations for eroups and clusters (e.¢.. BBLP02). and therefore one might expect that the MfgasTx relation should also be modified.," Preheating has previously been shown to bring consistency between a number of other observed and predicted scaling relations for groups and clusters (e.g., BBLP02), and therefore one might expect that the $M_{gas} 
- T_X$ relation should also be modified."
 The preheating model was originally put forward by Isaiscr (1991) and has subsequently been iuvestigated by a nunuber of authors (e.g... Evrard Temy 1991. Bower 1997. Cavaliere et al.," The preheating model was originally put forward by Kaiser (1991) and has subsequently been investigated by a number of authors (e.g., Evrard Henry 1991, Bower 1997, Cavaliere et al."
 1997: 1998: 1999: Dalogh ct al., 1997; 1998; 1999; Balogh et al.
 1999. Wi. Fabian. Nulseun 2000: Loewenstein 2000. Tozzi Novinan 2001: Dorgani ct al.," 1999, Wu, Fabian, Nulsen 2000; Loewenstein 2000, Tozzi Norman 2001; Borgani et al."
 2001: Thomas ct al., 2001; Thomas et al.
 2002: DDLDP02)., 2002; BBLP02).
 If the ICAL is injected with enough thermal euergv. the hot. N-orav. ciuitting eas will become decoupled from the dark halo potential aud break the selfsimular scaling relations.," If the ICM is injected with enough thermal energy, the hot X-ray emitting gas will become decoupled from the dark halo potential and break the self-similar scaling relations."
 The best estimates sugecst that a substantial amount of euerev ( 1 keV per particle] is required to reproduce the observed relatious (mainly the LyTy relation)., The best estimates suggest that a substantial amount of energy $\sim$ 1 keV per particle) is required to reproduce the observed relations (mainly the $L_X - T_X$ relation).
 It is not vet known what source(s) could inject such a large amount of energy iuto the ICAL, It is not yet known what source(s) could inject such a large amount of energy into the ICM.
 Both ealactic winds (driven by supernovac) and ejecta frou active galactic nuclei have been proposed. but because of the complexity of the plivsics. the exact details have vot to be worked out.," Both galactic winds (driven by supernovae) and ejecta from active galactic nuclei have been proposed, but because of the complexity of the physics, the exact details have yet to be worked out."
 For an in-depth discussion of poteutial sources of preheating aud of alternative possibilities for reproducing the observed relations we refer the reader to DDBLDP02., For an in-depth discussion of potential sources of preheating and of alternative possibilities for reproducing the observed relations we refer the reader to BBLP02.
" Iu this paper. we adopt the phvsicallv motivated. analytic model developed in BBEÉPU2 to explore the inpact of cluster preheating ou the M,Ty rolation."," In this paper, we adopt the physically motivated, analytic model developed in BBLP02 to explore the impact of cluster preheating on the $M_{gas} - T_X$ relation."
 In conrparisou with the LyFy aud ALFy relations. it has drawn very little attention by theoretical studies.," In comparison with the $L_X - T_X$ and $M - T_X$ relations, it has drawn very little attention by theoretical studies."
 The oulv studies to have examined the effects of cutropy injection ou the Ayo.Ly relation to date are Locweusteia (2000) and Dialek et al. (," The only studies to have examined the effects of entropy injection on the $M_{gas} - 
T_X$ relation to date are Loewenstein (2000) and Bialek et al. ("
2001).,2001).
 To be specific. Loewensteiu (20003. considered inodels where the eutropy injection occurs at the ceuters of eroups and clusters. after the latter have formed whereas Bialels et al (2001). like DBLP02. investigated preheated models.," To be specific, Loewenstein (2000) considered models where the entropy injection occurs at the centers of groups and clusters, after the latter have formed whereas Bialek et al (2001), like BBLP02, investigated preheated models."
 For reasons that will be described below Gn $85). we believe our work ercatly inproves upon both of these studies.," For reasons that will be described below (in 5), we believe our work greatly improves upon both of these studies."
 The prevailing apathy bv theorists is perhaps due in part to a near absence of published observational studies ou the gas masses of clusters., The prevailing apathy by theorists is perhaps due in part to a near absence of published observational studies on the gas masses of clusters.
 Wowever. iu light of the recent müiportaut observations discussed. above (e.g... Vikbhliniu et al.," However, in light of the recent important observations discussed above (e.g., Vikhlinin et al."
 1999: AIMEO99: Neuimiaun Arnaud 2001) aud the new iufiux of hieh resolution data from the aud N-rav satellites. which will likely provide even tighter constraints. we believe that a thorough examination of the MoasTx relation is timely.," 1999; MME99; Neumann Arnaud 2001) and the new influx of high resolution data from the and X-ray satellites, which will likely provide even tighter constraints, we believe that a thorough examination of the $M_{gas} 
- T_X$ relation is timely."
" The models we cousider below were developed im a flat A-CDAL cosmology with ο=0.3. h=0.75. aud a uucleosvuthesis value 9,=(0.0195? (Dues Tytler 1998)."," The models we consider below were developed in a flat $\Lambda$ -CDM cosmology with $\Omega_m = 0.3$, $h = 0.75$, and a nucleosynthesis value $\Omega_b = 0.019 h^{-2}$ (Burles Tytler 1998)."
 They are computed for a ummber of differcut preheating levels. corresponding to eutropy constants of Ay = 100. 200. 300. 127 keV οι”.," They are computed for a number of different preheating levels, corresponding to entropy constants of $K_0$ = 100, 200, 300, 427 keV $^2$."
 These span the range required to match the observed Ly relatious of eroups and hot clusters (e.g. Pomman ct al.," These span the range required to match the observed $L_X - T_X$ relations of groups and hot clusters (e.g., Ponman et al."
 1999: Llovd-Davies et al., 1999; Lloyd-Davies et al.
 2000: Tozzi Norman 2001: BBLP02)., 2000; Tozzi Norman 2001; BBLP02).
 For the purposes of comparison. we also implement au “isothermal” model (see Section 2.3 iu BBLPO2). which nündücs the sclfsimilar result deduced from nmuuerncal simulations (c.e.. Evrard et al.," For the purposes of comparison, we also implement an “isothermal” model (see Section 2.3 in BBLP02), which mimics the self-similar result deduced from numerical simulations (e.g., Evrard et al."
 1996)., 1996).
 Since an in-depth discussion of the preheated cluster uodels can be found iu BBLPO2. we present oulv a brief description of the models here.," Since an in-depth discussion of the preheated cluster models can be found in BBLP02, we present only a brief description of the models here."
 The preheated models can be παπασος. as follows: he donunaut dark matter component. which is unaffected w the cherev injection. collapses and virializes to form und halos.," The preheated models can be summarized as follows: the dominant dark matter component, which is unaffected by the energy injection, collapses and virializes to form bound halos."
 The distribution of the dark matter im such alos is asstuned to be the same as found in receut ultra- resolution nunerical simulations (Moore ct al., The distribution of the dark matter in such halos is assumed to be the same as found in recent ultra-high resolution numerical simulations (Moore et al.
 1998: IxIvpiu et al., 1998; Klypin et al.
 1999: Lewis et al., 1999; Lewis et al.
 2000) aud is described by where n — Lhe pany is the profile normalization. aud ris the scale radius.," 2000) and is described by where n = 1.4, $\rho_{dm,0}$ is the profile normalization, and $r_s$ is the scale radius."
 While the dark compoucut is unaffected by energy injection. the collapse of the barvouic component is hindered by the pressure forces induced by preheating.," While the dark component is unaffected by energy injection, the collapse of the baryonic component is hindered by the pressure forces induced by preheating."
 If the maxi iufall velocity duc purely to eravitv of the dark halo is subsonic. the flow will be strouely affected by the pressure and it will not muclereo accretion shocks.," If the maximum infall velocity due purely to gravity of the dark halo is subsonic, the flow will be strongly affected by the pressure and it will not undergo accretion shocks."
 It is assumed that the barvous will accumulate outo the halostsentropically at the adiabatic Boudi accretion rate (as described in Balogh et al., It is assumed that the baryons will accumulate onto the halos at the adiabatic Bondi accretion rate (as described in Balogh et al.
 1999)., 1999).
 This treatineut. however. is only appropriate for low mass halos.," This treatment, however, is only appropriate for low mass halos."
 If the exavitv of the dark halos is strong enough (as it is expected to be in the hot clusters being considered here) so that the maxima mfall velocity is transonic or supersonic. the eas will experience an additional (generally dominant) cutropy increase due to accretion shocks.," If the gravity of the dark halos is strong enough (as it is expected to be in the hot clusters being considered here) so that the maximum infall velocity is transonic or supersonic, the gas will experience an additional (generally dominant) entropy increase due to accretion shocks."
 Iu order to trace the shock history of the eas. a detailed kuowledge of the merger history of the cluster/seroup is required but is not considered by BDBLDPU2.," In order to trace the shock history of the gas, a detailed knowledge of the merger history of the cluster/group is required but is not considered by BBLP02."
 ILustead. it Is asstuned that at some earlier time the most massive cluster progenitor will have had a mass low euoueh such that shocks were ueglieible in its formation. simular to the," Instead, it is assumed that at some earlier time the most massive cluster progenitor will have had a mass low enough such that shocks were negligible in its formation, similar to the"
" This work is part of the Helmholtz Alliance's ""Planetary evolution and Life’, which RB and AM thank for financial support."," This work is part of the Helmholtz Alliance's 'Planetary evolution and Life', which RB and AM thank for financial support."
 Exchanges between Nice and Thessaloniki have been funded by a PICS programme of France's CNRS., Exchanges between Nice and Thessaloniki have been funded by a PICS programme of France's CNRS.
 HFL is grateful to NASA's Origins of Solar Systems and Outer Planets Program for funding., HFL is grateful to NASA's Origins of Solar Systems and Outer Planets Program for funding.
" Allégere, C. J.. Manhess. G.. Gopel. 22008, Earth and Planetary Science Letters, 267. 386 Bailey. B. L.. Malhotra, 22009, Icarus. 203, 155 Batygin, K.. Brown. M. 22010,ApJ.. 716, 1323"," Allèggre, C. J., Manhèss, G., Gopel, 2008, Earth and Planetary Science Letters, 267, 386 Bailey, B. L., Malhotra, 2009, Icarus, 203, 155 Batygin, K., Brown, M. 2010, 716, 1323"
On 2010 July 27. a microflare. with the Solar Object Locator of SOL2010-07-27T0s:16:00L223C108. occurred to the west side of NOAA active region (AR) 11089 (S21W21).,"On 2010 July 27, a microflare, with the Solar Object Locator of SOL2010-07-27T08:46:00L223C108, occurred to the west side of NOAA active region (AR) 11089 (S24W21)."
 The microllare was located at S18W[8 in the heliocentric coordinates., The microflare was located at S18W48 in the heliocentric coordinates.
 The GOES 1-3 soft X-ray lighteurve shows a sinall Dump less than A2 level from 08:16 UT to 08:56 UT., The GOES 1-8 soft X-ray lightcurve shows a small bump less than A2 level from 08:46 UT to 08:56 UT.
 Accompanying the occurrence of the tiny flare. EUV waves and dinmings cau be ideutified iu tlie base difference images. extencding from the flare site to the north lor a short distance.," Accompanying the occurrence of the tiny flare, EUV waves and dimmings can be identified in the base difference images, extending from the flare site to the north for a short distance."
 No coronagraph observatious from theSOHRO satellite are available., No coronagraph observations from the satellite are available.
 The microflare was located at S22E31 iu the field of view of theObservatory A. and the its coronagraphs did uot detect any CME during this event.," The microflare was located at S22E31 in the field of view of the A, and the its coronagraphs did not detect any CME during this event."
 Two factors do not favor the detection of a possible CME for this event: oue is that the source region is too close to the solar disk center for theSTEREO satellite. aud the other is that the CME. if existing.D> should have a small angular extension as üunplied by the EUV clitmmines. aud therefore is prone to be missed by coronagraplis 2010).," Two factors do not favor the detection of a possible CME for this event: one is that the source region is too close to the solar disk center for the satellite, and the other is that the CME, if existing, should have a small angular extension as implied by the EUV dimmings, and therefore is prone to be missed by coronagraphs ."
. The microllare and the associated “EIT wave” were well documented by the Atmospheric Imaging Assembly (ALA) on theObservatory SDO)mainBodyCitationEnd673][titlhoQ06.," The microflare and the associated “EIT wave"" were well documented by the Atmospheric Imaging Assembly (AIA) on the ."
 The AIA instrument has 10 EUV and UV channels with a spatial resolution of 1.72., The AIA instrument has 10 EUV and UV channels with a spatial resolution of $1.\arcsec 2$.
 Since “EIT waves” are most evident at the 193 channel. we mainly use this channel to study the dynatics of the “EIT wave’.," Since “EIT waves"" are most evident at the 193 channel, we mainly use this channel to study the dynamics of the “EIT wave""."
 The cadeuce of the observation is 12 s. which allows the detection of the detailed features of “EIT waves” that had not been seen before2010).," The cadence of the observation is 12 s, which allows the detection of the detailed features of “EIT waves"" that had not been seen before."
. Iu order to clearly slow the propagation of the faint “EIT waves”. runuing or base clifference techuique was often used.," In order to clearly show the propagation of the faint “EIT waves"", running or base difference technique was often used."
 Throughout this paper. we rely ou the percentage clillerence images. which are obtained by the base ciffereuce tages divided by the pre-event image1999).," Throughout this paper, we rely on the percentage difference images, which are obtained by the base difference images divided by the pre-event image."
. The rotation of the Sun is corrected., The rotation of the Sun is corrected.
 For simplicity. the percentage clillereuce linages are meutioned as difference images or difference intensity hereafter.," For simplicity, the percentage difference images are mentioned as difference images or difference intensity hereafter."
 The evolution of the “EIT wave” event ou 2010 July 27 is displayed in Figure L.. which shows he 193A difference images at [our moments.," The evolution of the “EIT wave"" event on 2010 July 27 is displayed in Figure \ref{fig1}, which shows the 193 difference images at four moments."
 The pre-event intensity map at 08:11:55 UT is chosen as the base image that is subtracted., The pre-event intensity map at 08:44:55 UT is chosen as the base image that is subtracted.
 In this figure. the bright. (clark) pixels tuclicate the intensity increase (decrease).," In this figure, the bright (dark) pixels indicate the intensity increase (decrease)."
" Bt is seen that at 08:10:13. UT. a small-sized brighteniug occurred 1θα(670. —350"")."," It is seen that at 08:46:43 UT, a small-sized brightening occurred near, $-350\arcsec$ )."
" Note that a small coronal loop with relatively stronger magnetic field was o the north at around(720"".. —200""). which is clearly seen in the second panel."," Note that a small coronal loop with relatively stronger magnetic field was to the north at around, $-200\arcsec$ ), which is clearly seen in the second panel."
 Shortly later. wieght frouts appeared and began to propagate to the north.," Shortly later, bright fronts appeared and began to propagate to the north."
 It is revealed from the attached movie (basediff.inpg) that the western aud eastern flanks Dormecd first., It is revealed from the attached movie ) that the western and eastern flanks formed first.
 At 08:19:55 UT. the northern flank jecame to be domiuanut iu brightuess.," At 08:49:55 UT, the northern flank became to be dominant in brightness."
 Until this stage. only one bright strip was evident along the iorthern Hank.," Until this stage, only one bright strip was evident along the northern flank."
 At 08:53:31 UT. the northern flauk of the brighteniugs swept over the sinall coronal," At 08:53:31 UT, the northern flank of the brightenings swept over the small coronal"
"A joint fit to both datasets reveals e = icon V = 344 and 42, 2 49 for 33 RV points and 5 filled parameters (including two systemate velocities, ?sopiju and τα). indicaüng the eccentricily 1s non-zero at the 2.76 level.","A joint fit to both datasets reveals $e$ = $^{+0.022}_{-0.020}$, $\chi^{2}$ = 34.4 and $\chi^{2}_{\rm circ}$ = 49 for 33 RV points and 5 fitted parameters (including two systematic velocities, $\gamma_{\rm SOPHIE}$ and $\gamma_{\rm FIES}$ ), indicating the eccentricity is non-zero at the $\sigma$ level."
" We also apply the test to the TIO Keck RVs, including their 7.3 m ! jitter term."," We also apply the test to the T10 Keck RVs, including their 7.3 m $^{-1}$ jitter term."
" We find the eccentricity is detected at the 4o level, € = 0.104 + 0.012, with V = 12.3. AZ.> = 83.9. 14 RVs and 4 fitted parameters."," We find the eccentricity is detected at the $\sigma$ level, $e$ = 0.104 $\pm$ 0.012, with $\chi^{2}$ = 12.3, $\chi^{2}_{\rm circ}$ = 83.9, 14 RVs and 4 fitted parameters."
" By combining all three datasets, our analysis suggests that the orbit is not circular with a 60 confidence. e = 0.095 + 0.011 (4? 2 55. 42cue = 133. 47 RVs and 6 fitted parameters (including three systematic velocities, ?xoppu. SELES and o)."," By combining all three datasets, our analysis suggests that the orbit is not circular with a $\sigma$ confidence, $e$ = 0.095 $\pm$ 0.011 $\chi^{2}$ = 55, $\chi^{2}_{\rm circ}$ = 133, 47 RVs and 6 fitted parameters (including three systematic velocities, $\gamma_{\rm SOPHIE}$, $\gamma_{\rm FIES}$ and $\gamma_{\rm Keck}$ )."
 We report the independent discovery and follow-up observations of the transiung planet HAT-P-14b found by T10., We report the independent discovery and follow-up observations of the transiting planet HAT-P-14b found by T10.
" We have analysed the follow-up data used to secure the detection, and derived an independent set of planetary parameters which are consistent with those found by T10."," We have analysed the follow-up data used to secure the detection, and derived an independent set of planetary parameters which are consistent with those found by T10."
" We provide additional evidence, through AO observations and additional bisector span measurements, that the planetary signal is not caused by astronomical false positives."," We provide additional evidence, through AO observations and additional bisector span measurements, that the planetary signal is not caused by astronomical false positives."
" We find the properties of the parent star to be in good agreement with those found by T10, which is reassuring given the different techniques used in their derivation."," We find the properties of the parent star to be in good agreement with those found by T10, which is reassuring given the different techniques used in their derivation."
" The star is a hot F5 dwarf with == 6583 + 100 K and appears to be relatively young at 0.10.5 Gyr, with a slightly super-solar metallicity |M/H]| = 0.08."," The star is a hot F5 dwarf with = 6583 $\pm$ 100 K and appears to be relatively young at $^{+0.4}_{-0.3}$ Gyr, with a slightly super-solar metallicity [M/H] = 0.08."
" Using the empirical calibrations of 2?) we determine the stellar mass to be 1.30 + 0.03., which agrees well with the values we derive [rom theoretical stellar models."," Using the empirical calibrations of \citet{Torres10star} we determine the stellar mass to be 1.30 $\pm$ 0.03, which agrees well with the values we derive from theoretical stellar models."
" By combining data [rom both groups, we are able to refine the orbital ephemeris /? = 4.627652 + 0.000004 d, Ti, = 2455134.43823 -Γ 0.00022 d, and planetary parameters."," By combining data from both groups, we are able to refine the orbital ephemeris $P$ = 4.627652 $\pm$ 0.000004 d, $T_{\rm 0}$ = 2455134.43823 $\pm$ 0.00022 d, and planetary parameters."
 HAT-P-14b has == 1.20 £ 0.05Rj. == 2.20 + 0.04 aand == 1.28 + 0.12p).," HAT-P-14b has = 1.20 $\pm$ 0.05, = 2.20 $\pm$ 0.04 and = 1.28 $\pm$ 0.12."
". It appears to blur the bi-modal distribution noted by ? in which massive (M, 2 Mj) planets appear either inflated (2,~ 1.3 Rj) or compact (A, 1] Rj)).", It appears to blur the bi-modal distribution noted by \citet{Bakos10b} in which massive $\sim$ 2 ) planets appear either inflated $\sim$ 1.3 ) or compact $\sim$ 1 ).
" We find the orbit to be significandy eccentric, e = 0.095 + 0.011. which hints at the possibility of another body in the system."," We find the orbit to be significantly eccentric, $e$ = 0.095 $\pm$ 0.011, which hints at the possibility of another body in the system."
" As noted by T10, the system promises excellent opportunities for further follow-up and characterisation, in particular due to the brightness of the host star, Έλις = 9.98."," As noted by T10, the system promises excellent opportunities for further follow-up and characterisation, in particular due to the brightness of the host star, $V_{\rm mag}$ = 9.98."
 The star is relatively [ast rotating for a planet host. ==84+1.0kmsL. making ita good candidate for the detection of the Rossiter-McLaughlin effect 2?)..," The star is relatively fast rotating for a planet host, = 8.4 $\pm$ 1.0, making it a good candidate for the detection of the Rossiter-McLaughlin effect \citep[RM,][]{Rossiter24,McLaughlin24}."
" However, due to the very high impact parameter, b = 0.905 + 0.005. the RM amplitude is only expected to be ~ 20 m/s. Despite this, it is süll an attractive target as ? and ? have shown that misaligned planets appear to preferentially orbit stars with high aand ssuch as HAT-P-14."," However, due to the very high impact parameter, b = 0.905 $\pm$ 0.005, the RM amplitude is only expected to be $\sim$ 20 m/s. Despite this, it is still an attractive target as \citet{Winn10} and \citet{Schlaufman10} have shown that misaligned planets appear to preferentially orbit stars with high and such as HAT-P-14."
discussed previously. they generally appear for frequencies for which the turning point of the £2 waves corresponds to the edge of the core.,"discussed previously, they generally appear for frequencies for which the turning point of the $\ell=2$ waves corresponds to the edge of the core."
 The relation between the frequencies of the waves and their turning points is given by (see Soriano et al. 2007) , The relation between the frequencies of the waves and their turning points is given by (see Soriano et al. \cite{soriano07}) )
where v is the frequency of the considered mode. r the radius. and ο) the sound speeed.," where $\nu$ is the frequency of the considered mode, $r$ the radius, and $c(r)$ the sound speeed."
 So that. for the =2 waves. the eritical frequency for which the turning point is at the core edge is approximately Although ος) may slightly vary from model to model. it is not surprising to find that fixed values of v. corresponding to models with similar values of Αι.," So that, for the $\ell=2$ waves, the critical frequency for which the turning point is at the core edge is approximately Although $c(R_{cc})$ may slightly vary from model to model, it is not surprising to find that fixed values of $\nu_c$ corresponding to models with similar values of $R_{cc}$."
 As a consequence. this behaviour appears earlier for more massive stars and for models with similar sizes. but different luminosities. effective temperatures. and gravities.," As a consequence, this behaviour appears earlier for more massive stars and for models with similar sizes, but different luminosities, effective temperatures, and gravities."
 Tables | and 4 present the parameters of the transition models computed with the solar composition (here we use the Grevesse Noels 1993 composition. compatible with helioseismology).," Tables \ref{tab1} and \ref{tab4} present the parameters of the transition models computed with the solar composition (here we use the Grevesse Noels \cite{grevesse93} composition, compatible with helioseismology)."
 As will be discussed below. Table 4 includes overshooting. which ts not the case for Table 1..," As will be discussed below, Table \ref{tab4} includes overshooting, which is not the case for Table \ref{tab1}."
 Tables 2 and 5 correspond to models computed with a larger abundance of metals (|Fe/H] = 0.30) and a helium value following the metallicity-helium relation obtained for the chemical evolution of galaxies (Izotov Thuan 2004)), Tables \ref{tab2} and \ref{tab5} correspond to models computed with a larger abundance of metals ([Fe/H] = 0.30) and a helium value following the metallicity-helium relation obtained for the chemical evolution of galaxies (Izotov Thuan \cite{izotov04}) ).
 Finally Tables 3. and 6 correspond to models with the same overmetallicity but with a solar helium abundanceY., Finally Tables \ref{tab3} and \ref{tab6} correspond to models with the same overmetallicity but with a solar helium abundance.
.. The strongest effect of overmetallicity is to move the evolutionary tracks to lower effective temperatures., The strongest effect of overmetallicity is to move the evolutionary tracks to lower effective temperatures.
 This is true for both helium values., This is true for both helium values.
 In the case of the large helium mass fraction. the ages of the overmetallic transition models are similar to the case of solar values; however. this corresponds to different stages of evolution.," In the case of the large helium mass fraction, the ages of the overmetallic transition models are similar to the case of solar values; however, this corresponds to different stages of evolution."
 For example. the models of 1.15 presented in Tables | and 2 have similar ages. but the one with overmetallicity and large helium abundance (Table 2)) is still on the main sequence with a convective core. whereas the one with solar abundances (Table 1)) is already on the subgiant branch. with a helium core but no convective core anymore.," For example, the models of 1.15 presented in Tables \ref{tab1} and \ref{tab2} have similar ages, but the one with overmetallicity and large helium abundance (Table \ref{tab2}) ) is still on the main sequence with a convective core, whereas the one with solar abundances (Table \ref{tab1}) ) is already on the subgiant branch, with a helium core but no convective core anymore."
 As can be checked in the tables. their luminosities. effective temperatures. and log g are consequently quite ditferent.," As can be checked in the tables, their luminosities, effective temperatures, and log $g$ are consequently quite different."
" Tig ~ AL. A > of5 11.7.13. 713 M4, gan. µια. Thi R,; των. ALjn,  ΑΕ. he dividing line between planets and brown dwarfs. in which case these mass estimates indicate that IIR. 5799 josts a fascinating planetary svsteun."," $_{\rm eff}$ $\sim$ $_*$ $_\odot$ $\lambda$ $>$ $30 - 160$ $5 - 11$ $7 - 13$ $7 - 13$ $_{Jup}$ $\mu$ $\mu$ $_{\rm pl}$ $_{pl}$ $_{Jup}$ $_{Jup}$ $\sim$ $_{Jup}$ the dividing line between planets and brown dwarfs, in which case these mass estimates indicate that HR 8799 hosts a fascinating planetary system."
 However. Mova et al. (," However, Moya et al. ("
2010a) arene that the stellar age is still uuconstraiued. pointing out that some of he age estimators are compromised by the A Dootis character of the star or are statistical du nature.,"2010a) argue that the stellar age is still unconstrained, pointing out that some of the age estimators are compromised by the $\lambda$ Bootis character of the star or are statistical in nature."
 They used astroseisnologv to show that anu older age of ~ Garis possible. if the inclination of the equatorial aue of the star is Z.>36°.," They used astroseismology to show that an older age of $\sim$ 1 Gyr is possible, if the inclination of the equatorial plane of the star is $I_* > 36 \degr$."
 Because the masses of the companions are derived fron) the conmgpurisou of the observed bhDmuinositv to theoretical evolutionary racks of giaut planets and brown dwarfs. the age of he star ds critical in deteriuinius these masses.," Because the masses of the companions are derived from the comparison of the observed luminosity to theoretical evolutionary tracks of giant planets and brown dwarfs, the age of the star is critical in determining these masses."
 At this older age. the companion huninositics would have faded sieuificautlv aud their assigued masses would be well iuto he brown dwarf rauge.," At this older age, the companion luminosities would have faded significantly and their assigned masses would be well into the brown dwarf range."
 Thus. the nature of the IIR. 5799 system would change significautly.," Thus, the nature of the HR 8799 system would change significantly."
 This would affect the onuation scenarios ποος: discussion. and the ability to use these three planets (of simular age aud imoetallicitv) to calibrate evolutionary and atmospheric models of eiaut auets.," This would affect the formation scenarios under discussion, and the ability to use these three planets (of similar age and metallicity) to calibrate evolutionary and atmospheric models of giant planets."
 Iu this letter. we assess the possibility of nore massive conrpaiions by considering the limiting case in which the nasses of the planets are set to the largest values that ανασα. stability cau allow (Section 2)).," In this letter, we assess the possibility of more massive companions by considering the limiting case in which the masses of the planets are set to the largest values that dynamical stability can allow (Section \ref{planet_conf}) )."
 Caven the observed planetary huuimositioes and using evolutionary nodels. we estimate the age corresponding to these Πιο masses.," Given the observed planetary luminosities and using evolutionary models, we estimate the age corresponding to these limiting masses."
 We then evaluate the stability of the ylanets and the dust-producing plauetesinals during hat timescale to assess whether an older age could still cad to a planetary system (planets | planetcsimals) consistent with the observations (Section 33)., We then evaluate the stability of the planets and the dust-producing planetesimals during that timescale to assess whether an older age could still lead to a planetary system (planets + planetesimals) consistent with the observations (Section \ref{stability}) ).
 We discuss he results in Section. L.., We discuss the results in Section \ref{discussion}.
 We now consider whether the configuration of the colpanious can help resolve the uncertainty over the age of IIR. 8799., We now consider whether the configuration of the companions can help resolve the uncertainty over the age of HR 8799.
 The companions have been detected in HST/NICMOS archival data taken LO vears prior to the discovery tage (Lafreniere et al., The companions have been detected in HST/NICMOS archival data taken 10 years prior to the discovery image (Lafreniere et al.
 2009). however. their long orbital periods make the determination of the orbital parametersstilluncertain.," 2009), however, their long orbital periods make the determination of the orbital parametersstilluncertain."
While there are severalconfigurationsthat wouldfit(1) theastrometry.(2) the observed luminosities. aud (3) the requirement,"While there are severalconfigurationsthat wouldfit(1) theastrometry,(2) the observed luminosities, and (3) the requirement"
Wis: similarly the absolutely normalization of the reflection component (Jinorm) anticorrelates with P.,$W_{\rm K\alpha}$; similarly the absolutely normalization of the reflection component $Rnorm$ ) anticorrelates with $F_{\rm K\alpha}$.
" We note that contrary to findings from simple power law fits. Pj, is seen to decrease with flux. when non-unity abundances are assumed."," We note that contrary to findings from simple power law fits, $F_{\rm K\alpha}$ is seen to decrease with flux, when non-unity abundances are assumed."
 We caution however about the large degeneracies associated with complex fits when only the 20 keV data are considered. (, We caution however about the large degeneracies associated with complex fits when only the $-$ 20 keV data are considered. (
The case for requiring supersolar abundances is discussed in Lee et al.,The case for requiring supersolar abundances is discussed in Lee et al.
 1999)., 1999).
 With the exception of the behaviour of P. all other parameters as e.g. LP display similar trends using complex and simple power law fits.," With the exception of the behaviour of $F_{\rm K\alpha}$ , all other parameters as e.g. $\Gamma$ display similar trends using complex and simple power law fits."
" The apparent discrepancy between £j, from simple and complex fits in conjunction with the large errors associated with Pj, would suggest that an ambiguity exists in measurements of the iron line itself. which is unable to resolve."," The apparent discrepancy between $F_{\rm K\alpha}$ from simple and complex fits in conjunction with the large errors associated with $F_{\rm K\alpha}$ would suggest that an ambiguity exists in measurements of the iron line itself, which is unable to resolve."
 A detailed investigation of the time intervals surrounding the bright and flare reveal further complexities., A detailed investigation of the time intervals surrounding the bright and flare reveal further complexities.
" While changes to LF,iy are consistent with flux-correlated studies. we find that there is evidence to suggest that a change in 1; occurs during the time intervals immediately associated with the flare event: in other words. tentative evidence for changes to Zi, are apparent on short time scales."," While changes to $\Gamma_{3-10}$ are consistent with flux-correlated studies, we find that there is evidence to suggest that a change in $F_{\rm K\alpha}$ occurs during the time intervals immediately associated with the flare event; in other words, tentative evidence for changes to $F_{\rm K\alpha}$ are apparent on short time scales."
" In particular. a large increase in the line flux (e.g A2) is evident in the interval immediately following the flare: Vs, increases similarly."," In particular, a large increase in the line flux (e.g ) is evident in the interval immediately following the flare; $W_{\rm K\alpha}$ increases similarly."
" It is curious that ἐν, shows a significant increase after the flare rather than during. and may be an indication that are we witnessing some type of response to the flare (e.g. Fig. 10)."," It is curious that $F_{\rm K\alpha}$ shows a significant increase after the flare rather than during, and may be an indication that are we witnessing some type of response to the flare (e.g. Fig. \ref{fig11-ascaflareplt}) ),"
 but also caution that the evidence is very tentative., but also caution that the evidence is very tentative.
" We note also that j, is comparably high (~ factorof 1.7 increase) during the times surrounding the flare. and times following the flare. in contrast to values presented in Tables 2.. 3.. and 8∎∎"," We note also that $F_{\rm K\alpha}$ is comparably high $\sim$ factor of 1.7 increase) during the times surrounding the flare, and times following the flare, in contrast to values presented in Tables \ref{tab1-pl}, \ref{tab2-pexrav}, and \ref{tab6-intgammachange}."
 Motivated by the temporal tindings of e.g. Miyamoto et al. (, Motivated by the temporal findings of e.g. Miyamoto et al. (
1988) and Cui et al. (,1988) and Cui et al. (
1997) for Cygnus X-I. we next wish to investigate whether the collecting area of coupled with this long observation is sufficient to discern time lags. time leads. and in particular whether reverberation effects can finally be seen.,"1997) for Cygnus X-1, we next wish to investigate whether the collecting area of coupled with this long observation is sufficient to discern time lags, time leads, and in particular whether reverberation effects can finally be seen."
 We acknowledge that a good assessment can be hampered by unevenly sampled data. the nature of which has been investigated by a number of workers (e.g. Edelson Krolik 998: Yaqoob et al.," We acknowledge that a good assessment can be hampered by unevenly sampled data, the nature of which has been investigated by a number of workers (e.g. Edelson Krolik 1998; Yaqoob et al."
 1997) for AG o Int Zand Fenimore (1996) for application to gamma-ray bursts., 1997) for AGN to In`t Zand Fenimore (1996) for application to gamma-ray bursts.
 We adopt a method similar to that presented by Edelson Krolik (1998)., We adopt a method similar to that presented by Edelson Krolik (1998).
 Accordingly. we detine the following formalism. for our calculations of the autocorrelation function (ACF). and cross correlation function (CCF).," Accordingly, we define the following formalism for our calculations of the autocorrelation function (ACF), and cross correlation function (CCF)."
 The ime interval 7 is incremented by n multiples of 64s (7=n. 64): the subsequent definition of ης is the numberof bins for which the difference in time between consecutive time intervals satisfy the present value for 7., The time interval $\tau$ is incremented by n multiples of 64s $\tau=n\times64$ ); the subsequent definition of $n_\tau$ is the number of bins for which the difference in time between consecutive time intervals satisfy the present value for $\tau$.
 The variables {ν and A/ correspond respectively to the ACF and CCF value at 7=0: these terms are used in order to normalize the ACF and CCF so that coherent noise addition at 7=0 is eliminated., The variables $K$ and $M$ correspond respectively to the ACF and CCF value at $\tau=0$; these terms are used in order to normalize the ACF and CCF so that coherent noise addition at $\tau=0$ is eliminated.
 In order to better understand the nature of our findings shown in Figs. [5--17..," In order to better understand the nature of our findings shown in Figs. \ref{fig16-CCFe1e2e3}- \ref{fig18-e1e2e3e4},"
 we run our CCF algorithm on simulated light curves., we run our CCF algorithm on simulated light curves.
 The light curves over which the CCFs are evaluated are identical except that one has a specitied fraction shifted in phase or time., The light curves over which the CCFs are evaluated are identical except that one has a specified fraction shifted in phase or time.
 In principle. time- and phase-shifted light curves are identical at certain Fourier frequencies as for example. for the case in which the power spectrum can be represented by a single sine curve.," In principle, time- and phase-shifted light curves are identical at certain Fourier frequencies as for example, for the case in which the power spectrum can be represented by a single sine curve."
However. if the light curve is the addition of e.g. several sine curves (more representative of actual physical systems). then a time shift,"However, if the light curve is the addition of e.g. several sine curves (more representative of actual physical systems), then a time shift"
single pointing.,single pointing.
" The width of the slit was 1.6"" which, in combination with the 600RI grism and the 2x binned readout, correspondsA». to a resolution of ~ 485 at Ha (6563 "," The width of the slit was $\arcsec$ which, in combination with the 600RI grism and the $2\times 2$ binned readout, corresponds to a resolution of $\sim$ 485 at $\alpha$ (6563 )."
"This resolution prevents us from resolving the narrow component of the Ha line, but it increases the signal-to-noise sufficiently to detect and resolve the broad component."," This resolution prevents us from resolving the narrow component of the $\alpha$ line, but it increases the signal-to-noise sufficiently to detect and resolve the broad component."
" The data were corrected for bias, flatfielded, and the skylines and cosmic rays were removed."," The data were corrected for bias, flatfielded, and the skylines and cosmic rays were removed."
" The wavelength calibration was done by fitting 8 fourth order polynomial to the spectral lines of He, HgCd, Ar and Ne lamps, obtained during daytime."," The wavelength calibration was done by fitting a fourth order polynomial to the spectral lines of He, HgCd, Ar and Ne lamps, obtained during daytime."
" In addition, we checked our calibration against the position of 3 skylines1996)."," In addition, we checked our calibration against the position of 3 skylines."
". The absolute wavelength calibration has a systematic uncertainty of 0.5 pixel ( 37 kms~')) and the errors on the datapoints were determined by tracing the error propagation through the reduction steps, starting from the raw data."," The absolute wavelength calibration has a systematic uncertainty of 0.5 pixel ( 37 ) and the errors on the datapoints were determined by tracing the error propagation through the reduction steps, starting from the raw data."
" For each observation block, obvious outliers, probably caused by erroneous skyline subtraction, were removed."," For each observation block, obvious outliers, probably caused by erroneous skyline subtraction, were removed."
" The resulting spectra were fitted with two superimposed Gaussians and an offset, convolved with the slit width."," The resulting spectra were fitted with two superimposed Gaussians and an offset, convolved with the slit width."
 The best fitting parameters for both spectra are listed in Table 2.., The best fitting parameters for both spectra are listed in Table \ref{fit}.
" In order to determine the significance of the broad component in the NE spectrum, we compared the obtained x? value (464) with a fit in which we only fitted a single Gauss to the spectrum(525)."," In order to determine the significance of the broad component in the NE spectrum, we compared the obtained $\chi^2$ value (464) with a fit in which we only fitted a single Gauss to the spectrum(525)."
 The difference indicates a significance of 7.80., The difference indicates a significance of $\sigma$.
 The x2. values are 2.03 and 2.56 for the SW and NE fits respectively., The $\chi^2_{\rm red}$ values are 2.03 and 2.56 for the SW and NE fits respectively.
" These high values are mainly caused by substructure of the narrow line: a fit to the broad component with a single Gauss, excluding the central region (between -100 and 800 convolved with the resolution, gives of 0.51 for kms-!)),the SW and 0.76 for the NE."," These high values are mainly caused by substructure of the narrow line: a fit to the broad component with a single Gauss, excluding the central region (between -100 and 800 ), convolved with the resolution, gives $\chi^2_{\rm red}$ of 0.51 for the SW and 0.76 for the NE."
" This substructurex2,, might either be caused by spatial surface brightness fluctuations of the remnant within the slit, or by errors in the skyline subtraction."," This substructure might either be caused by spatial surface brightness fluctuations of the remnant within the slit, or by errors in the skyline subtraction."
 We determine the lc errors on the parameters for the broad components by using the errors we determined in our data reduction., We determine the $1\sigma$ errors on the parameters for the broad components by using the errors we determined in our data reduction.
" For the narrow component we first increased the errors such that x2,,=1, and then determined the 1 c errors, using Ax?— 1."," For the narrow component we first increased the errors such that $\chi^2_{\rm red}=1$, and then determined the 1 $\sigma$ errors, using $\Delta \chi^2=1$ ."
" 4mm To determine the shock velocity, we use the X-ray line width of ogas=4900+420kms-!,, as observed by the Reflection Grating Spectrometer (RGS) onboard XMM-Newton2008)."," 4mm To determine the shock velocity, we use the X-ray line width of $\sigma_{\rm RGS} = 4900 \pm 420$, as observed by the Reflection Grating Spectrometer (RGS) onboard XMM-Newton."
. This line width is caused by both the thermal and bulk broadening of the plasma., This line width is caused by both the thermal and bulk broadening of the plasma.
" We follow the method of to disentangle the bulk broadening from the thermal broadening to obtain the forward shock velocity, as described below."," We follow the method of to disentangle the bulk broadening from the thermal broadening to obtain the forward shock velocity, as described below."
" As shown in(2010),, σπας and v, are related as follows: where ra,=Urs/Us: the ratio ot the reverse shock velocity to the forward shock velocity and rpuk=Ubulk,ej/Vbulk, csm: the gradient in the plasma bulk velocity from reverse to forward shock."," As shown in, $\sigma_{\rm RGS}$ and $v_{\rm s}$ are related as follows: where $r_{\rm sh} = v_{\rm rs}/v_{\rm s}$: the ratio ot the reverse shock velocity to the forward shock velocity and $r_{\rm bulk}=v_{\rm bulk,\ ej}/v_{\rm bulk,\ CSM}$ : the gradient in the plasma bulk velocity from reverse to forward shock."
" We obtain ray~1 (whole remnant) and rs,~0.5 (NE), from analytical models with M.j=1.4Mo, E—1.4x10°! erg and an age of 400 years"," We obtain $r_{\rm sh} \simeq 1$ (whole remnant) and $r_{\rm sh} \simeq 0.5$ (NE), from analytical models with $M_{\rm ej} = 1.4\ \rm{M_\odot}$, $E = 1.4\times10^{51}$ erg and an age of 400 years."
" Additionally, we (Badenesconstrain the models with a forward 2008)..shock radius of 15.9+0.8” and 16.340.3” for the entire remnant and the NE respectively and a reverse shock radius of 11.4+2.1” and 13.5+0.2” for the entire remnant and the NE respectively, based on a deprojection of Chandra images (48.9 ks, obs ID:776), following the procedure of(2010)."," Additionally, we constrain the models with a forward shock radius of $15.9 \pm 0.8 \arcsec$ and $16.3 \pm 0.3 \arcsec$ for the entire remnant and the NE respectively and a reverse shock radius of $11.4 \pm 2.1\arcsec$ and $13.5 \pm 0.2\arcsec$ for the entire remnant and the NE respectively, based on a deprojection of Chandra images (48.9 ks, obs ID:776), following the procedure of."
". We obtain Τρι=0.95 with numerical simulations 2006),, using the above parameters."," We obtain $r_{\rm bulk} = 0.95$ with numerical simulations , using the above parameters."
" We use equation (1)) to estimate v,=6000+300 ((whole remnant) and vs=6600+400 ((NE).", We use equation\ref{vRGS}) ) to estimate $v_{\rm s} = 6000\pm300$ (whole remnant) and $v_{\rm s} = 6600\pm 400$ (NE).
" In the remainder of this paper, we use"," In the remainder of this paper, we use"
"of CH;CN. H""CO- (C97. C99) and NH; (Zhang et al.","of $_3$ CN, $^{13}$ $^+$ (C97, C99) and $_3$ (Zhang et al."
 1998)., 1998).
 These molecular line observations all identify a velocity gradient along the proposed disk. and the aand OH maser velocities appear generally. consistent. with these velocity gradients.," These molecular line observations all identify a velocity gradient along the proposed disk, and the and OH maser velocities appear generally consistent with these velocity gradients."
 The northern OH masers. all right-hand circularly polarised. and have velocities ranging from -14 km s! to -12 km s'. while the southern right-hand polarised components have velocities of 2.1 km s'!.," The northern OH masers, all right-hand circularly polarised, and have velocities ranging from -14 km $^{-1}$ to -12 km $^{-1}$, while the southern right-hand polarised components have velocities of 2.1 km $^{-1}$."
 The remaining two maser spots. both associated with the southern clump. have velocities of ~—4 km s!. although it should be noted that one of these spots is clearly à component of a Zeeman pair with one of the right-hand spots.," The remaining two maser spots, both associated with the southern clump, have velocities of $\sim-4$ km $^{-1}$, although it should be noted that one of these spots is clearly a component of a Zeeman pair with one of the right-hand spots."
 The suggestion that the OH masers arise from material in a circumstellar disk around the source is also supported by comparing the OH masers with a recent high resolution infrared image of this source., The suggestion that the OH masers arise from material in a circumstellar disk around the source is also supported by comparing the OH masers with a recent high resolution infrared image of this source.
 Figure 5. shows the integrated OH emission overlaid on a 2.244m K band image of the central region of the source (Sridharan et al 2004)., Figure \ref{fig:ir} shows the integrated OH emission overlaid on a $\mu$ m K band image of the central region of the source (Sridharan et al 2004).
 The image shows a small bipolar K nebula with the emission lobes separated by a dark extinction lane., The image shows a small bipolar K nebula with the emission lobes separated by a dark extinction lane.
 This structure is identical to the kinds of bipolar infrared nebulae see1 around young low mass stars where the K emission ts due tc» scattering off the walls of the cavity cleared by the outflow from the central star and the extinction is tracing the locatio1 of a cireumstellar disk around the source., This structure is identical to the kinds of bipolar infrared nebulae seen around young low mass stars where the K emission is due to scattering off the walls of the cavity cleared by the outflow from the central star and the extinction is tracing the location of a circumstellar disk around the source.
 Clearly the OH masers are associated with the material responsible for the extinction. the material presumably in a circumstellar disk.," Clearly the OH masers are associated with the material responsible for the extinction, the material presumably in a circumstellar disk."
 A position velocity diagram along the disk axis is shown in Figure 6.., A position velocity diagram along the disk axis is shown in Figure \ref{fig:kepler}.
 As indicated on the diagram. most of the OH and mmaser velocities are all consistent with a Keplerian velocity gradient about a central source of mass 5 or «10M... but some of them suggest a mass of ~20 M..," As indicated on the diagram, most of the OH and maser velocities are all consistent with a Keplerian velocity gradient about a central source of mass $\sim$ 5 or $\stackrel{<}{\sim}$ 10, but some of them suggest a mass of $\sim$ 20."
. The first value is similar to the 7 derived by Cesaroni et al. (, The first value is similar to the $\sim$ 7 derived by Cesaroni et al. (
Cesaroni priv.,Cesaroni priv.
 comm.).," comm.),"
 while the last value is similar to the 24 dderived by Cesaroni et al. (, while the last value is similar to the 24 derived by Cesaroni et al. (
1999) and the 20 dderived by Zhang et al. (,1999) and the 20 derived by Zhang et al. (
1998).,1998).
 It therefore appears that the observations suggest that these masers are arising on or. in the surface of a circumstellar disk around this massive young star.," It therefore appears that the observations suggest that these masers are arising on or, in the surface of a circumstellar disk around this massive young star."
 The presence of a OH Zeeman pair provides a measurement of the magnetic field in this disk., The presence of a OH Zeeman pair provides a measurement of the magnetic field in this disk.
 The 6.3km s7! splitting of this pair implies à magnetic field strength of 10.7 mG (Elitzur 1996). pointing away from the observer.," The 6.3km $^{-1}$ splitting of this pair implies a magnetic field strength of 10.7 mG (Elitzur 1996), pointing away from the observer."
 This is relatively high for typical magnetic field strengths towards OH masers in star forming regions (e.g. Garcia-Barreto et al., This is relatively high for typical magnetic field strengths towards OH masers in star forming regions (e.g. Garcia-Barreto et al.
 1988: Hutawarakorn. Cohen 2003 and references therein).," 1988; Hutawarakorn, Cohen 2003 and references therein)."
 The spatial association of the OH and 6.7-GHz methanol masers confirms the close association of the OH and class II methanol masers proposed by Caswell (1996) and modelled by Cragg et al. (, The spatial association of the OH and 6.7-GHz methanol masers confirms the close association of the OH and class II methanol masers proposed by Caswell (1996) and modelled by Cragg et al. (
2002).,2002).
 The flux ratio of these masers. $(6668)/S(1665)= 20.3/2.7 = 7.5. places IRAS 2012644104 in OH-favoured sources.," The flux ratio of these masers, S(6668)/S(1665)= 20.3/2.7 = 7.5, places IRAS $20126+4104$ in OH-favoured sources."
 Note that although the OH and mmasers are closely associated. there 1s à clear difference in position. amounting to 0.1 aresee (170 AU). suggesting that the OH and mmasers are not co-propagating.," Note that although the OH and masers are closely associated, there is a clear difference in position, amounting to 0.1 arcsec (170 AU), suggesting that the OH and masers are not co-propagating."
 Nevertheless the models of Cragg et al. (, Nevertheless the models of Cragg et al. (
2002) can be used to place some constraints on the physical conditions in the regions where the masers are observed.,2002) can be used to place some constraints on the physical conditions in the regions where the masers are observed.
" The presence of 1665 MHz maser emissionfrom OH. but the absence of 1667 MHz maser emission suggests that the material has a relatively high gas temperatures. Z4,> 30K and relatively high H» gas densities. 1> IO0em-.. unless the dust temperature is ~ 300K."," The presence of 1665 MHz maser emissionfrom OH, but the absence of 1667 MHz maser emission suggests that the material has a relatively high gas temperatures, $T_{\rm gas} > 30$ K and relatively high $_2$ gas densities, $n>10^6$ , unless the dust temperature is $\stackrel{>}{\sim}300$ K,"
star.,star.
 We adopt the binary ephemeris frou Wolf ct al. (, We adopt the binary ephemeris from Wolf et al. (
"1993) Π.Ο)=2115615.1156(1)|O.1582061601)E.where TZ, is the inferior conjunction of the white dwarf.","1993) $ T_{o}(HJD) = 2445615.4156(4) + 0.15820616(4)~E$,where $T_{o}$ is the inferior conjunction of the white dwarf."
 The trailed specra over the full wavelength rauge are shown in Fig., The trailed spectra over the full wavelength range are shown in Fig.
 2 with the ain to display the motion of the weak hues (the intensity scale is adjusted so that line appears saturaed)., 2 with the aim to display the motion of the weak lines (the intensity scale is adjusted so that line appears saturated).
 The dise and red star enission conmponuents are seen in the lines of σδ. 1172. ΤΙΝ. the Bowen blend and the ΙΤ].," The disc and red star emission components are seen in the lines of 4388, 4472, 4481, the Bowen blend and the 4713."
 The red star componenut is tje sharp “S-wave moving from red to blue at phase 0.5., The red star component is the sharp `S'-wave moving from red to blue at phase 0.5.
 It can also be traced in the £512 aud the £118 lines., It can also be traced in the 4542 and the 4418 lines.
 Note that the 7Swave conrponeut disappears earlier (binary pliase 0.7) than that of the neighbouring £172 (binary phase 0.75)., Note that the `S'-wave component disappears earlier (binary phase 0.7) than that of the neighbouring 4472 (binary phase 0.75).
 We reconstruct the Doppler images of the oeissiou lines using the trailed spectra (Marsh and Horne 1955)., We reconstruct the Doppler images of the emission lines using the trailed spectra (Marsh and Horne 1988).
 A Doppler image is the reconstruction of the emission line distribution in velocity space and has been particularly successful in resolving the location of enission componcuts such as the red star (IP Pee: Tarlaftis et al., A Doppler image is the reconstruction of the emission line distribution in velocity space and has been particularly successful in resolving the location of emission components such as the red star (IP Peg; Harlaftis et al.
 199D). the eas streams (OY Car in outburst: IEbulaftis aud ανα 1996). the bright spot (CS2000|25: IIarlaftis et al.," 1994), the gas stream (OY Car in outburst; Harlaftis and Marsh 1996), the bright spot (GS2000+25; Harlaftis et al."
 1996) aud spiral waves in the outer accretion dise (Steeghls. ITarlaftis aud TWorne 1997: Ihulaftis ct al.," 1996) and spiral waves in the outer accretion disc (Steeghs, Harlaftis and Horne 1997; Harlaftis et al."
 1999)., 1999).
 We built the Doppler tages of the emission lines (see above references for the procedure) and. after subtracting the axisvunuetric dise enüssion. we can zoom outo the Roche lobe of the red star (Fig.," We built the Doppler images of the emission lines (see above references for the procedure) and, after subtracting the axisymmetric disc emission, we can zoom onto the Roche lobe of the red star (Fig."
 3 from left to right. high-ionization to low-ionization Lunes. 1656.P. £388. 1172. 1713. List).," 3 from left to right, high-ionization to low-ionization lines, 4686, 4388, 4472, 4713, 4481)."
 The line. emissiou. ou the Doppler nuages of Fig.," The line emission, on the Doppler images of Fig."
 3. is stronger on the side of the red star facing the white dwart and further it weakens towards the equator of the red star. indicative of screcuime by the disc.," 3, is stronger on the side of the red star facing the white dwarf and further it weakens towards the equator of the red star, indicative of screening by the disc."
 The maps show, The maps show
Fast electrons. typically generated by high-energy photons or cosmic ray collisions. are crucial for a wide range of astrophysical problems.,"Fast electrons, typically generated by high-energy photons or cosmic ray collisions, are crucial for a wide range of astrophysical problems."
" For example. cosmic ray ionization is important for the thermal balance of interstellar clouds (2).. electrons produced by the Comptonization of -ray photons affect the early history of supernova remnants (2).. and broad-line emission regions around quasars are likely ionized by hard photons from those sources. making the fate of the liberated electrons of paramount importance (2),"," For example, cosmic ray ionization is important for the thermal balance of interstellar clouds \citep{spitzer69}, electrons produced by the Comptonization of $\gamma$ -ray photons affect the early history of supernova remnants \citep{xu91b}, and broad-line emission regions around quasars are likely ionized by hard photons from those sources, making the fate of the liberated electrons of paramount importance \citep{shull85}."
 As such. there have been numerous studies of the interactions between these electrons and a background gas. through the processes of collisional excitation. ionization. and. electron-electron scattering.," As such, there have been numerous studies of the interactions between these electrons and a background gas, through the processes of collisional excitation, ionization, and electron-electron scattering."
 These calculations have included both analytic approaches (22222) and numerical explorations (2222)...," These calculations have included both analytic approaches \citep{spitzer68, spitzer69, jura71, bergeron73, xu91} and numerical explorations \citep{habing71, shull79, shull85, valdes08}."
 Most have focused on interactions with atomic or ionic gases. relevant especially to low-density material in the interstellar tor intergalactic) media.," Most have focused on interactions with atomic or ionic gases, relevant especially to low-density material in the interstellar (or intergalactic) media."
 Others have considered the additional effects of molecules (e.g.. 25).," Others have considered the additional effects of molecules (e.g., \citealt{dalgarno99}) )."
 Recently. the fate of X-rays in the high-redshift intergalactic medium (GM) has become an important question.," Recently, the fate of X-rays in the high-redshift intergalactic medium (IGM) has become an important question."
 Before the reionization of HI. ultraviolet ionizing photons are trapped near their sources. but X-rays can travel much larger distances through the IGM.," Before the reionization of HI, ultraviolet ionizing photons are trapped near their sources, but X-rays can travel much larger distances through the IGM."
 As such. they are thought to provide the most important radiation background for the bulk of the IGM. slowly ionizing the mostly neutral gas far from stars.," As such, they are thought to provide the most important radiation background for the bulk of the IGM, slowly ionizing the mostly neutral gas far from stars."
 However. most of their energy is deposited indirectly. through the fast electrons generated by photo-ionization.," However, most of their energy is deposited indirectly, through the fast electrons generated by photo-ionization."
 Por example. the electrons produced by X-rays from the first hard sources (quasars. supernova remnants. or even Population III stars: ??)) are most likely responsible for heating the IGM to T~1000K before reionization (22)..," For example, the electrons produced by X-rays from the first hard sources (quasars, supernova remnants, or even Population III stars; \citealt{oh01, venkatesan01}) ) are most likely responsible for heating the IGM to $T \sim 1000 \kel$ before reionization \citep{kuhlen05-sim, furl06-glob}."
 This heating and ionization could influence future generations of structure (e.g. 223) and it has important observational consequences.," This heating and ionization could influence future generations of structure (e.g., \citealt{ricotti02, oh03-entropy}) ) and it has important observational consequences."
 For example. it can dramatically change the redshifted 21 em signal from the early IGM (222)... and it could even provide a signature of exotic physics during the Dark Ages before the first sources appear (2222)..," For example, it can dramatically change the highly-redshifted 21 cm signal from the early IGM \citep{furl06-review, kuhlen06-21cm, pritchard07}, and it could even provide a signature of exotic physics during the Dark Ages before the first sources appear \citep{chen04-decay, furl06-dm, mapelli06, valdes08}."
 Moreover. HI photons produced as the secondary electrons collisionally excite atoms in the background gas can greatly affect the 21 em signal by changing the hypertine level populations of HI ¢???).. and the secondary ionizations may play an important role in the HI (and Hel) reionization of the IGM (?2??)..," Moreover, HI photons produced as the secondary electrons collisionally excite atoms in the background gas can greatly affect the 21 cm signal by changing the hyperfine level populations of HI \citep{chuzhoy06-first, pritchard07, chen08}, and the secondary ionizations may play an important role in the HI (and HeI) reionization of the IGM \citep{ricotti04_a, ricotti05, volonteri09}."
 Here we revisit this problem using a Monte Carlo model (described in 32)) that incorporates the mostrecent cross-sections for interactions between X-rays and primordial gas., Here we revisit this problem using a Monte Carlo model (described in \ref{mc}) ) that incorporates the mostrecent cross-sections for interactions between X-rays and primordial gas.
 We review the relevant physics of collisional ionization. excitation. and electron scattering in §$3--S and comment on some of our approximations in $6..," We review the relevant physics of collisional ionization, excitation, and electron scattering in \ref{ion}- \ref{heat} and comment on some of our approximations in \ref{ignored}. ."
 We present our principal results. together with a comparison to previous work. in $7. and a simple example of their utility for," We present our principal results, together with a comparison to previous work, in \ref{results} and a simple example of their utility for"
sullers a close eucounter with a planet.,suffers a close encounter with a planet.
 Iu this situation. since the gravitational effect of the planet ou tlie particle exceeds that of the star. the roles of planet aid star are reversed and the Ixepleriau portion of the Hamiltonian is taken to occur around the planet aud the star is relegated to the role of disturbanceLevison&Duncan(199D).," In this situation, since the gravitational effect of the planet on the particle exceeds that of the star, the roles of planet and star are reversed and the Keplerian portion of the Hamiltonian is taken to occur around the planet and the star is relegated to the role of disturbance\citet{ld94}."
. Although a smaller timestep is used duriug a close eucounter. the duration of au encounter is typically small enough that the simulation proceeds rapidly.," Although a smaller timestep is used during a close encounter, the duration of an encounter is typically small enough that the simulation proceeds rapidly."
 In the intermediate region where planetary. aud stellar forces are comparable. the timestep is also reduced but the integration remains heliocentric.," In the intermediate region where planetary and stellar forces are comparable, the timestep is also reduced but the integration remains heliocentric."
 The magnitudee of the radiatiou force on a particle is eeiven by where L is stellar lunuiinosity. 4 is the particle cross-sectional area. (pj; is the radiation pressure coellicient. ¢ is the speed of light and r is the heliocentric cistauce.," The magnitude of the radiation force on a particle is given by where $L$ is stellar luminosity, $A$ is the particle cross-sectional area, $Q_{PR}$ is the radiation pressure coefficient, $c$ is the speed of light and $r$ is the heliocentric distance."
 It is staudard practice to describe the streugth of the radiation force on a given particle as a fraction of the gravitational force ou a particle. such that The constant 9 is depeuds ou the particle size aud Composition. as well as the stellar ass ancl luminosity. but is inclependent of the particles heliocentric distance.," It is standard practice to describe the strength of the radiation force on a given particle as a fraction of the gravitational force on a particle, such that The constant $\beta$ is depends on the particle size and composition, as well as the stellar mass and luminosity, but is independent of the particle's heliocentric distance."
 The magnitude of ddrag is proportional to 5»., The magnitude of drag is proportional to $\beta$.
 The ratio of solar wind drag to ddrag is depeudeut on particle size. but is relatively coustant over the rauge of particles usually considered (s>0.5jun.Burusetal.1979).. so the combined drag effects of PR and solar wind cau be expressed using a sinele parameter given by where ste is the ratio of solar wind drag to ddrag (Burusetal.1979).," The ratio of solar wind drag to drag is dependent on particle size, but is relatively constant over the range of particles usually considered \citep[$s > 0.5 \mu$m, ][]{burns79}, so the combined drag effects of PR and solar wind can be expressed using a single parameter given by where $sw$ is the ratio of solar wind drag to drag \citep{burns79}."
. Following Moro-Martin&Malhotra(2002).. we use a constant. value of sw=0.35 in all our sinulatious.," Following \citet{mm02}, we use a constant value of $sw=0.35$ in all our simulations."
 The acceleration on a particle due to solar wind drag aud radiation. iguoriug terms of order (ο)/i> aand higher.. is. given. by The first term in tliis expression is the result of the outwardly-directed radiation pressure. which causes the vet acceleration in the radial direction to be reduced to (1 τρως.," The acceleration on a particle due to solar wind drag and radiation, ignoring terms of order $(v/c)^{2}$ and higher, is given by The first term in this expression is the result of the outwardly-directed radiation pressure, which causes the net acceleration in the radial direction to be reduced to $(1-\beta) F_{grav}$ ."
 The remaining two, The remaining two
different components of the universe: matter. dark energy. curvature and radiation.,"different components of the universe: matter, dark energy, curvature and radiation."
 Phe gravitational growth index + will be a parameter of Κον interest., The gravitational growth index $\gamma$ will be a parameter of key interest.
 It can distinguish other heories from Einstein. gravity (see for example Lincer(2005):Lincleretal.(2007):Guzzo (2008))).," It can distinguish other theories from Einstein gravity – (see for example \citet{growth,cahn-gamma,guzzo}) )."
 The merit ofa large scale structure survey in terms of its gravitational owobative power may be conveniently quantified by the uncertainty (5)., The merit of a large scale structure survey in terms of its gravitational probative power may be conveniently quantified by the uncertainty $\sigma(\gamma)$.
 The Figure of Merit. Science. Working Group (Albrechtetal.2009). found: that for the suite of uture Stage LLL experiments. expected. to. be completed oore the proposed. Joint Dark Lnerey Mission. (JDEAL)program. the anticipated uncertainty is (75)=0.21.," The Figure of Merit Science Working Group \citep{StageIII} found that for the suite of future Stage III experiments, expected to be completed before the proposed Joint Dark Energy Mission (JDEM)program, the anticipated uncertainty is $\sigma(\gamma) = 0.21$."
 The standard technique for making such parameter estimation predictions is the Fisher matrix Clegmark 1907)., The standard technique for making such parameter estimation predictions is the Fisher matrix \citep{tth}.
". For a survey covering a volume Vy where the mean ealaxy number density is n. the element of the Fisher matrix [or parameters p; and p; is obtained as an integral over the space of modes k Clegmark1997). by: pu "" ""n y:5 The accessible modes are weightedyr), )due Üp;to shot Op;noise lín according to an cllective volume (Feldmanetal.1994): Voth) = . Mott yo2("," For a survey covering a volume $V_0$ where the mean galaxy number density is $\bar{n}$ the element of the Fisher matrix for parameters $p_i$ and $p_j$ is obtained as an integral over the space of modes $\mathbf{k}$ \citep{teg}, by: = d^3 k ( )^2 The accessible modes are weighted due to shot noise $1/\bar n$ according to an effective volume \citep{fkp}: ) = V_0 ( )^2."
"12) asThe constraint. everage comes mostly from regionspe), where nP(k.ji)1. hat is V,2Vy."," The constraint leverage comes mostly from regions where $\bar{n} P(k,\mu) 
\gtrsim 1$, that is $V_{e} \approx V_0$."
 In order to avoid the uncertainties associated with reatment of non-linearities. we truncate the Fisher matrix integral at a maxiniuun value A.," In order to avoid the uncertainties associated with treatment of non-linearities, we truncate the Fisher matrix integral at a maximum value $k_{+}$."
 We take A=OLA /Alpe. which is the scale where departures from linear theory begin o become significant (see. e.g. the analysis of Percivaletal. (2009))).," We take $k_{+} = 0.1\,h$ /Mpc, which is the scale where departures from linear theory begin to become significant (see, e.g., the analysis of \citet{dep}) )."
 Sce Sec., See Sec.
 4.2. [or a further investigation of non- cllects., \ref{snonlin} for a further investigation of non-linear effects.
 The information about 5. comes from. two dilferent ruts of the power spectrum., The information about $\gamma$ comes from two different parts of the power spectrum.
 The real space. isotropic part. corresponding to no redshift space distortions. or js=0 (observations transverse to the line of sight) in the linear regime. is proportional to the growth factor squared: PA) 2 PU DialE(1," The real space, isotropic part, corresponding to no redshift space distortions, or $\mu=0$ (observations transverse to the line of sight) in the linear regime, is proportional to the growth factor squared: (k) = b^2 P(k) D(a)^2."
3) Note that surveys lacking sullicient. redshift resolution are only sensitive to the transverse modes due to smearing along the line of sight. (see. e.g... Padmanabhan: (2008))).," Note that surveys lacking sufficient redshift resolution are only sensitive to the transverse modes due to smearing along the line of sight (see, e.g., \citet{nikhil}) )."
 Using Eq. (8)).," Using Eq. \ref{eq:gro}) ),"
 the information carried by this part involves ss 5oes The redshift space distortions thatin the power spectrum eive further information through the parameter. f. which with Eqs. (5))," the information carried by this part involves = 2 _0^a The redshift space distortions in the power spectrum give further information through the parameter $f$, which with Eqs. \ref{eq:fdef}) )"
 and (8)) reads [955, and \ref{eq:gro}) ) reads f =.
-405) Therefore. if we define the anisotropic part alone as p) 2bfp os gu HG] it carries information on  through= ," Therefore, if we define the anisotropic part alone as ) 2 b f ^2 + f^2 ^4, it carries information on $\gamma$ through = _m(a) +."
This factor eives a sense of the information281 fromOia theᾗ redshift distortions., This factor gives a sense of the information from the redshift distortions.
 Because the measurements become noisier when subclividecl into angular bins. and because a substantial majority of the information resides in the spherically averaged. power spectrum (Shojietal.2009)... analyses frequently use the one dimensional. spherically averaged power spectrum bf | ," Because the measurements become noisier when subdivided into angular bins, and because a substantial majority of the information resides in the spherically averaged power spectrum \citep{martincom}, analyses frequently use the one dimensional, spherically averaged power spectrum b f + f^2."
This incorporates information from boththe original isotropic power spectrum and the redshift distortion anisotropies. and may be most familiar in terms of the DyxDAHo)ox(lk)U7 factor of Eisensteinal.(2005).," This incorporates information from boththe original isotropic power spectrum and the redshift distortion anisotropies, and may be most familiar in terms of the $D_V\propto [D_A^2/H(z)]^{1/3}\propto 
(k_\perp^2 k_\parallel)^{-1/3}$ factor of \citet{eis05}."
". In. particular. the sensitivity to , arises from — |."," In particular, the sensitivity to $\gamma$ arises from = +."
 In Sec. ?7?7..," In Sec. \ref{ssolve},"
 we will investigate the relative importance of the transverse. anisotropic. spherically averaged. as well as full versions of the power spectrum for constraints on the eravitational growth index and other parameters.," we will investigate the relative importance of the transverse, anisotropic, spherically averaged, as well as full versions of the power spectrum for constraints on the gravitational growth index and other parameters."
 The constraints on ~ expected from nearer termi (Stage LIL) surveys are not that informative. as mentioned. with c()=0.21 compared to a clillerence As=0.13 (Lueetal.2004:Linder2005:et2007). between general relativity and DCGDP sravitv (Dvalietal.2000:Dellavetetal.2002). for cxample.," The constraints on $\gamma$ expected from nearer term (Stage III) surveys are not that informative, as mentioned, with $\sigma(\gamma)=0.21$ compared to a difference $\Delta\gamma=0.13$ \citep{luess,growth,cahn-gamma} between general relativity and DGP gravity \citep{dgp,ddg} for example."
 We therefore turn to Stage IV experiments ancl assess their potential for a more accurate test of the standard cosmological mocel., We therefore turn to Stage IV experiments and assess their potential for a more accurate test of the standard cosmological model.
 We consider two versions of Stage IV. power spectrum experiments: BigBOSS (Schlegeletal.2009). is à proposed erouncd-based wide field spectroscopic survey and JDIZM-PS (Gehrelsetal.2009). is a proposed. space-based. wide field evism survey., We consider two versions of Stage IV power spectrum experiments: BigBOSS \citep{bigboss} is a proposed ground-based wide field spectroscopic survey and JDEM-PS \citep{jdem} is a proposed space-based wide field grism survey.
 Both aim at measuring the three dimensional spatial distribution of galaxies to study barvon.: acoustic oscillations and the growth of structure., Both aim at measuring the three dimensional spatial distribution of galaxies to study baryon acoustic oscillations and the growth of structure.
 Both experiments would use the Stage HE experiment BOSS (Schlegeletal. 2009).. detecting Luminous red galaxies (LRG) out to z= 0.7. as à springboard to higher redshifts.," Both experiments would use the Stage III experiment BOSS \citep{boss}, detecting luminous red galaxies (LRG) out to $z=0.7$ , as a springboard to higher redshifts."
 BigBOSS would extend mapping of LRGout to z= Land to the southern sky and both experiments would supplement LI: with dillerent classes of emission line galaxies (EL) out to z 2., BigBOSS would extend mapping of LRGout to $z=1$ and to the southern sky and both experiments would supplement LRG with different classes of emission line galaxies (EL) out to $z\approx2$ .
 Following Schlegeletal.(2009):Gehrelsct (2000).we give in Table 1. the redshift range. survey solid angle Quy.expected target galaxy bias [actors byne; and bey. mean galaxy number density n. anc wavelength resolution &=AfAA of the spectrographs to be used. (so the redshift resolution o;—93z/(1|]z)—Ιἲ ly," Following \citet{bigboss,jdem,anze}, ,we give in Table \ref{spec} the redshift range, survey solid angle $\Omega_{sky}$ ,expected target galaxy bias factors $b_{LRG}$ and $b_{EL}$ , mean galaxy number density $\bar{n}$ , and wavelength resolution $R=\lambda/\Delta\lambda$ of the spectrographs to be used (so the redshift resolution $\sigma_z=\delta z/(1+z)=R^{-1}$ )."
 We consider variations in redshift. number density. and bias in Sec. 7?..," We consider variations in redshift, number density, and bias in Sec. \ref{ssurvey}. ."
variations due to pulsations (IxXilkennvοἱal.1988:Lawson&Cottrell1997).,"variations due to pulsations \citep{kil88, law97}."
. None of the Πας stars has shown anv evidence for dust formation., None of the HdC stars has shown any evidence for dust formation.
 As described in (1993).. RCD and dC stars are spectroscopically similar. having Che same abundances. effective temperatures. and absolute luminosities.," As described in \citet{bru98}, RCB and HdC stars are spectroscopically similar, having the same abundances, effective temperatures, and absolute luminosities."
 However. the HdC stars lack the large visible brightness variations. [Ro excesses and emission lines seen in (he RCB stars. which are associated with mass-loss and dust formation.," However, the HdC stars lack the large visible brightness variations, IR excesses and emission lines seen in the RCB stars, which are associated with mass-loss and dust formation."
 Only five Πας stars are known 1967).. but as thev are bright stars (all are in the IID catalogue). it is likely that there are manv more in the Galaxy. (Warner1967:Drilling1986).," Only five HdC stars are known \citep{war67}, but as they are bright stars (all are in the HD catalogue), it is likely that there are many more in the Galaxy \citep{war67,dri86}."
. The latter is also true of the RCB stars. of which only about 50 are known in the Galaxy Tisserandοἱal. 2008).," The latter is also true of the RCB stars, of which only about 50 are known in the Galaxy \citep{cla96,zan05,tis08}."
. The origin of RCD stars has long been debated., The origin of RCB stars has long been debated.
 The leading candidates appear to be (1) the merger of a C/O white dwarl ancl a Ile white dwarL in which the merged object undergoes an episode of helium-burning (hat creates an extended atmosphere around the degenerate core (the so-called. double-degenerate (DD) scenario Webbink 1984).. and (2) a late or very late IHe-shell flash that blows up an isolated incipient white dwarf into a “born-again” cool supergiant (Ibenetal.1996).," The leading candidates appear to be (1) the merger of a C/O white dwarf and a He white dwarf, in which the merged object undergoes an episode of helium-burning that creates an extended atmosphere around the degenerate core (the so-called double-degenerate (DD) scenario \citep{ibe84, web84}, and (2) a late or very late He-shell flash that blows up an isolated incipient white dwarf into a “born-again"" cool supergiant \citep{ibe96}."
 The similar chemical abundances of HdC and RCB stars have long suggested that the two classes of stars are evolutionarily linked., The similar chemical abundances of HdC and RCB stars have long suggested that the two classes of stars are evolutionarily linked.
" Stronger evidence for a link was established recently when Clavtonetal.(2005.2007). discovered (that both HdC stars and RCD stars have huge overabundances of O. with ""O/!5O close to and in some cases less than unity (the solar and interstellar ratios are 2:500)."," Stronger evidence for a link was established recently when \citet{cla05, cla07} discovered that both HdC stars and RCB stars have huge overabundances of $^{18}$ O, with $^{16}$ $^{18}$ O close to and in some cases less than unity (the solar and interstellar ratios are $\sim$ 500)."
 Garceia-Ilernándezetal.(2009). have recently confirmed the low values of !O/!O in several of the stars observed by (2007)., \citet{gar09} have recently confirmed the low values of $^{16}$ $^{18}$ O in several of the stars observed by \citet{cla07}.
. Such values are not approached in anv other class of stars., Such values are not approached in any other class of stars.
 For several reasons. including the observed low !9O0/PO ratios. the late He-shell flash scenario seems unlikely to be the origin of RCB stars. and thus also unlikely to be the origin of the IldC stars," For several reasons, including the observed low $^{16}$ $^{18}$ O ratios, the late He-shell flash scenario seems unlikely to be the origin of RCB stars, and thus also unlikely to be the origin of the HdC stars"
relatively strong coupliug leads to a large azimuthal twisting of the maeguetic field.,relatively strong coupling leads to a large azimuthal twisting of the magnetic field.
 This. together with the requirement that the magnetic field remain closed. leads to very stroug spiu-dowan torques ou the star.," This, together with the requirement that the magnetic field remain closed, leads to very strong spin-down torques on the star."
 For this reason. the CC93 model predicts slower equilibrimn spins than other models for a eiven field streneth 2007).. ," For this reason, the CC93 model predicts slower equilibrium spins than other models for a given field strength \citep[e.g., see][]{johnskrull3ea99, johnskrull07}."
However. as shown iu section 3.. when the opening of the field is properly taken into account. the spin-down torques are weaker. and the magnitude of this effect is larger for cases with stroug maguetic coupling (as in the CC'93 model).," However, as shown in section \ref{sec_results}, when the opening of the field is properly taken into account, the spin-down torques are weaker, and the magnitude of this effect is larger for cases with strong magnetic coupling (as in the CC93 model)."
function of time as the warp diffuses.,function of time as the warp diffuses.
" Indeed, this is the reason for the large error bars in Fig."," Indeed, this is the reason for the large error bars in Fig."
" 8 when attempting to fit a single az value to the /, profile, particularly at low a where in the ? theory the dependency of a2 on wv is much stronger."," \ref{fig:largeA} when attempting to fit a single $\alpha_{2}$ value to the $l_{x}$ profile, particularly at low $\alpha$ where in the \citetalias{ogilvie99} theory the dependency of $\alpha_{2}$ on $\psi$ is much stronger."
" What is possible is to check whether the deviations expected from such non-linear theory follow the observed trend, i.e. a general decrease value of the diffusion coefficient."," What is possible is to check whether the deviations expected from such non-linear theory follow the observed trend, i.e. a general decrease value of the diffusion coefficient."
 The short-dashed line in Fig., The short-dashed line in Fig.
" 8 shows the relation between o? and obased on the theory of ? for a fixed value of i»=0.55, computed numerically based on a routine kindly provided to us by Gordon Ogilvie."," \ref{fig:largeA} shows the relation between $\alpha_2$ and $\alpha$based on the theory of \citetalias{ogilvie99} for a fixed value of $\psi=0.55$, computed numerically based on a routine kindly provided to us by Gordon Ogilvie."
 We can indeed see that the theory does reproduce qualitatively the observed trend., We can indeed see that the non-linear theory does reproduce qualitatively the observed trend.
" As an additional test of the theory, we have also evolved the standard equation for diffusive evolution (Eq. 5)),"," As an additional test of the theory, we have also evolved the standard equation for diffusive evolution (Eq. \ref{eq:pringle}) ),"
" where the diffusion coefficients o1 (now different from o) and o» are computed at each radius and at each timestep, based on Ogilvie's non-linear theory."," where the diffusion coefficients $\alpha_{1}$ (now different from $\alpha$ ) and $\alpha_2$ are computed at each radius and at each timestep, based on Ogilvie's non-linear theory."
" In this case, we do not have any free parameter to fit: α is the nominal value of the viscosity coefficient, and αι and 02 are prescribed functions of a and 7 which we compute at each radius and time by evaluation of the relevant integrals in ? using the routine provided."," In this case, we do not have any free parameter to fit: $\alpha$ is the nominal value of the viscosity coefficient, and $\alpha_{1}$ and $\alpha_2$ are prescribed functions of $\alpha$ and $\psi$ which we compute at each radius and time by evaluation of the relevant integrals in \citetalias{ogilvie99} using the routine provided."
 The resulting profiles of {ας are shown in Fig., The resulting profiles of $l_x$ are shown in Fig.
 for the case where o—0.43., \ref{fig:profilelargeA} for the case where $\alpha=0.43$.
" The black solid lines show the results of the SPH simulations at£—0 andt—1000 (in code units), while the red dashed line show the corresponding profiles computed from Eq. (5))."," The black solid lines show the results of the SPH simulations at $t=0$ and $t=1000$ (in code units), while the red dashed line show the corresponding profiles computed from Eq. \ref{eq:pringle}) )."
 The agreement of the non-linear theory with the SPH results is excellent., The agreement of the non-linear theory with the SPH results is excellent.
 Note that the warp diffusion coefficient in the theory of ? depends also on the amount of bulk viscosity present in the disc., Note that the warp diffusion coefficient in the theory of \citetalias{ogilvie99} depends also on the amount of bulk viscosity present in the disc.
" Indeed, in order to get the good match shown in Fig."," Indeed, in order to get the good match shown in Fig."
" we have computed the warp diffusion coefficient assuming a bulk viscosity with magnitude a,=5a/3, as expected from the SPH formalism (see sect. 3.2.2))."," we have computed the warp diffusion coefficient assuming a bulk viscosity with magnitude $\alpha_{\rm b}=5\alpha/3$, as expected from the SPH formalism (see sect. \ref{sec:av}) )."
" At low disc viscosities (small o), the simulations show a steepening effect in the warp profile — the physical result of different parts of the warp propagating at different speeds."," At low disc viscosities (small $\alpha$ ), the simulations show a steepening effect in the warp profile — the physical result of different parts of the warp propagating at different speeds."
 We find that the steepening becomes stronger as the disc viscosity is reduced (evident from the increase in the error bars seen in Fig., We find that the steepening becomes stronger as the disc viscosity is reduced (evident from the increase in the error bars seen in Fig.
 as a— 0) and in the limit of extremely small viscosity can result in a near-complete break in the disc., \ref{fig:largeA} as $\alpha \to 0$ ) and in the limit of extremely small viscosity can result in a near-complete break in the disc.
 This is demonstrated in Fig., This is demonstrated in Fig.
 10 which shows the equivalent of Fig., \ref{fig:profilelargeAsmallalpha} which shows the equivalent of Fig.
 9 for a low viscosity calculation (o= 0.03) shown at intervals of t=500 in code units and evolved as far as £=3500., \ref{fig:profilelargeA} for a low viscosity calculation $\alpha=0.03$ ) shown at intervals of $t=500$ in code units and evolved as far as $t=3500$.
 A three dimensional rendering of the disc structure is shown in Fig., A three dimensional rendering of the disc structure is shown in Fig.
 11 at t=1500 in code units., \ref{fig:discbreak} at $t=1500$ in code units.
" Despite the apparent break the calculations are able to resolve a thin but steady stream of material continuing to accrete across the discontinuity, indicating that the disc regions remain connected even in this extreme case."," Despite the apparent break the calculations are able to resolve a thin but steady stream of material continuing to accrete across the discontinuity, indicating that the disc regions remain connected even in this extreme case."
 The simulation results in this case are more difficult to compare with theory because of the development of large unphysical oscillations around the warp in the one dimensional code., The simulation results in this case are more difficult to compare with theory because of the development of large unphysical oscillations around the warp in the one dimensional code.
" This is not unexpected from the non-linear theory, since in this regime the evolution of the disc is not well described by a diffusion equation."," This is not unexpected from the non-linear theory, since in this regime the evolution of the disc is not well described by a diffusion equation."
" In particular, ? (Eq."," In particular, \citetalias{ogilvie99} (Eq."
" 141) shows that and thus for large w (more specifically, when 7)> 24a, which is the case here) the effective viscosity becomes negative, suggesting that diffusion is no longer an appropriate description."," 141) shows that and thus for large $\psi$ (more specifically, when $\psi>\sqrt{24}\alpha$ , which is the case here) the effective viscosity becomes negative, suggesting that diffusion is no longer an appropriate description."
 Whilst for small amplitude warps one expects a transition to the wave-like regime at low a (based on the dispersion relation derived, Whilst for small amplitude warps one expects a transition to the wave-like regime at low $\alpha$ (based on the dispersion relation derived
scattering from the CBR. which at this redshift) produces energy loss euivalent to that due to a magnetic field. of 40 n'T. Using a model with ellective pitch-angle scattering of electrons (Jalle Perola 1973). the plot of age agains oojected. distance along the source was well fitted by a straight line with gradient ~1.2O0! km !|. inferre ages being of the order of 10 vears (as found by Jággers De Grijp 1985).,"scattering from the CBR, which at this redshift produces energy loss equivalent to that due to a magnetic field of 0.40 nT. Using a model with effective pitch-angle scattering of electrons (Jaffe Perola 1973), the plot of age against projected distance along the source was well fitted by a straight line with gradient $\sim 1.2 \times 10^4$ km $^{-1}$, inferred ages being of the order of $10^7$ years (as found by Jäggers De Grijp 1985)."
 The intercept was non-zero. reflecting the oesence of steeper-spectrum material surrounding the je ermination: derived. velocities were similar if the intercep was mace zero by choosing a steeper initial power-law index (2.53).," The intercept was non-zero, reflecting the presence of steeper-spectrum material surrounding the jet termination; derived velocities were similar if the intercept was made zero by choosing a steeper initial power-law index (2.53)."
 These inferred. ages ancl velocities in the tails are comparable to those found by spectral age methods in other WAT sources F'Iavlor 11990: ODonoghue 11993) and would imply outflow which is considerably faster than the sound. speed. in the external medium. given the temperature of the gas around. 1130. (discussed. below): this is perhaps surprising in view of the absence of any evidence for post-hot spot shock structures in the tails an of their. generally relaxed: appearance.," These inferred ages and velocities in the tails are comparable to those found by spectral age methods in other WAT sources Taylor 1990; O'Donoghue 1993) and would imply outflow which is considerably faster than the sound speed in the external medium, given the temperature of the gas around 130 (discussed below); this is perhaps surprising in view of the absence of any evidence for post-hot spot shock structures in the tails and of their generally relaxed appearance."
 The usual caveats apply to velocities determined by spectral-ageing methods. but it should be noted that most factors that can alloc the velocity (including a contribution to the energy. density from relativistic protons. a particle filling factor less than unity. ancl significant projection of the radio source) wil produce velocities higher than the value given above.," The usual caveats apply to velocities determined by spectral-ageing methods, but it should be noted that most factors that can affect the velocity (including a contribution to the energy density from relativistic protons, a particle filling factor less than unity, and significant projection of the radio source) will produce velocities higher than the value given above."
 Only if the assumptions involved in the spectral ageing analysis are seriously wrong — for example. if there is significant particle acceleration in the tails or significant magnetic lield inhomogeneity can the tail velocity be much lower than this value.," Only if the assumptions involved in the spectral ageing analysis are seriously wrong – for example, if there is significant particle acceleration in the tails or significant magnetic field inhomogeneity – can the tail velocity be much lower than this value."
 Eviclenee for such processes is discussed in Eilek (1996) and references therein., Evidence for such processes is discussed in Eilek (1996) and references therein.
 Ailey ((1983) report on IPC observations of130., Miley (1983) report on IPC observations of.
. Serendipitouslv. the source is also included: in the [field of a 39.4 ksROSAT PSPC pointed. observation. taken rom the public archives. of the X-ray emitting supernova remnant OOL59L|5147 (Plelfermann. Aschenbach νοο 1991: Reich 1992).," Serendipitously, the source is also included in the field of a 39.4 ks PSPC pointed observation, taken from the public archives, of the X-ray emitting supernova remnant 04591+5147 (Pfeffermann, Aschenbach Predehl 1991; Reich 1992)."
 Although the X-ray source associated with is 32 arcmin away from the pointing centre of the PSPC. and is thus badly. vignetted. the observations have superior signal-to-noise to the data and show details of the X-ray structure of the source.," Although the X-ray source associated with is 32 arcmin away from the pointing centre of the PSPC, and is thus badly vignetted, the observations have superior signal-to-noise to the data and show details of the X-ray structure of the source."
 The cluster is detected. at 2200X100 PSPC counts between 0.1:2.4 keV. (derived from a circle of 11 aremin radius about the centre of the X-ray source. using a background annulus between 11 and 17.5 arcmin). in spite of the reduced sensitivity of the PSPC at this oll-axis distance.," The cluster is detected at $2200 \pm 100$ PSPC counts between 0.1–2.4 keV (derived from a circle of 11 arcmin radius about the centre of the X-ray source, using a background annulus between 11 and 17.5 arcmin), in spite of the reduced sensitivity of the PSPC at this off-axis distance."
 Because of the dillicultv of measuring the |ackeround in the presence of extended. emission from the SNR. and because the shadows of the ring and one of the racial struts pass close to the source. the derived count rate of ~6«107 + is uncertain.," Because of the difficulty of measuring the background in the presence of extended emission from the SNR, and because the shadows of the ring and one of the radial struts pass close to the source, the derived count rate of $\sim 6 \times 10^{-2}$ $^{-1}$ is uncertain."
 A rough correction for vignetting would imply an on-axis count rate of 9.-102ὃν 4., A rough correction for vignetting would imply an on-axis count rate of $9 \times 10^{-2}$ $^{-1}$.
 Using the Post-Recuetion Ollline Software (PROS) within IRA. I made spectral fits to the data. correcting for the olf-axis location of the source.," Using the Post-Reduction Offline Software (PROS) within IRAF, I made spectral fits to the data, correcting for the off-axis location of the source."
 A single. Iavmond-Smith model provided a good [it (47=15.8 with 26 degrees of reedom) eiving a best-fit temperature AY=2.94 keV: 1e fitted galactic Ny was 0.9025107 ef," A single Raymond-Smith model provided a good fit $\chi^2 =
15.8$ with 26 degrees of freedom), giving a best-fit temperature $kT =
2.9^{+9}_{-2}$ keV; the fitted galactic $N_H$ was $0.9^{+0.5}_{-0.2}
\times 10^{22}$ $^{-2}$ [cf."
 the value [0.2 1077 em. ? predicted by interpolation [rom Stark ((1992)]., the value of $0.4 \times 10^{22}$ cm $^{-2}$ predicted by interpolation from Stark (1992)].
 Errors quoted for Na and AT are le lor wo interesting pareuneters., Errors quoted for $N_H$ and $kT$ are $1\sigma$ for two interesting parameters.
 Abuncanees were poorly 'onstrained: TO per cent solar abundance gave marginally re best Gt., Abundances were poorly constrained; 70 per cent solar abundance gave marginally the best fit.
 With this best-fit model. the 0.1-2.4. keV uminositv of the cluster is 51021 W. consistent with the uminosity derived. on crude spectral assumptions. from the data by Miley ((1983): the cluster is thus comparable in X-ray luminosity o rich Abell clusters. and the temperature consistent with he temperature-Iuminosity relation DDavic 11993).," With this best-fit model, the 0.1-2.4 keV luminosity of the cluster is $5 \times 10^{37}$ W, consistent with the luminosity derived, on crude spectral assumptions, from the data by Miley (1983); the cluster is thus comparable in X-ray luminosity to rich Abell clusters, and the temperature consistent with the temperature-luminosity relation David 1993)."
 There is no evidence for a lower temperature in he central regions of the source. and so no evidence that a cooling Low is present: this appears to be normal for the yost clusters of WATs (Norman. Burns Sulkanen 1988: Commez 11997) although Schindler Prieto (1997) suggest that a weak cooling Low is present in Abell 2634. the host cluster of 4465. and Livelra A inhabits a cooling How with hieh mass deposition rates (David 11990.," There is no evidence for a lower temperature in the central regions of the source, and so no evidence that a cooling flow is present; this appears to be normal for the host clusters of WATs (Norman, Burns Sulkanen 1988; Gómmez 1997) although Schindler Prieto (1997) suggest that a weak cooling flow is present in Abell 2634, the host cluster of 465, and Hydra A inhabits a cooling flow with high mass deposition rates (David 1990)."
 The best-fit Gaussian to the oll-axis point-spread funetion (PSE) of ROSAT at this distance from the pointing centre has σzz10 aresec (Hasinger 11995).," The best-fit Gaussian to the off-axis point-spread function (PSF) of ROSAT at this distance from the pointing centre has $\sigma \approx
70$ arcsec (Hasinger 1995)."
 For a radio-N-rav comparison | have smoothed the broad-band (Q.1-2.4 keV) X-ray image with a Gaussian of this size: this should allow the coma-induced asymmetry of the PSE to be neglected., For a radio-X-ray comparison I have smoothed the broad-band (0.1-2.4 keV) X-ray image with a Gaussian of this size; this should allow the coma-induced asymmetry of the PSF to be neglected.
 Phe X-ray images re[3€ 130-xrav)) show an extended. structure on. scales comparable to the length of the radio source (ic. 1 Alpe)., The X-ray images \\ref{3C130-xray}) ) show an extended structure on scales comparable to the length of the radio source (i.e. $\sim 1$ Mpc).
 The cluster gas seems reasonably svmumetrical. about the raclio source. in contrast to the ebumpy structures. with olfset racio sources. seen in some lower-Iuminosity WAT hosts even at lower spatial resolution BBurns 11994: Gomme 11997).," The cluster gas seems reasonably symmetrical about the radio source, in contrast to the clumpy structures, with offset radio sources, seen in some lower-luminosity WAT hosts even at lower spatial resolution Burns 1994; Gómmez 1997)."
 The distortion of the X-ray. isophotes to. the northeast coincides with. and may be related to. the kink (~ 150300 kpe from the nucleus) in the northern tail: there is no structure in the X-ray. emission which can be related to the sudden change in direction at the end of the southern racio tail. however.," The distortion of the X-ray isophotes to the northeast coincides with, and may be related to, the kink $\sim
150$ –300 kpc from the nucleus) in the northern tail; there is no structure in the X-ray emission which can be related to the sudden change in direction at the end of the southern radio tail, however."
 Given the large and asvmametrical PSE. ] have not attempted to fit radial profiles to the X-ray data.," Given the large and asymmetrical PSF, I have not attempted to fit radial profiles to the X-ray data."
 Approaches to the source dynamics of WATS in the literature BBurns 1981: ος 11984: ODonoghue, Approaches to the source dynamics of WATs in the literature Burns 1981; Eilek 1984; O'Donoghue
"wwe have removed the velocity-dependent observer corrections from the Boltzmann transport equation in the appropriate, Lagrangian approach as described by refsec:noc,, (12)), but retain the Newtonian gravity, O(v/c)-hydrodynamics,equation and reduced opacities of the mmodel.","we have removed the velocity-dependent observer corrections from the Boltzmann transport equation in the appropriate, Lagrangian approach as described by \\ref{sec:noc}, , equation \ref{eq:lagnoc}) ), but retain the Newtonian gravity, $\mathcal{O}(v/c)$ -hydrodynamics, and reduced opacities of the model."
" Dropping the observer corrections in model rresults in a dramatic change in the properties of the core at bounce, as can be seen in Figure 1 and Table 2.."," Dropping the observer corrections in model results in a dramatic change in the properties of the core at bounce, as can be seen in Figure \ref{fig:bounce} and Table \ref{tab:models}."
" The homologous core mass drops from 0.717 ffor the mmodel to 0.427 ffor the mmodel; the latter of which is virtually indistinguishable from the homologous core mass,Λ4ςμ., for the most physically complete model (0.429 ffor GR-FullOp))."," The homologous core mass drops from 0.717 for the model to 0.427 for the model; the latter of which is virtually indistinguishable from the homologous core mass, for the most physically complete model (0.429 for )."
" This coincidental alignment of the bounce shock positions for the most and least physically complete models should be contrasted with the lower electron and lepton fractions (Y;, Yr) and density (see Table 2)) in the homologous core for the mmodel relative to the mmodel, as well as the larger inflow velocities and densities and lower Y, and ΥΠ, outside the shock, which imply, among other things, an increased ram pressure against which the shock must propagate."," This coincidental alignment of the bounce shock positions for the most and least physically complete models should be contrasted with the lower electron and lepton fractions $Y_e, Y_L$ ) and density (see Table \ref{tab:models}) ) in the homologous core for the model relative to the model, as well as the larger inflow velocities and densities and lower $Y_e$ and $Y_L$ outside the shock, which imply, among other things, an increased ram pressure against which the shock must propagate."
" The shock (Figure 2)) in the mmodel starts more vigorously than in the mmodels, but does not show the large overshoot of the other reduced opacity model,"," The shock (Figure \ref{fig:shock}) ) in the model starts more vigorously than in the models, but does not show the large overshoot of the other reduced opacity model,."
" All of the shock trajectories cross near 35 ms after N-ReducOp..bounce, with the mmodel having the deepestshock throughout the rest of the"," All of the shock trajectories cross near 35 ms after bounce, with the model having the deepestshock throughout the rest of the"
90101).,.
 To account for this. we use the analytic function described by IKCUO which was in turn derived from the mocels of Sudarskyetal.(2005).," To account for this, we use the analytic function described by KG10 which was in turn derived from the models of \citet{sud05}."
 This function results ina time-depeudence of the albedo as the strong ivradiation of the atmospheres of eiaut plauets removes reflective condensates frou the upper atmospheres during periastron passage., This function results in a time-dependence of the albedo as the strong irradiation of the atmospheres of giant planets removes reflective condensates from the upper atmospheres during periastron passage.
 This fiction can have a dramatic effect iu daipemime the flux ratio at simall starplanet separations., This function can have a dramatic effect in dampening the flux ratio at small star–planet separations.
 The phase angle à is defined to be zero when the plauct is af superior conjunction aud is described by where wis the aremueut of periastrou aud. f is the true anomaly., The phase angle $\alpha$ is defined to be zero when the planet is at superior conjunction and is described by where $\omega$ is the argument of periastron and $f$ is the true anomaly.
" Thus. in terms of orbital parameters. niuiuiuuu phase occurs when kw|f=90° and maxinuun phase occurs when w|f=270""."," Thus, in terms of orbital parameters, minimum phase occurs when $\omega + f = 90\degr$ and maximum phase occurs when $\omega + f
= 270\degr$."
 The exact nature of a planetary phase function depends upon the assumptions reearding the scattering properties of the atinospherce., The exact nature of a planetary phase function depends upon the assumptions regarding the scattering properties of the atmosphere.
 Rather than assuming isotropic scattering (Lambert sphere). we adopt the clupirically derived phase function of Wilton(1992) which is based upou observations of Jupiter auc Venus and incorporates substantially wore back-scatteriug due to cloud-covering.," Rather than assuming isotropic scattering (Lambert sphere), we adopt the empirically derived phase function of \citet{hil92} which is based upon observations of Jupiter and Venus and incorporates substantially more back-scattering due to cloud-covering."
 This approach contains a correction to the planctary visual iiaguitude of the form leading to a phase function given by which is used throughout the remainder of this paper., This approach contains a correction to the planetary visual magnitude of the form leading to a phase function given by which is used throughout the remainder of this paper.
 As shown by Sudarskyetal.(2005) and I&CUO. the imaxinmnui flux ratio docs not uecessarily occur at zero phase anele for à uou-circular orbit.," As shown by \citet{sud05} and KG10, the maximum flux ratio does not necessarily occur at zero phase angle for a non-circular orbit."
 This is because the starplauet separation r is constantlv changing and is eiven bv where eis the semi-major axis aud e is the orbital eccentricity., This is because the star–planet separation $r$ is constantly changing and is given by where $a$is the semi-major axis and $e$ is the orbital eccentricity.
 The Ir couponent of Equation 1 becomes dominant for highly eccentric orbits which is especially inportaut when the orbital inchnation is considered., The $R_p^2/r^2$ component of Equation \ref{fluxratio} becomes dominant for highly eccentric orbits which is especially important when the orbital inclination is considered.
" To add the effect of inchuation angle to the phase ""nuctiou. the phase angle (Equation 2)) is modified as ollows: At Ist and 3rd quarter (a=907 and a= 2707). the Hux ratio is completely independent of inclination anele."," To add the effect of inclination angle to the phase function, the phase angle (Equation \ref{phaseangle}) ) is modified as follows: At 1st and 3rd quarter $\alpha = 90\degr$ and $\alpha = 270\degr$ ), the flux ratio is completely independent of inclination angle."
" However, at all other pliase angles there arises a complex xittern. of flux ratios from the inclination aud arewet of periastrou."," However, at all other phase angles there arises a complex pattern of flux ratios from the inclination and argument of periastron."
 Tere we describe these dependencies., Here we describe these dependencies.
 The two cases worth considering first are the extremes of edge-on (/= 907) aud face-on (/= 07) orbits.," The two cases worth considering first are the extremes of edge-on $i
= 90\degr$ ) and face-on $i = 0\degr$ ) orbits."
 As noted earlier. the radial velocity method is biased towards the detection of edge-on orbits due to the increase in the seimui-auuplitude of the signal.," As noted earlier, the radial velocity method is biased towards the detection of edge-on orbits due to the increase in the semi-amplitude of the signal."
 Tn addition. most of the planets monitored for phase signatures are kuown to transit.," In addition, most of the planets monitored for phase signatures are known to transit."
 Thus the case of edge-on orbits is currently the dominant form of investigated svstemis., Thus the case of edge-on orbits is currently the dominant form of investigated systems.
" This guarantees the observabilitv of both zero plase aud full phase. the contribution of which to the fiux ratio depends upon the starplauet separation at à=0"","," This guarantees the observability of both zero phase and full phase, the contribution of which to the flux ratio depends upon the star–planet separation at $\alpha = 0\degr$."
 The case of face-on orbits mneaus that the phase function becomes completely flat since onlv half phase will be visible at any oue time., The case of face-on orbits means that the phase function becomes completely flat since only half phase will be visible at any one time.
 Thus the fux ratio is completely determined by the ecceutricity of the orbit which drives both the starplanct separation and the changing albedo of the upper atinospliere., Thus the flux ratio is completely determined by the eccentricity of the orbit which drives both the star–planet separation and the changing albedo of the upper atmosphere.
 Bevoud the cases of edec-on and face-on orbits. the interaction of inclination aud periastron argumneut becomes iore complex.," Beyond the cases of edge-on and face-on orbits, the interaction of inclination and periastron argument becomes more complex."
" Figure 1 demonstrates this for four eccentric orbits. cach with fixed periastrou argunents, which are inclined from edge-on to face-on."," Figure \ref{phases} demonstrates this for four eccentric orbits, each with fixed periastron arguments, which are inclined from edge-on to face-on."
 Within the range of 15Sw£z135° the flux ratio actually increases with decreasing orbital inclination as the davside of the planet becomes more visible where the planet is the hottest. most noticable for the case of av=907.," Within the range of $45\degr \lesssim \omega \lesssim
135\degr$ the flux ratio actually increases with decreasing orbital inclination as the dayside of the planet becomes more visible where the planet is the hottest, most noticable for the case of $\omega =
90\degr$."
 Outside of this regine. the peak flux from the planet generally declines as the inclination iucreases aud thus the access to the full phase diuinishes.," Outside of this regime, the peak flux from the planet generally declines as the inclination increases and thus the access to the full phase diminishes."
 A further level of detail to the description of inclination dependence may be added by calculating the peak flux ratio for the full rauge of iuclinatious (07<7<907) aud periastron arguineuts (07«iw 360°) for a given period and eccentricity., A further level of detail to the description of inclination dependence may be added by calculating the peak flux ratio for the full range of inclinations $0\degr < i < 90\degr$ ) and periastron arguments $0\degr < \omega < 360\degr$ ) for a given period and eccentricity.
 The resulting intensity maps then show the optimal orbital configuration for detection and how the distribution of peak flux ratios smoothly varies witli these coufleurations., The resulting intensity maps then show the optimal orbital configuration for detection and how the distribution of peak flux ratios smoothly varies with these configurations.
 Shown in Figure 2. are two such examples of these iutensitv maps for eccentricities of 0.2 aud 0.5., Shown in Figure \ref{array} are two such examples of these intensity maps for eccentricities of 0.2 and 0.8.
 Iu each case. the strougest flax ratio occurs where the full phase of the planet coincides with the smallest starplanet separation (/=907. w= 2707).," In each case, the strongest flux ratio occurs where the full phase of the planet coincides with the smallest star–planet separation $i = 90\degr$, $\omega =
270\degr$ )."
 Tlowever. notice the interaction which occurs where w=907.," However, notice the interaction which occurs where $\omega
= 90\degr$."
 The chaugiug peak planctary flux along this line of the intensity maps is due to the competing componcuts of the plase function and starplauet separation as the planet moves from a crescent phase near poriastron to the full phase of the plauct at apastrou., The changing peak planetary flux along this line of the intensity maps is due to the competing components of the phase function and star–planet separation as the planet moves from a crescent phase near periastron to the full phase of the planet at apastron.
 For highly ecceutric orbits. the planet flux at i=90° becomes the brightest when the orbit is viewed face-on since the flux ratio is dominated by the starplauet separation.," For highly eccentric orbits, the planet flux at $\omega =
90\degr$ becomes the brightest when the orbit is viewed face-on since the flux ratio is dominated by the star–planet separation."
 The starplanet projected separation is a component which will iufiueuce the target selection for coronagrapl experiueuts for direct detection of the reflected planetary flux., The star–planet projected separation is a component which will influence the target selection for coronagraph experiments for direct detection of the reflected planetary flux.
" Beichmanetal.(2010) tabulate the inner working augle of selected future erouncd-based inaeiug instrunients which rauge from 0.037 to 0.17"". compared to 0.035"" to O.s50 for James Webb Space Telescope GNST) iustruunceuts."," \citet{bei10} tabulate the inner working angle of selected future ground-based imaging instruments which range from $0.03\arcsec$ to $0.17\arcsec$ , compared to $0.035\arcsec$ to $0.850\arcsec$ for James Webb Space Telescope (JWST) instruments."
 The angular projected separation of the planet from the star is eiven by, The angular projected separation of the planet from the star is given by
"Uicrstauding the natire of dark matter remains a ""uudeainenutal problem inm astroplivsics aud cosimnologv.",Understanding the nature of dark matter remains a fundamental problem in astrophysics and cosmology.
 Since the discovery of neutriünos LnOlr-Zeoro rest lass (Fukudactal.1998:DileikvetL998).. the possibility hat neutrinos coutribute to cosmological «ark matter as become a hot topic again.," Since the discovery of neutrinos' non-zero rest mass \citep{Fukuda,Bilenky}, the possibility that neutrinos contribute to cosmological dark matter has become a hot topic again."
 In particular. the sterile jeutrinos belong to a class of candidate dark matter yaricles with no standard model interaction.," In particular, the sterile neutrinos belong to a class of candidate dark matter particles with no standard model interaction."
 Although he recext MiniBooNE data challenges the LSND result hat sugecsts the existence of CV scale sterie neutrimos (Aeuilar-Arevaloctal.2007). iLOLC lInassive sterile 101trinos (og.," Although the recent MiniBooNE data challenges the LSND result that suggests the existence of eV scale sterile neutrinos \citep{Aguilar}, more massive sterile neutrinos (eg."
 keV. MeV) may still exist.," keV, MeV) may still exist."
 Tjo fact that active neutrinos have mass nauplies that right-handed neutrinos should exist which may indeed be massive sterile neutrinos., The fact that active neutrinos have mass implies that right-handed neutrinos should exist which may indeed be massive sterile neutrinos.
 The existence of the sterile neutrinos has been invoked to explain many phenomena such as nüssius nass (DodelsouandWidrow1991:ShiaudFuller1999) and the hieh teuiperatire of the hot eas in Alilkv Wav and clusters (ChanaudChu2007.2005).," The existence of the sterile neutrinos has been invoked to explain many phenomena such as missing mass \citep{Dodelson,Fuller} and the high temperature of the hot gas in Milky Way and clusters \citep{Chan,Chan2}."
" ""Therefore. it js worthwhile to discuss observational consequeuces if massive sterile neuvinos exist. which may decay iuto light neutrinos. positron-clectron pairs and plotous."," Therefore, it is worthwhile to discuss observational consequences if massive sterile neutrinos exist, which may decay into light neutrinos, positron-electron pairs and photons."
" Iu this article. we consider the possibility that sterile neutrino decays eive rise othe 511 keV lines in Milkv. Wav and thus otain bounds ou the mass nmi, and total decay rate D of the sterile neutrimos using relevant obscrvational data."," In this article, we consider the possibility that sterile neutrino decays give rise to the 511 keV lines in Milky Way and thus obtain bounds on the mass $m_s$ and total decay rate $\Gamma$ of the sterile neutrinos using relevant observational data."
 The bright 511 keV annihilation line from λμ Wax iis been observed for a few decades (Leventhalet197s:Iknódlsederetal. 2005).. and its origin has en mmeh debated.," The bright 511 keV annihilation line from Milky Way has been observed for a few decades \citep{Leventhal,Knodlseder}, and its origin has been much debated."
 Recent values. of the 511 keV photon fhx from the bulee aud disk are (1.05£0.06)< phi 7s E Laud (Kidd«10 Spl ὃνἆ respectively (I&nódlsederetal. 90051.," Recent values of the 511 keV photon flux from the bulge and disk are $(1.05 \pm 0.06) \times 
10^{-3}$ ph $^{-2}$ $^{-1}$ and $(0.7 \pm 0.4) \times 
10^{-3}$ ph $^{-2}$ $^{-1}$ respectively \citep{Knodlseder}. ."
.. Assimune a )ositrounuu fraction of fp=0.93. oue can rauslate jese duteusities to amuililajon rates of (1.5+40.1)xes ! aud (0.3+0.2)«1057 st 105)octivelv IznódlseDual.2005).," Assuming a positronium fraction of $f_p=0.93$, one can translate these intensities to annihilation rates of $(1.5 \pm 0.1) \times 10^{43}$ $^{-1}$ and $(0.3 \pm 0.2) \times 
10^{43}$ $^{-1}$ respectively \citep{Knodlseder}."
. The zuimnililation rate iu re bulee is several fiues larecr than hat in the disk., The annihilation rate in the bulge is several times larger than that in the disk.
 The source of positrons in the «isk can he explained by t1ο decay of ΑΙ., The source of positrons in the disk can be explained by the decay of $^{26}$ Al.
 Using a diss model. he . ≻∐∪↑∪∐∏∏⊼↕↴∖↴↸⊳⋜↧↕↸⊳∏↕⋜↧↑↸∖≼ to be 5↽ ∖↓∩↓↻⊔⊳⋯−↴∖↴↓∙⋅≻ which can account for 601005€ of the disk eiissiou (Iuódlsederetal.2005).," Using a disk model, the photon flux is calculated to be $5 \times 
10^{-4}$ ph $^{-2}$ $^{-1}$, which can account for $60-100 \%$ of the disk emission \citep{Knodlseder}."
. However. the origin of the nlee source is still an openquestion in astrophysics.," However, the origin of the bulge source is still an openquestion in astrophysics."
 Recent observation by INTECRAL/IBIS indicates tlat he upper lini of photon fiux for resolved single poiut, Recent observation by INTEGRAL/IBIS indicates that the upper limit of photon flux for resolved single point
"The refraction index of the electromagnetic waves (N,-) in a cold magnetized plasma satisfies the dispersion equation where wand k are the wave frequency and wave vector. c, and wy are the electron plasma and cyclotron frequencies respectively. and @ is the angle between the wave vector and the magnetic field.","The refraction index of the electromagnetic waves $N_{\sigma}$ ) in a cold magnetized plasma satisfies the dispersion equation where $\omega$ and $\mathbf{k}$ are the wave frequency and wave vector, $\omega_{\mathrm{p}}$ and $\omega_{\mathrm{B}}$ are the electron plasma and cyclotron frequencies respectively, and $\theta$ is the angle between the wave vector and the magnetic field."
 In Eq. (AI).," In Eq. \ref{N2}) ),"
 c=-1 for the X-mode (fast extraordinary) and the Z-mode (slow extraordinary): 7=+1 for the O-mode (fast ordinary) and the W-mode (whistlers)., $\sigma=-1$ for the X-mode (fast extraordinary) and the Z-mode (slow extraordinary); $\sigma=+1$ for the O-mode (fast ordinary) and the W-mode (whistlers).
" The polarization state of the magnetoionic modes is described by the following parameters: 7;- is the axial 1-U-V+UVratio cos?of 4.the polarization ellipse. andwhere L,- tsthe longitudinal part of the polarization.", The polarization state of the magnetoionic modes is described by the following parameters: where $T_{\sigma}$ is the axial ratio of the polarization ellipse and $L_{\sigma}$ is the longitudinal part of the polarization.
" The group velocity of the magnetotonic waves equals N The derivative of the refraction index with respect to the propagation direction can be caleulated using the relation The different magnetoionic modes exist in the following frequency ranges: The cutoff and resonance frequencies are given by: In vacuum. the dispersion parameters become: By expanding the polarization parameters (A4--A5)) in the small parameters VU-Land. VV and retaining only the zero order terms in the expansions. we obtain the polarization state of the magnetiomic modes near the cyclotron frequency in a low-density plasma: According to Melrose Dulk (1982)) and Aschwanden (1990) ). the growth rate of the magnetotonic oscillations equals (the wave-mode index c is hereafter omitted for brevity) where v and p are the electron velocity and momentum. the indices z and L indicate the longitudinal and transversal components of the vectors with respect to the magnetic field (in particular. V-=Neos@ and N_= Nsin@). B=efc. T=(1—87)7 is the relativistic factor, J,CU and J/GU are the Bessel function and its derivative over the argument Jt. and Hopf(p) is the electron distribution function satisfying the normalization condition where n is the total electron concentration."," The group velocity of the magnetoionic waves equals where The derivative of the refraction index with respect to the propagation direction can be calculated using the relation The different magnetoionic modes exist in the following frequency ranges: The cutoff and resonance frequencies are given by: In vacuum, the dispersion parameters become: By expanding the polarization parameters \ref{Ts}- \ref{Ls}) ) in the small parameters $\sqrt{U}-1$ and $\sqrt{V}$ and retaining only the zero order terms in the expansions, we obtain the polarization state of the magnetioinic modes near the cyclotron frequency in a low-density plasma: According to Melrose Dulk \cite{mel82}) ) and Aschwanden \cite{asc90}) ), the growth rate of the magnetoionic oscillations equals (the wave-mode index $\sigma$ is hereafter omitted for brevity) where $\mathbf{v}$ and $\mathbf{p}$ are the electron velocity and momentum, the indices $z$ and $\bot$ indicate the longitudinal and transversal components of the vectors with respect to the magnetic field (in particular, $N_z=N\cos\theta$ and $N_{\bot}=N\sin\theta$ ), $\beta=\varv/c$, $\Gamma=(1-\beta^2)^{-1/2}$ is the relativistic factor, $J_s(\lambda)$ and $J'_s(\lambda)$ are the Bessel function and its derivative over the argument $\lambda$, and $f(\mathbf{p})$ is the electron distribution function satisfying the normalization condition where $n$ is the total electron concentration."
 The growth rates of the vacuum modes can be obtained by substituting the corresponding dispersion parameters (AT11--A12)), The growth rates of the vacuum modes can be obtained by substituting the corresponding dispersion parameters \ref{NV}- \ref{TV}) ).
 If we introduce the dimensionless parameters and. in addition. use the polar coordinates (αια) for the distribution function (where « is the angle between the electron velocity and the magnetic field). then the expression for the growth rate takes the form," If we introduce the dimensionless parameters and, in addition, use the polar coordinates $(u, \alpha)$ for the distribution function (where $\alpha$ is the angle between the electron velocity and the magnetic field), then the expression for the growth rate takes the form"
(SOLI/MDI) on the Solar and Heliospheric Observatory (SOMO).,(SOI/MDI) on the Solar and Heliospheric Observatory (SOHO).
 The MDI project is supported by NASA grant. NAG5-8878 to Stanford University., The MDI project is supported by NASA grant NAG5-8878 to Stanford University.
 SOIIO is a project of international cooperation between ESA and NASA., SOHO is a project of international cooperation between ESA and NASA.
 This work was supported in part bv NASA Grant NAG5-LO0912 to SB., This work was supported in part by NASA Grant NAG5-10912 to SB.
(it covered most of the range to which the experiment was sensitive).,(it covered most of the range to which the experiment was sensitive).
 LE NWMLADP. had. confirmed. a flat. patch in. the CODE range ((—4. 20) but not elsewhere. then we would iàve the strange result that data would. the significance of a detection.," If WMAP had confirmed a flat patch in the COBE range $\ell=4-20$ ) but not elsewhere, then we would have the strange result that data would the significance of a detection."
 This obviously means that we rave excessively favoured a Gaussian. prior., This obviously means that we have excessively favoured a Gaussian prior.
 Alore interesting is what actually happened: WALAP ziled to confirm the COBE Uat patch., More interesting is what actually happened: WMAP failed to confirm the COBE flat patch.
 This showed that indeed. the observed CODIS flat. patch was a fluke: it was à svstematic., This showed that indeed the observed COBE flat patch was a fluke: it was a systematic.
 Magueijo&Mecdeiros(2003). explained in detail the origin of the systematic error., \cite{jooes} explained in detail the origin of the systematic error.
 The rejection of the Dat patch described in Section ?? as a selection elleet could be a further example of the dangers of blind use of a Gaussian prior., The rejection of the flat patch described in Section \ref{flat} as a selection effect could be a further example of the dangers of blind use of a Gaussian prior.
 Let us consider a simplified model in which IN is the number of observed. independent bispectrum components and p be the probability that cach one of them is within » sigma of the mean., Let us consider a simplified model in which $N$ is the number of observed independent bispectrum components and $p$ be the probability that each one of them is within $n$ sigma of the mean.
 Lf we define a Lat patch with η=1 we have pz0.68., If we define a flat patch with $n=1$ we have $p\approx 0.68$.
 The probability of a Lat patch longer than L starting at a given Lixed multipole is p., The probability of a flat patch longer than $L$ starting at a given fixed multipole is $p^L$.
 In WAIAD a flat patch occurs with length £=16. starting at f=32.," In WMAP a flat patch occurs with length $L=16$, starting at $\ell=32$."
 Its. probability for a Gaussian is pj=0.002. seeminely ruling out. Ciaussianitv.," Its probability for a Gaussian is $p_L=0.002$, seemingly ruling out Gaussianity."
 In Section ?? we argued that a selection cllect was at play., In Section \ref{flat} we argued that a selection effect was at play.
 Using our simplified model. the probability of a Gaussian generating a flat pateh longer than L starting ds plnL| 1).," Using our simplified model, the probability of a Gaussian generating a flat patch longer than $L$ starting is $p^L(N-L+1)$ ."
 In this case. considering that IN5ς (number of signal dominated bispectrum components). we get a probability 0.18. i.e. consistency with Ciaussianitv.," In this case, considering that $N\approx 100$ (number of signal dominated bispectrum components), we get a probability $0.18$, i.e. consistency with Gaussianity."
 ln Section 77. we refined this argument with Monte Carlo simulations., In Section \ref{flat} we refined this argument with Monte Carlo simulations.
 The problem with this argument is that it is reversed if there is a theory. predicting a flat patch in the range (= 62. and this is used as a prior.," The problem with this argument is that it is reversed if there is a theory predicting a flat patch in the range $\ell=32-62$ , and this is used as a prior."
 Then Caussianity is ruled out at the 99.6% confidence level., Then Gaussianity is ruled out at the $99.6\%$ confidence level.
 So the prior is all important., So the prior is all important.
 As the example of COBE-DAIR. shows the threat of systematics should make us beware of simply dismissing anomalies as Iukes., As the example of COBE-DMR shows the threat of systematics should make us beware of simply dismissing anomalies as flukes.
 We stress that none of these remarks apply to the results in Section ??.. concerning North-South asvmamoetries in the normalised inter-£ bispectrum.," We stress that none of these remarks apply to the results in Section \ref{asyminter}, concerning North-South asymmetries in the normalised $\ell$ bispectrum."
 We were particularly careful to compare like with like - letting the Gaussian simulations choose the axis of maximal symmetry. just as we did with the data.," We were particularly careful to compare like with like - letting the Gaussian simulations choose the axis of maximal symmetry, just as we did with the data."
 Phe fact that the maximal asymmetry axis correlates well with that found for the power spectrum asvmmeltry reinforces the detection., The fact that the maximal asymmetry axis correlates well with that found for the power spectrum asymmetry reinforces the detection.
 For a Ciaussian process power spectrum and normalised bispectrum are independent random variables: one would need two independent. Dukes instead. of one., For a Gaussian process power spectrum and normalised bispectrum are independent random variables: one would need two independent flukes instead of one.
 In this paper we have performed tests of Ciaussianity ancl explored possible asvnimetries of the WALAD first vear data., In this paper we have performed tests of Gaussianity and explored possible asymmetries of the WMAP first year data.
 We calculated the inter-f. 47. and. single-£. £7. normalised bispectrum ancl compared. the WALAP results with those from simulations.," We calculated the $\ell$, $J^3_\ell$, and $\ell$, $I^3_\ell$, normalised bispectrum and compared the WMAP results with those from simulations."
 We found the whole sky results o be consistent with Ciaussianity. barring a flat. patch’ anomaly in the £7 for (=3262.," We found the whole sky results to be consistent with Gaussianity, barring a 'flat patch' anomaly in the $I^3_\ell$ for $\ell=32-62$."
 Depending on our attituce oward priors this may or may not be seen as rejecting CGaussianityv at the 99.6% confidence level., Depending on our attitude toward priors this may or may not be seen as rejecting Gaussianity at the $99.6\%$ confidence level.
 We also. searched. for. hemisphere asvmametries. in he bispectrum. and using the AJ; found a North-South asvinmetry. with significance. level maximised at [or a North pole at. (f.6) (57.10).," We also searched for hemisphere asymmetries in the bispectrum, and using the $J^3_\ell$ found a North-South asymmetry, with significance level maximised at for a North pole at $(\ell,b)\approx$ (57,10)."
 This is the same North bole Chat maximises angular-power spectrum. asymmetry (Eriksen&al2004a:Hansen.Bandayιόν2004).," This is the same North pole that maximises angular-power spectrum asymmetry \citep{erik1,hbg}."
. Ao similar analysis with £5 showed consisteney— with Caussianity., A similar analysis with $I^3_\ell$ showed consistency with Gaussianity.
 We noted that asvmametries in the 3-point. correlation function. C(8). (Eriksen&al2004a) do not diminish when one excludes low-f contributions.," We noted that asymmetries in the 3-point correlation function, $C^3(\theta)$, \citep{erik1} do not diminish when one excludes $\ell$ contributions."
 However such asvmmetries are unlikely to be related to those reported in this paper., However such asymmetries are unlikely to be related to those reported in this paper.
 Our quacrupole analysis failed to find any correlation with previously reported anomalies (Schwarz&al2004:Oliveira-Costa&al 2004).," Our quadrupole analysis failed to find any correlation with previously reported anomalies \citep{copi1,teg}."
. Whether these features are real or the hallmark of further svstematic errors remains to be seen., Whether these features are real or the hallmark of further systematic errors remains to be seen.
 We would liketo thank AJ. Banday.. M. Kuntz ancl J. Aledeiros for help with this project.," We would liketo thank A.J. Banday, M. Kuntz and J. Medeiros for help with this project."
 Phe results in this paper have been derived. using the LEALPix package 1905)., The results in this paper have been derived using the HEALPix package \citep{healp}. .
"63212, where we could not the emission above 0.09 K inTx.",", where we could not the emission above 0.09 K in."
" At the front end, we used the 25-BEam Array Receiver System (BEARS) in double-sideband (DSB) mode, which has 5x5 beams separated by 41"".1 in the plane of the sky (Sunadaet ."," At the front end, we used the 25-BEam Array Receiver System (BEARS) in double-sideband (DSB) mode, which has $\times$ 5 beams separated by $''$ .1 in the plane of the sky \citep{sun00, yam00} ."
 The 25 beams have beam-to-beam variations of about in both beam efficiency and sideband ratio., The 25 beams have beam-to-beam variations of about in both beam efficiency and sideband ratio.
" To correct for the beam-to-beam gain variations, we calibrated the intensity scale of each beam by observing the W3 region with a 100 GHz SIS receiver (8100) with a single-sideband filter."," To correct for the beam-to-beam gain variations, we calibrated the intensity scale of each beam by observing the W3 region with a 100 GHz SIS receiver (S100) with a single-sideband filter."
" At the back end, we used 25 sets of 1024 channel autocorrelators (ACs), which have a velocity resolution of 0.104 km sl at 110 GHz (Soraietal.2000)."," At the back end, we used 25 sets of 1024 channel autocorrelators (ACs), which have a velocity resolution of 0.104 km $^{-1}$ at 110 GHz \citep{sor00}."
". Since the data dumping time of the ACs was 0.1 s, the spatial data sampling interval on the sky plane was 3.5""."," Since the data dumping time of the ACs was 0.1 s, the spatial data sampling interval on the sky plane was $''$."
" The spatial interval corresponds to0.3A@yppw,, and is small enough to satisfy the Nyquist theorem."," The spatial interval corresponds to, and is small enough to satisfy the Nyquist theorem."
" The telescope pointing was checked every 1.5 hours by observing the SiO (v—1, J—1-0; 43.122 GHz) maser source 'T-Cep."," The telescope pointing was checked every 1.5 hours by observing the SiO $v$ =1, $J$ =1–0; 43.122 GHz) maser source T–Cep."
 Since the pointing uncertainty of the telescope remains as small as a few arcseconds below a wind speed of 5 m sl. we rejected the data taken under the condition that the wind speed averaged over one minute exceeds 6 m s!.," Since the pointing uncertainty of the telescope remains as small as a few arcseconds below a wind speed of 5 m $^{-1}$, we rejected the data taken under the condition that the wind speed averaged over one minute exceeds 6 m $^{-1}$."
" Consequently, the pointing accuracy is better than 3"" for the data to be analyzed."," Consequently, the pointing accuracy is better than $''$ for the data to be analyzed."
" To construct an a-ó-vrpsg data cube with spatial and velocity resolutions of 22"" and 0.1 in full width half maximum (FWHM), we used a Gaussian function as a gridding convolution function (GCF) to integrate the spectra which were taken with the very high spatial sampling rate of 3.5""."," To construct an $\alpha$ $\delta$ $v_{\rm LSR}$ data cube with spatial and velocity resolutions of $'' $ and 0.1 in full width half maximum (FWHM), we used a Gaussian function as a gridding convolution function (GCF) to integrate the spectra which were taken with the very high spatial sampling rate of $''$."
" We adopted 17"".0 as the size of the GCF in FWHM.", We adopted $''$ .0 as the size of the GCF in FWHM.
" The resultant effective spatial resolution of the cube, becomes 22"".0, corresponding to 0.1 pc at the distance to the $140 region."," The resultant effective spatial resolution of the cube, becomes $''$ .0, corresponding to 0.1 pc at the distance to the S140 region."
" Since the DSB system noise temperature ranged from 285 to 357 K with a mean value of 321 K during the observations, we have an rms noise level of the data cube of 0.12 K in T3."," Since the DSB system noise temperature ranged from 285 to 357 K with a mean value of 321 K during the observations, we have an rms noise level of the data cube of 0.12 K in ."
" (5) where 2, is the assumed redshift at which star formation begins.","(z) = 1 -, where $z_\star$ is the assumed redshift at which star formation begins."
 Equations| (2)). (2?)). and (5)) can be combined to geive Ilere yyy is the star formation rate of the MW today and fpiiw is the MW. gas mass fraction today.," Equations \ref{general}) ), \ref{ng}) ), and \ref{mu}) ) can be combined to give Here $\psi_{\rm MW}$ is the star formation rate of the MW today and $\mu_{\rm 0,MW}$ is the MW gas mass fraction today."
 The observational determination of the CSFR is based on observations of the cosmic luminosity density in various wavebands which are dominated by emission from short-lived stars: these include the UV (e(e.g.. Lillyefal.(1996);Macau.Pozzetti.&Dickinson (1993))). the IR (e.g.. Chary&Elbaz((2001):Coleetal. (2001))) and IIa (e.g.. Gallegoe£al(1995):Tresse&Maddox(1998);Hopkins.Connolly.Szalayv (2000))).," The observational determination of the CSFR is based on observations of the cosmic luminosity density in various wavebands which are dominated by emission from short-lived stars; these include the UV (e.g., \cite{lilly96, madau98}) ), the IR (e.g., \cite{chary01, cole01}) ) and $\alpha$ (e.g., \cite{gallego95, tm98, hopkins00}) )."
 For 2<1 there is a relatively clear picture of (he CSFR evolution. and results from different. wavebancds generally agree that the CSFR increases to peak around 2=1 reaching a maxinum more than an orcer of magnitude larger than its present value.," For $z<1$ there is a relatively clear picture of the CSFR evolution, and results from different wavebands generally agree that the CSFR increases to peak around $z=1$ reaching a maximum more than an order of magnitude larger than its present value."
 For 2=1—4. however. the uncertainty in the CSER reaches about an order of magnitude. as the derived CSETR values depend stronely on (he assumed amount of dust exünction (see. e.g.. (he compilation of CSFR estimates using different. (racers by Hopkins.efαἱ (2001))).," For $z = 1-4$, however, the uncertainty in the CSFR reaches about an order of magnitude, as the derived CSFR values depend strongly on the assumed amount of dust extinction (see, e.g., the compilation of CSFR estimates using different tracers by \cite{hopkins01}) )."
" In computing the normal galaxy. contribution to the EGRB we will use the analytic fit ol the CSFR evolution given bv. Coleetal.(2001).. based on near-infrared observations. a Salpeter mass Function and an O4=0.7. O,,=0.3 cosmology."," In computing the normal galaxy contribution to the EGRB we will use the analytic fit of the CSFR evolution given by \cite{cole01}, based on near-infrared observations, a Salpeter mass function and an $\Omega_\Lambda=0.7$, $\OmegaM =0.3$ cosmology."
" We will refer to their fit of data points not corrected for dust extinction as the “uncorrected CSFR™ and to (heir fit of data points corrected for dust absorption as the ""dust-corrected CSER.", We will refer to their fit of data points not corrected for dust extinction as the “uncorrected CSFR” and to their fit of data points corrected for dust absorption as the “dust-corrected CSFR.”
The FSV solar-like pulsator HD 49933 (HR 2530. HIP 32851) is one of the primary targets of the CoRoT space mission (?)..,"The F5V solar-like pulsator HD 49933 (HR 2530, HIP 32851) is one of the primary targets of the CoRoT space mission \citep{baglin06}."
 The discovery of its oscillation was made from the ground by observing highly resolved time-series spectra (2).., The discovery of its oscillation was made from the ground by observing highly resolved time-series spectra \citep{mosser05}.
 Since the launch of the spacecraft in December 2006. two runs of 60 and 137 days have been performed leading to a rich p-mode spectrum of 51 frequencies corresponding to the (=0.1.2 eigenmodes (???)..," Since the launch of the spacecraft in December 2006, two runs of 60 and 137 days have been performed leading to a rich p-mode spectrum of 51 frequencies corresponding to the $\ell=0, 1, 2 $ eigenmodes \citep{appourchaux08, benomar09, kallinger10}."
 The interpretation of these eigenmodes to derive constraints on the stellar interior. depends on the accuracy of the determination. of the fundamental parameters of the star (Tay.logg. |[Fe/H]).," The interpretation of these eigenmodes to derive constraints on the stellar interior depends on the accuracy of the determination of the fundamental parameters of the star $\teff, \logg,\feh$ )."
 All chemical abundance analyses of HD 49933 converge to that of a slightly metal-poor star with |Fe/H] between -0.3 and -0.5 (e.g.??).," All chemical abundance analyses of HD 49933 converge to that of a slightly metal-poor star with ${\rm
[Fe/H}]$ between -0.3 and -0.5 \citep[e.g.][]{gillon06, bruntt09}."
 The determination of its effective temperature and gravity have been revisitec several times., The determination of its effective temperature and gravity have been revisited several times.
 As reported in ?.. the Το obtained by different methods ranges from 6450 to 6780 K. which is by far too uncertain to tightly constrain the stellar evolution models in the HR diagram.," As reported in \cite{bruntt09}, , the $\teff$ obtained by different methods ranges from $6450$ to $6780$ K, which is by far too uncertain to tightly constrain the stellar evolution models in the HR diagram."
 Similarly. loge is also not very well constrained with values ranging from 4.0 to 4.3.," Similarly, $\logg$ is also not very well constrained with values ranging from $4.0$ to $4.3$."
 Several authors have attempted to determine Το. and logg by fitting the Ημ and Mg lines. respectively (??)..," Several authors have attempted to determine $\teff$ and $\logg$ by fitting the ${\rm
H_\beta}$ and Mg lines, respectively \citep{ryab09, kallinger10}."
 Their solutions are most consistent with a cooler temperature (=6500 K) and a lower gravity (= 4.0)., Their solutions are most consistent with a cooler temperature $\approx 6500$ K) and a lower gravity $\approx 4.0$ ).
 Since spectrometry and photometry lead to stellar parameters with a large uncertainty. in this work we use a different approach by determining its diameter using interferometry.," Since spectrometry and photometry lead to stellar parameters with a large uncertainty, in this work we use a different approach by determining its diameter using interferometry."
 Combined with seismology. it provides a very good constraint of the stellar parameter space (?)..," Combined with seismology, it provides a very good constraint of the stellar parameter space \citep{creevey07}."
 This approach was proposed for the first time by ? for the binary star e/Cen A&B and later with similar success for other asteroseismic targets (9090909029)," This approach was proposed for the first time by \cite{kervella03} for the binary star $\alpha$ Cen $\&$ B and later with similar success for other asteroseismic targets \citep{cunha07,teixeira09,bruntt09b,mazumdar09,bazot11}."
 In this letter we present the first. campaign. of interferometrie measurements of HD 49933 obtained using the VEGA instrument (??) at the CHARA Array (?)..," In this letter we present the first campaign of interferometric measurements of HD 49933 obtained using the VEGA instrument \citep{vega,vega2} at the CHARA Array \citep{chara}."
 We used state-of-art 3D hydrodynamical simulations of the surface of the star to calculate the limb darkening and to compute the visibility curves prior extracting the angular diameter., We used state-of-art 3D hydrodynamical simulations of the surface of the star to calculate the limb darkening and to compute the visibility curves prior extracting the angular diameter.
 The derived stellar radius is then used to constrain the other fundamental stellar parameters by fitting the oscillation frequency separations., The derived stellar radius is then used to constrain the other fundamental stellar parameters by fitting the oscillation frequency separations.
 Observations of HD 49933 were performed on October 16. 2010 using the VEGA instrument of the CHARA Array and the instrument CLIMB (?) for 3T group delay tracking.," Observations of HD 49933 were performed on October 16, 2010 using the VEGA instrument of the CHARA Array and the instrument CLIMB \citep{climb} for 3T group delay tracking."
The telescopes ΕΙ. E2 and W2 were used giving access to ground baselines of 66. 156. and 221m. The seeing was stable with a value of rO between 10 and 13 em.,"The telescopes E1, E2 and W2 were used giving access to ground baselines of 66, 156, and 221m. The seeing was stable with a value of r0 between $10$ and $13$ cm."
 For the absolute calibration of the squared visibilities. we observed two different calibrators. CI = HD 55185 and C2 = HD 46487.," For the absolute calibration of the squared visibilities, we observed two different calibrators, C1 = HD 55185 and C2 = HD 46487."
 The observing sequence was C1-T-C2-T-C2-T-C1. each block being almost 20mn.," The observing sequence was C1-T-C2-T-C2-T-C1, each block being almost 20mn."
 The data were processed using the standard V2 procedure of the VEGA instrument as described in ?.., The data were processed using the standard $V^2$ procedure of the VEGA instrument as described in \cite{vega2}.
 To validate the absolute calibration. we first considered CI às à target and thus estimated its uniform disk (UD) diameter using C2 as the calibrator.," To validate the absolute calibration, we first considered C1 as a target and thus estimated its uniform disk (UD) diameter using C2 as the calibrator."
 A spectral band of Ati=30 nm centered around ο=735 nm was used for the data processing., A spectral band of $\Delta\lambda=30$ nm centered around $\lambda=735$ nm was used for the data processing.
 The equivalent UD diameter of C2 is taken àsOupjais;= mas according to the SearchCal of the JMMC (?).., The equivalent UD diameter of C2 is taken as$\theta_\mathrm{HD46487}=0.180\pm0.013$ mas according to the SearchCal of the JMMC \citep{Bonneau06}.
 Fitting à UD model to the three squared visibilities provides an estimate of the angular diameter of HD 55185 of, Fitting a UD model to the three squared visibilities provides an estimate of the angular diameter of HD 55185 of
in the scenario of WR wind-bubbles stalled by a high interstellar pressure discussed. by Chevalier (2004).,in the scenario of WR wind-bubbles stalled by a high interstellar pressure discussed by Chevalier (2004).
 For this scenario to explain the general propertics of the radio and optical emissions of the afterglows 990123 and 021211. the magnetic field and electron energy. parameters would have to decrease and increase. respectively. by a factor of about LOO at /=/.. when the wind termination shock is reached. (figures 4 and 5)). a contrived feature without a physical foundation.," For this scenario to explain the general properties of the radio and optical emissions of the afterglows 990123 and 021211, the magnetic field and electron energy parameters would have to decrease and increase, respectively, by a factor of about 100 at $t=t_*$, when the wind termination shock is reached (figures \ref{janwb} and \ref{decwb}) ), a contrived feature without a physical foundation."
 We also find that the wind-bubble scenario requires winds which are as tenuous as those for the reverse-forward shock scenario., We also find that the wind-bubble scenario requires winds which are as tenuous as those for the reverse-forward shock scenario.
 Thus the reverse-forwarcl shock scenario. provides a more natural explanation than the wind-bubble scenario for the steep carly decay of the optical emission of the afterglows 990123 and 021211., Thus the reverse-forward shock scenario provides a more natural explanation than the wind-bubble scenario for the steep early decay of the optical emission of the afterglows 990123 and 021211.
 Given that the GRB ejecta can be initially magnetized and that the RS is less relativistic than he FS. the microphysical parameters might diller behind he two shocks.," Given that the GRB ejecta can be initially magnetized and that the RS is less relativistic than the FS, the microphysical parameters might differ behind the two shocks."
 If their equality is required. for a. simpler scenario. with fewer assumptions. then a wind-like CDM is avoured by the reverse-lorward shock scenario. though a woblen still exists: the low wind density inferred. in each case (ch.< 0.1). which is similar to that derived by Chevalier (2004) for the afterglows 020405 and 021211 and by ice (2002) for the afterglow 011211.," If their equality is required for a simpler scenario, with fewer assumptions, then a wind-like CBM is favoured by the reverse-forward shock scenario, though a problem still exists: the low wind density inferred in each case $A_* < 0.1$ ), which is similar to that derived by Chevalier (2004) for the afterglows 020405 and 021211 and by Price (2002) for the afterglow 011211."
" In the sample of 64 Galactic WR. stars analyzed by Nugis Lamers (2000) here isonly one star with ch,<0.1. the majority. of he other stars having a mass-loss rate ALC(0.5.7) a wind velocity (à,(10003000)kms +"," In the sample of 64 Galactic WR stars analyzed by Nugis Lamers (2000) there isonly one star with $A_* < 0.1$, the majority of the other stars having a mass-loss rate $\dot{M} \in (0.5 - 7) 
\times 10^{-5} \Msunyear$ , a wind velocity $v_w \in (1000-3000)\, \kms$ ,"
" In the sample of 64 Galactic WR. stars analyzed by Nugis Lamers (2000) here isonly one star with ch,<0.1. the majority. of he other stars having a mass-loss rate ALC(0.5.7) a wind velocity (à,(10003000)kms +."," In the sample of 64 Galactic WR stars analyzed by Nugis Lamers (2000) there isonly one star with $A_* < 0.1$, the majority of the other stars having a mass-loss rate $\dot{M} \in (0.5 - 7) 
\times 10^{-5} \Msunyear$ , a wind velocity $v_w \in (1000-3000)\, \kms$ ,"
"The solar wind is a weakly collisional plasma(e.g.,Kasperal.2008) that is ubiquitously observed to be in a turbulent state (Tu&Marsch1995;Goldsteinetal.HorburyMatthaeus&Velli 2011).","The solar wind is a weakly collisional plasma\citep[e.g.,][]{kasper08} that is ubiquitously observed to be in a turbulent state \citep{tu95,goldstein95a,horbury05,bruno05a,petrosyan10,matthaeus11}."
" Much progress has been made in understanding the nature of this turbulence since the first direct spacecraft observations (e.g.,Siscoeetal.1968;Cole-man1968) but many aspects remain to be fully understood."," Much progress has been made in understanding the nature of this turbulence since the first direct spacecraft observations \citep[e.g.,][]{siscoe68,coleman68} but many aspects remain to be fully understood."
" In particular, the three-dimensional (3D) structure has been poorly characterized."," In particular, the three-dimensional (3D) structure has been poorly characterized."
" Here, we use a new single-spacecraft to measure the 3D structure of turbulence in the fast solar techniquewind."," Here, we use a new single-spacecraft technique to measure the 3D structure of turbulence in the fast solar wind."
" Turbulence is usually modeled as a local cascade of fluctuations from large to small scales, forming an inertial "," Turbulence is usually modeled as a local cascade of fluctuations from large to small scales, forming an inertial range."
"In the solar wind, most of the at scales is in range.Alfvénnic fluctuations (Belcher&energyDavis1971;largeBrunoetal.1985;Horbury1995;Bale 2005),, which have magnetic field and velocity fluctuations perpendicular to the magnetic field direction (Alfvén1942)."," In the solar wind, most of the energy at large scales is in Alfvénnic fluctuations \citep{belcher71,bruno85,horbury95,bale05}, which have magnetic field and velocity fluctuations perpendicular to the magnetic field direction \citep{alfven42}."
 Early isotropic magnetohydrodynamic (MHD) turbulence theories (Iroshnikov1963;Kraichnan1965) based on Kolmogorov scaling arguments(Kolmogorov1941) predict that the energy of weak Alfvénnic turbulence is E(k)~k3?. where kis the spectrumwavenumber of the fluctuations.," Early isotropic magnetohydrodynamic (MHD) turbulence theories \citep{iroshnikov63,kraichnan65} based on Kolmogorov scaling arguments\citep{kolmogorov41a} predict that the energy spectrum of weak Alfvénnic turbulence is $E(k)\sim k^{-3/2}$, where $k$ is the wavenumber of the fluctuations."
" Although 1D velocity power spectra in the solar windat 1 AU display this scaling (Mangeneyetal.2001;Podesta2007;Salem2009;2011a;Boldyrevetal. 2011),, the magnetic field has a k?/? scaling (e.g.,Matthaeus&Goldstein1982;etal.2006;Chen2011a;Boldyrev 2011).."," Although 1D velocity power spectra in the solar windat 1 AU display this scaling \citep{mangeney01,podesta07a,salem09,chen11b,boldyrev11}, the magnetic field has a $k^{-5/3}$ scaling \citep[e.g.,][]{matthaeus82a,smith06a,chen11b,boldyrev11}."
 It was later realized (Montgomery&Turner1981;She-balinetal.1983) that the magnetic field direction can induce anisotropy in plasma turbulence.," It was later realized \citep{montgomery81,shebalin83} that the magnetic field direction can induce anisotropy in plasma turbulence."
" It was then proposed (Higdon1984;Goldreich&Sridhar1995) that Alfvénnic turbulence tends towards a state of critical balance, in which the timescale of the Alfvénnic fluctuations propagating along the magnetic field is equal to the timescale of their nonlinear decay."," It was then proposed \citep{higdon84,goldreich95} that Alfvénnic turbulence tends towards a state of critical balance, in which the timescale of the Alfvénnic fluctuations propagating along the magnetic field is equal to the timescale of their nonlinear decay."
" This produces a spectrum perpendicular to the local magnetic field of E(k.)~ke? Pa parallel spectrum of E(k)~ and local wavevector scaling Kj~RP, "," This produces a spectrum perpendicular to the local magnetic field of $E(k_\perp)\sim k_\perp^{-5/3}$, a parallel spectrum of $E(k_\parallel)\sim k_\parallel^{-2}$ and local wavevector scaling $k_\parallel\sim k_\perp^{2/3}$."
Solar wind turbulenceki? measurements show evidence for both wavevector anisotropy of the form Κι> (Crookeretal.Wicksetal.2011;Horbury2011) and a steeper spectral1b; index parallel to the local magnetic field (Horburyetal.2008;2011b;Wicksetal.2011;Horbury 2011). ," Solar wind turbulence measurements show evidence for both wavevector anisotropy of the form $k_\perp>k_\parallel$ \citep{crooker82,bieber96,leamon98a,horbury08,podesta09a,wicks10a,chen11a,wicks11a,horbury11} and a steeper spectral index parallel to the local magnetic field \citep{horbury08,podesta09a,luo10,wicks10a,chen11a,wicks11a,horbury11}. ."
Critical balance theory was later extended to allow for the possibility that Alfvénnic turbulence is 3D anisotropic (Boldyrev2006)., Critical balance theory was later extended to allow for the possibility that Alfvénnic turbulence is 3D anisotropic \citep{boldyrev06}.
". The two special orthogonal directions are the mean magnetic field Bo and the perpendicular magnetic field fluctuation óB, .", The two special orthogonal directions are the mean magnetic field $\mathbf{B}_0$ and the perpendicular magnetic field fluctuation $\delta\mathbf{B}_\perp$ .
" The theory assumes that the magnetic fieldand velocity fluctuations align to within a scale dependent angle 0,5, which makes them 3D anisotropic: l2€» A, where /, € and A are their correlation lengths"," The theory assumes that the magnetic fieldand velocity fluctuations align to within a scale dependent angle $\theta_{vb}$ , which makes them 3D anisotropic: $l>\xi>\lambda$ , where $l$ , $\xi$ and $\lambda$ are their correlation lengths"
The predictions of the cold dark matter (CDM). models successfully account for several observational facts such as the clistribution of matter on large scales. the uniformity of the cosmic microwave radiation and its small temperature anisotropies. and the measured. values of the cosmological porameters.,"The predictions of the cold dark matter (CDM) models successfully account for several observational facts such as the distribution of matter on large scales, the uniformity of the cosmic microwave radiation and its small temperature anisotropies, and the measured values of the cosmological parameters."
 Llowever. on small scales the predictions seem to be in conflict with observations sugeesting that modifications to the standard scenario should be introduced.," However, on small scales the predictions seem to be in conflict with observations suggesting that modifications to the standard scenario should be introduced."
 One of the problems of the CDM. theory is that. the inner density profile of the virialized haloes is too steep with respect to what is inferred. Crom rotation curves of dwarf spiral galaxies (Moore 1994: Flores Primack 1994: δικοί 1995)., One of the problems of the CDM theory is that the inner density profile of the virialized haloes is too steep with respect to what is inferred from rotation curves of dwarf spiral galaxies (Moore 1994; Flores Primack 1994; Burkert 1995).
 In fact it is well known that these galaxies are systems stronely dominated. by dark matter. so that heir rotation curves are good tracers of the dark halo eravitational potential.," In fact it is well known that these galaxies are systems strongly dominated by dark matter, so that their rotation curves are good tracers of the dark halo gravitational potential."
 A similar situation is expected. for ow surface brightness (LSB) galaxies (de Blok AleGaugh 1997). (although the observational evidence of a soft core in hese cases has been challenged by some authors e.g... van den Bosch et al 1999: Swaters. Macdore & Trewhella 2000).," A similar situation is expected for low surface brightness (LSB) galaxies (de Blok McGaugh 1997), (although the observational evidence of a soft core in these cases has been challenged by some authors e.g., van den Bosch et al 1999; Swaters, Madore $\&$ Trewhella 2000)."
 Sumilarlv. Hernánndez & Cilmore (1998) showed that. the IHE rotation curves of both LSB and normal large ealaxies can be characterized by a significant soft core inner region.," Similarly, Hernánndez $\&$ Gilmore (1998) showed that the HI rotation curves of both LSB and normal large galaxies can be characterized by a significant soft core inner region."
 The Tullvy-Eisher relation can be also better. predicted. by galaxy formation models when a density. profile shallower, The Tully-Fisher relation can be also better predicted by galaxy formation models when a density profile shallower
h,.
yperon polarization and the single-," However, they do not explain"
in (he profile ancl its (wo neighbors. using 200 phase bins: fitting a Lorentzian function to ihe phase range 0.98 to 1.00. using 400 phase bins: ancl calculating the first moment over that part of the pulse where the bins exceed of the highest bin. using 800 phase bins.,"in the profile and its two neighbors, using 200 phase bins; fitting a Lorentzian function to the phase range 0.98 to 1.00, using 400 phase bins; and calculating the first moment over that part of the pulse where the bins exceed of the highest bin, using 800 phase bins."
 This procedure allows us to determine the pulse phase in an individual observation with an accuracy ol 1 milliperiod., This procedure allows us to determine the pulse phase in an individual observation with an accuracy of 1 milliperiod.
 We did investigate whether (he time resolution of the observations gives rise (o an additional systematic error ancl found this not to be the case: Lorentzian fits to two back-to-back observations made in the fall of 2003 with resolutions of 250 ji and 16 yes. respectively. differed by less than 0.1 milliperiod.," We did investigate whether the time resolution of the observations gives rise to an additional systematic error and found this not to be the case: Lorentzian fits to two back-to-back observations made in the fall of 2003 with resolutions of 250 $\mu$ s and 16 $\mu$ s, respectively, differed by less than 0.1 milliperiod."
 The absolute phase of the peak of the X-ray main pulse (with respect to the peak of the radio main pulse. using the Lorentzian fits). as a function of lime (in Modified Julian Days) is shown in Fig. 2:," The absolute phase of the peak of the X-ray main pulse (with respect to the peak of the radio main pulse, using the Lorentzian fits), as a function of time (in Modified Julian Days) is shown in Fig. \ref{fig2};"
 the errors are a combination of the 1 milliperiod error mentioned above and (he rms deviations in the fits of the radio (imine ephemerices and (vpically. amount to 40 yrs. Independently from these statistical errors there is the svstematic error of up to 40 pes introduced by the radio receiver svstem and calibration. as mentionedin the previous seciion.," the errors are a combination of the 1 milliperiod error mentioned above and the rms deviations in the fits of the radio timing ephemerides and typically amount to 40 $\mu$ s. Independently from these statistical errors there is the systematic error of up to 40 $\mu$ s introduced by the radio receiver system and calibration, as mentionedin the previous section."
 Fi, Fig.
e.2 shows the history of (he main X-ray. pulse phase. with the occurrences of glitches marked.," \ref{fig2} shows the history of the main X-ray pulse phase, with the occurrences of glitches marked."
 Glitches 7 through 12 are taken [rom Wong.Backer.&Lyne(2001).. while the elitch numbered 8 corresponds to Note 12 on the Jodrell Bank web page and adding one to glitch numbers 12 through 18 vields (he corresponding Note numbers (13 through 19) on that page.," Glitches 7 through 12 are taken from \citet{wong2001}, while the glitch numbered 8 corresponds to Note 12 on the Jodrell Bank web page and adding one to glitch numbers 12 through 18 yields the corresponding Note numbers (13 through 19) on that page."
 We have monitored the X-ray. emission more closely folowing (wo glitches. but found no unusual behavior.," We have monitored the X-ray emission more closely following two glitches, but found no unusual behavior."
 The data points in this figure can be divided into four equality categories., The data points in this figure can be divided into four quality categories.
 First. Cae data points that are based on timing ephemerides prior to MJD 50870 obviously display. larger deviations (han the later ones: we attribute this to the poorer quality of Chose radio ephemeris records.," First, the data points that are based on timing ephemerides prior to MJD 50870 obviously display larger deviations than the later ones; we attribute this to the poorer quality of those radio ephemeris records."
 Second. there are a number of points with high error estimates (75 milliperiods). associated. with elitches: the monthly Gmine ephemeris records are nol sulliciently to handle elitehes properly.," Second, there are a number of points with high error estimates $>$ 5 milliperiods), associated with glitches; the monthly timing ephemeris records are not sufficiently fine-grained to handle glitches properly."
 Third. there are seven outlvers (below phase 0.9860) that are clearly well below the remaining data points.," Third, there are seven outlyers (below phase 0.9860) that are clearly well below the remaining data points."
 In all cases these represent either all observations covered by a single ephemeris record or observations at the edge of such a record: hence. we attribute these to inaccuracies in the timing ephemerides.," In all cases these represent either all observations covered by a single ephemeris record or observations at the edge of such a record; hence, we attribute these to inaccuracies in the timing ephemerides."
 Fourth. (he remaining 111 data points form a normal distribution with the expected rms scatter of 1 milliperiod.," Fourth, the remaining 111 data points form a normal distribution with the expected rms scatter of 1 milliperiod."
 This indicates that the data are consistent with a constant value., This indicates that the data are consistent with a constant value.
 The data points from the first three categories are represented by open circles in Fig. 2:, The data points from the first three categories are represented by open circles in Fig. \ref{fig2}; ;
 the high-quality data points, the high-quality data points
mean PSF thus derived was then applied to the expected true count. distribution. i.e. the product of the model of ealactic emission. GALDIF. and the exposure map. EX. to derive the expected count distribution iu the data. C(0..0))= GALDIF((a..;3)) where a aud 3 deuote the true augular coordinates of the incoming 5-ravs.. whereas 0 aud © are the measured angular coordinates of the5-rays.,"mean PSF thus derived was then applied to the expected true count distribution, i.e. the product of the model of galactic emission, GALDIF, and the exposure map, EX, to derive the expected count distribution in the data, )= ) where $\alpha$ and $\beta$ denote the true angular coordinates of the incoming , whereas $\theta$ and $\phi$ are the measured angular coordinates of the."
. The staudard maps of diffuse galactic enission. that cau be obtained (rom the CORO Scieuce Support Center. are calculatecl asstuning a single power-law spectrum with iudex 2.1. thus implying a PSF of different width that used iu my analysis.," The standard maps of diffuse galactic emission, that can be obtained from the CGRO Science Support Center, are calculated assuming a single power-law spectrum with index 2.1, thus implying a PSF of different width that used in my analysis."
 Table 1 displays a list of all viewing periods considered in this study., Table \ref{t1} displays a list of all viewing periods considered in this study.
 For each viewing period the data have been analyzed in the enerey baucds 100-300 MeV. 300-1.000 MeV. 2100 MeV. 2300 MeV. aud 1.000 MeV. I have also constructed a combined data set for all viewing periods with a view to derive a detailed spectrum aud to pinpoint the localization of the source.," For each viewing period the data have been analyzed in the energy bands 100-300 MeV, 300-1.000 MeV, $>$ 100 MeV, $>$ 300 MeV, and $>$ 1.000 MeV. I have also constructed a combined data set for all viewing periods with a view to derive a detailed spectrum and to pinpoint the localization of the source."
 The spectra of EGRET sources are usually determined ou the basis of a likelihood analysis in ten standard energy ranges., The spectra of EGRET sources are usually determined on the basis of a likelihood analysis in ten standard energy ranges.
" Table 2. lists the integrated. photon flux aud the pE, flux iu the teu uarrow energy baucds aud five wide energy bauds based ou a likelihood analysis of the combined data set."," Table \ref{t2} lists the integrated photon flux and the $\nu\,F_\nu$ flux in the ten narrow energy bands and five wide energy bands based on a likelihood analysis of the combined data set."
 The spectrum depeuds very weakly ou the assumed position of the source. if that is varied by lS0.1? in the Galactic Plane.," The spectrum depends very weakly on the assumed position of the source, if that is varied by $\delta l \lesssim 0.1^\circ$ in the Galactic Plane."
 Figure 1 shows the pΕν spectrum above 70 MeV. The parameters oL the diffuse foregrouud model. νι aud Gp. were allowed to vary lreely.," Figure \ref{f1} shows the $\nu\,F_\nu$ spectrum above 70 MeV. The parameters of the diffuse foreground model, $\gm$ and $\gb$, were allowed to vary freely."
 Table 2. also lists the integrated photon flux for all energy bauds that oue obtains using the standard maps ofdilluse [oregrouud emission., Table \ref{t2} also lists the integrated photon flux for all energy bands that one obtains using the standard maps of diffuse foreground emission.
 The combined data set has au exposure distribution that peaks within 2° of the Galactic Center αμα is fairly symmetric., The combined data set has an exposure distribution that peaks within $^\circ$ of the Galactic Center and is fairly symmetric.
 The staucdard maps aud the moclilied maps of galactic απο emission differ only. marginally., The standard maps and the modified maps of galactic diffuse emission differ only marginally.
 The photon flux for the enerey bauds below 100 MeV is strongly dependent ou the particulars of the likelihood rum. e.g. on the analysis parameter RANAL. and thus discrecdited.," The photon flux for the energy bands below 100 MeV is strongly dependent on the particulars of the likelihood run, e.g. on the analysis parameter RANAL, and thus discredited."
 An anticorrelation between the multiplier Gaya of the diffuse emission moclel aud the photon flux attributed to can be observed. which iudicates a similarity ofthe expected augular distributions of events for the diffuse emission aud16-2851. respectively.," An anticorrelation between the multiplier $\gm$ of the diffuse emission model and the photon flux attributed to can be observed, which indicates a similarity of the expected angular distributions of events for the diffuse emission and, respectively."
 The results of the likelihood analysis at energiesbelow 100 MeV do therefore not merit our cousicderation in the variability aud localization study., The results of the likelihood analysis at energiesbelow 100 MeV do therefore not merit our consideration in the variability and localization study.
As noted above. in the Poisson's equation. the density perturbations giving rise to the gravitational potential perturbations can be due to compressibility and. due to the stratification of the equilibrium: vertical structure.,"As noted above, in the Poisson's equation, the density perturbations giving rise to the gravitational potential perturbations can be due to compressibility and due to the stratification of the equilibrium vertical structure."
 As is evident. [rom Fig., As is evident from Fig.
 2c. in non-self-eravitating discs. the r-miocle is nearly incompressible.," 2c, in non-self-gravitating discs, the r-mode is nearly incompressible."
 So. it is interesting to see if it still remains incompressible in scll-eravitating disces and which out of these two sources of density perturbations 1s ultimately responsible for the gravitational instability.," So, it is interesting to see if it still remains incompressible in self-gravitating discs and which out of these two sources of density perturbations is ultimately responsible for the gravitational instability."
 In order to explore this. we again computed the dispersion curve of the basic even r-mode in the incompressible limit of equations (16-18). which take the form: supplemented with the same boundary. conditions.," In order to explore this, we again computed the dispersion curve of the basic even r-mode in the incompressible limit of equations (16-18), which take the form: supplemented with the same boundary conditions."
 Notice that on the right hand side of Poisson equation (25). only the density perturbation due to stratifications remains.," Notice that on the right hand side of Poisson's equation (25), only the density perturbation due to stratification remains."
 In Fig., In Fig.
 S. we compare the dispersion. curves of the basic even r-mode obtained in the incompressible limit to those computed for the compressible sell-gravitating and non-self-eravitating cases.," 8, we compare the dispersion curves of the basic even r-mode obtained in the incompressible limit to those computed for the compressible self-gravitating and non-self-gravitating cases."
" Lt is clear that in self-gravitating disces. the basic even remode becomes compressible at οAy, thereby providing density perturbation for the gravitational potential and. therefore. accounting for the large dip on the dispersion curve."," It is clear that in self-gravitating discs, the basic even r-mode becomes compressible at $k \sim k_m$ thereby providing density perturbation for the gravitational potential and, therefore, accounting for the large dip on the dispersion curve."
 “Lhe dispersion curve in the incompressible limit cleviates only very slightlv from that in the non-self-eravitating case. because of the second source of density perturbation. ie. stratification. which turns out to be much smaller than that associated. with compressibility.," The dispersion curve in the incompressible limit deviates only very slightly from that in the non-self-gravitating case, because of the second source of density perturbation, i.e. stratification, which turns out to be much smaller than that associated with compressibility."
 Such a comparison also shows that the primary cause of eravitational instability in incompressible discs. as described in GLB and 2.. is the displacements. of frec-surfaces.," Such a comparison also shows that the primary cause of gravitational instability in incompressible discs, as described in GLB and \cite{LP93b}, is the displacements of free-surfaces."
 In our case. the equilibrium density vanishes at the surface. so that the effect. of surface displacements on the growth rate of instability is null (see boundary. condition 22) and in the incompressible case only density perturbation due to stratification. as shown. is hardly sullicient for gravitational instability.," In our case, the equilibrium density vanishes at the surface, so that the effect of surface displacements on the growth rate of instability is null (see boundary condition 22) and in the incompressible case only density perturbation due to stratification, as shown, is hardly sufficient for gravitational instability."
 Llere we compute the vertical structure of the eigenfunctions of the eravitationally unstable basic even r-mode and. also [ind what type of motions it induces., Here we compute the vertical structure of the eigenfunctions of the gravitationally unstable basic even r-mode and also find what type of motions it induces.
 We take Q=0.1 and s=5. which gives f=1.22 for the disc height.," We take $Q=0.1$ and $s=5$, which gives $h=1.22$ for the disc height."
 For these parameters. the dispersion. curve of this moce in Fig.," For these parameters, the dispersion curve of this mode in Fig."
" δ has a minimum ""M=015. Le. gives the largest. &rowth rate of gravitational instability. at A,=1."," 8 has a minimum $\omega_{min}^2=-0.75$, i.e. gives the largest growth rate of gravitational instability, at $k_{m}=1$."
 In Fig., In Fig.
 9. we plot the corresponding cigenfunctions onlv in the upper half of clise’s full vertical extent.," 9, we plot the corresponding eigenfunctions only in the upper half of disc's full vertical extent."
 Because the mode has even parity. the vertical velocity ancl the derivative of the gravitational potential are odd. functions and. correspondingly. the pressure ancl potential are even functions of z.," Because the mode has even parity, the vertical velocity and the derivative of the gravitational potential are odd functions and, correspondingly, the pressure and potential are even functions of $z$."
 Only the pressure perturbation has one node in the interval 0<2xf. other quantities have no nodes anc vary over the whole vertical extent.," Only the pressure perturbation has one node in the interval $0 < z \leq h$, other quantities have no nodes and vary over the whole vertical extent."
 This demonstrates that the spatial structure of the gravitationally most. unstable moce is three-climensional with non-zero vertical velocity., This demonstrates that the spatial structure of the gravitationally most unstable mode is three-dimensional with non-zero vertical velocity.
 To see how sell-gravitv changes the form of the eigenfunctions. we also show the eigenfunctions of the same branch in the non-self-egravitating limit. which appear to be more," To see how self-gravity changes the form of the eigenfunctions, we also show the eigenfunctions of the same branch in the non-self-gravitating limit, which appear to be more"
hTable 5. reports the \? of. the two zones ofB the new physical. [unction.B. their. sum and 4?E ol the Schechter luminosity fnnction.,"Table \ref{chisquare_data} reports the $\chi^2$ of the two zones of the new physical function, their sum and $\chi^2$ of the Schechter luminosity function."
viscosities. respectively.,"viscosities, respectively."
 The droplet formation rate can be expressed as with μμ ds the quark chemical potential in unitsof 400 MeV. The total number No of droplets formed by nucleation is (he integrated rate over the volume of the neutron star ancl the time after (he nucleation process star(s at the moment Ky N=hlAndy[7INp)T()]dl. where R is the radius of the neutron star.," The droplet formation rate can be expressed as with $\mu _{400}$ is the quark chemical potential in unitsof $400$ MeV. The total number $N$ of droplets formed by nucleation is the integrated rate over the volume of the neutron star and the time after the nucleation process starts at the moment $t_0$, $N=\int_{0}^{R}4\pi r^{2}dr\int_{t_{0}}^{\infty }I\left[ \Delta
p\left( r\right) ,T\left( t\right) \right]dt$, where $R$ is the radius of the neutron star."
 The pressure and temperature depend sensitivelv on the equation of state of the nuclear and quark matter., The pressure and temperature depend sensitively on the equation of state of the nuclear and quark matter.
 To convert the core of the neutron star into quark matter al least one clroplet must be formed. i.e... No 1.," To convert the core of the neutron star into quark matter at least one droplet must be formed, i.e., $N>1$ ."
 Then. with the use of Eq. (3)).," Then, with the use of Eq. \ref{eqn}) ),"
 the condition lor the formation . ⋅ ⋅ ⊋∶⊾≱⋡↓∶⊾≱⋡ ∪≀↧↴↥↥≼↲≀↕⊔∖⊽↥∪∐≼↲≼∐⋅∪↕↽≻↥≼↲↥↕⋅≼↲≺⇂∏∐⋅≼↲⋟∖⊽⊔≤−≻∔⇀∖↕↩∖↓∐↓−⋟∆∍↿⋅⋡↓∣∣↕⇁↿⋅⋡↓∣∣⋅∖∖⇁↥∐↲↕⋅≼↲∫∍↿⋅≀↧↴∐≺⊔↴⋅≀↕↴↕⋅≼↲⊔∐↲≺∢∪↕⋅≼↲ ↗ ↽≀↧↴↥⋯↲⋟∖⊽∪↓⋟⊔∐↲↕↽≻↕⋅≼↲⋟∖⋱∖⊽∏↕⋅≼↲≀↕↴↕∐⇂∩↲∐↕↕↽≻≼↲↕⋅≀↧↴⊓∐⋅≼↲," the condition for the formation of at least one droplet requires $\sigma \leq 24$ MeV $^{-2}\Delta
P_{c,10}^{2/3}T_{c,10}^{1/3}$, where $P_{c}$ and $T_{c}$ are the core values of the pressure and temperature, respectively."
⋅↕⋅≼↲⋟∖⊽↕↽≻≼↲≺∢∐∖⇁≼↲↥∡∖↽⋖↜∐≼↲↕⋟∖⊽≼↲∐↽≻≼↲↕⋅≸≟⊥≤∍≤∍↱≻↕⋝⋅⋅↴⊺∐≼↲↕↽≻↕⋅≼↲⋟∖⋱∖⊽∏↕⋅≼↲≀↧↴∐≺⇂ ≼↲↥⊔⋯↴∏↽≻∡∖↽≺∐∐⋡≼↲↕⋅≼↲∐≺∢≼↲≼⇂≼↲↕↽≻≼↲∐≼⇂⋟∖⊽⋟∖⊽⊔⋅∪∐↖⊂↽↔↴↥∡∖↽∪∐⊔∐↲≼↲≺⇂∏≀↧↴∐∪∐∪↓⋟⋟∖⇁↥≀↧↴∩↲∪↓⋟∐∏≺∢↥≼↲≀↧↴↕⋅≀↧↴∐≺⇂≺⇂∏≀↧↴↕⋅↳↽∐⋯⊔≼↲↕⋅ xl hence the droplet formation rate cannot be reliably estimated., The pressure and enthalpy difference depends strongly on the equation of state of nuclear and quark matter and hence the droplet formation rate cannot be reliably estimated.
 The effect of subcritical hadron bubbles on an inhomogeneous first order euark-hadron phase (transition was studied in(2000)., The effect of subcritical hadron bubbles on an inhomogeneous first order quark-hadron phase transition was studied in.
. The transition from nuclear matter (consisting of neutrons. protons and electrons) to matter containing strangeness was considered. within a mean field type model and with the use of Langer nucleation (lieory. by(2002).," The transition from nuclear matter (consisting of neutrons, protons and electrons) to matter containing strangeness was considered, within a mean field type model and with the use of Langer nucleation theory, by."
. An estimate of the time the new phase to appear at various densities and times in the cooling history of a protoneutron star was also given., An estimate of the time the new phase to appear at various densities and times in the cooling history of a protoneutron star was also given.
 On the other hand. as à result of an inerease of the central density of the star. due lor exaniple to accretion or spinning down. a metastable. super-compressed neutron phase can appear. wilh a star consisting of the new phase surrounded by normal neutron matter1993).," On the other hand, as a result of an increase of the central density of the star, due for example to accretion or spinning down, a metastable, super-compressed neutron phase can appear, with a star consisting of the new phase surrounded by normal neutron matter."
. This situation is similar {ο a eas undergoing a phase transition to its liquid state if compressed (o a volume V. at fixed 7., This situation is similar to a gas undergoing a phase transition to its liquid state if compressed to a volume $V_{c}$ at fixed $T$.
 Iit is slowly compressed. it may stay in (he vapor state even for V.« (similar exaniples are liquicl water Ireed of dissolved air which may be heated above boiling temperature without the formation of vapor aud cooled below freezing temperature without solidification).," If it is slowly compressed, it may stay in the vapor state even for $V<V_{c}$ (similar examples are liquid water freed of dissolved air which may be heated above boiling temperature without the formation of vapor and cooled below freezing temperature without solidification)."
 The transition [rom (he super-compressed neutron pliase to the quark phase can take place ina very long time. which in some cases can be longer than (he estimated life of the neutron star1998).," The transition from the super-compressed neutron phase to the quark phase can take place in a very long time, which in some cases can be longer than the estimated life of the neutron star."
. Therefore an increase in the density οἱ ihe neutron star does not automatically lead (o the transition to quark matter inside the star., Therefore an increase in the density of the neutron star does not automatically lead to the transition to quark matter inside the star.
 The effect of the magnetic field on the quark star structure and on the nucleation process of the quark bubbles has been studied in 1996).., The effect of the magnetic field on the quark star structure and on the nucleation process of the quark bubbles has been studied in .
 In the presence of strong magnetic fields the equation of state of strange, In the presence of strong magnetic fields the equation of state of strange
sizescales.,sizescales.
 There are objects in these size ranges which have lower column densities and are not sampled by the objects in the catalogue., There are objects in these size ranges which have lower column densities and are not sampled by the objects in the catalogue.
" These plots therefore show the size limit completeness of the catalogue, for both IRDCs, Πο=55"", and fragments, Πο=9""."," These plots therefore show the size limit completeness of the catalogue, for both IRDCs, $R_c=55\arcsec$, and fragments, $R_c=9\arcsec$."
" Adopting these sizes, the average column density of clouds/fragments below these sizes gives average column density completeness limitsof «Ng,>t8=2.5x10?! cm? and <Ng,>1RPC=5x10?! cm7?."," Adopting these sizes, the average column density of clouds/fragments below these sizes gives average column density completeness limitsof $ <N_{H_2}>_{\rm c}^{\rm frag}=
2.5\times10^{21}$ $^{-2}$ and $ <N_{H_2}>_{\rm c}^{\rm IRDC}=5\times10^{21}$ $^{-2}$."
" Using Eq. 1,,"," Using Eq. \ref{eqn:mass},"
 these values give the mass completeness of the catalogue as MIRDC—800 Me and MP8=9 Mo.," these values give the mass completeness of the catalogue as $M_{\rm c}^{\rm
 IRDC}= 800$ $_{\odot}$ and $M_{\rm c}^{\rm frag}=9$ $_{\odot}$."
 The density distribution of IRDCs is difficult to interpret since the clouds are defined based on a column density threshold., The density distribution of IRDCs is difficult to interpret since the clouds are defined based on a column density threshold.
 Therefore we confine our discussion of the density distributions to the fragments., Therefore we confine our discussion of the density distributions to the fragments.
" Both sensitivity and angular resolution are important limiting factors in the context of density completeness of the sample: both compact low-mass fragmentsand large, diffuse, high-mass fragments could remain undetected."," Both sensitivity and angular resolution are important limiting factors in the context of density completeness of the sample: both compact low-mass fragmentsand large, diffuse, high-mass fragments could remain undetected."
" For any fragment mass, Mj there is a minimum radius for which the peak column density of the fragment becomes higher than the threshold for identifying fragments (3x 10?!?)."," For any fragment mass, $M_{\rm lim}$ there is a minimum radius for which the peak column density of the fragment becomes higher than the threshold for identifying fragments $3\times10^{21}$ $^{-2}$ )."
 If this minimum radius is larger than the angularcm resolution then such a fragment is detected., If this minimum radius is larger than the angular resolution then such a fragment is detected.
" Therefore, we can define a density completeness limit for all fragments more massive than Mj."," Therefore, we can define a density completeness limit for all fragments more massive than $M_{\rm lim}$ ."
" The minimum radius, Ryin, and the corresponding maximum density, pl”,are"," The minimum radius, $R_{\rm min}$ , and the corresponding maximum density, $\rho^{\rm low}_{\rm c}$ ,are"
The study of. Protoplanetary ancl Debris Disks is a significant area of research today.,The study of Protoplanetary and Debris Disks is a significant area of research today.
 Lt is uncertain whether the processes which formed our own solar system are typical of other planetary. systems., It is uncertain whether the processes which formed our own solar system are typical of other planetary systems.
 Both protoplanetary and. debris disks describe planetary svstems in the midst of formation/evolution and therefore are crucial to. our understanding of planetary systems as a whole., Both protoplanetary and debris disks describe planetary systems in the midst of formation/evolution and therefore are crucial to our understanding of planetary systems as a whole.
 In steady state. the mutual collisions of planetesimals replenish dust removed from the disk by radiation pressure and. Povnting-Robertson crag.," In steady state, the mutual collisions of planetesimals replenish dust removed from the disk by radiation pressure and Poynting-Robertson drag."
 Planetesimals and regenerated dust around older stars describe debris disks in contrast to the first-generation dust around stars vounger than LO Myr old., Planetesimals and regenerated dust around older stars describe debris disks in contrast to the first-generation dust around stars younger than 10 Myr old.
 Detection of a disk may be gained if infrared emission fron an object is in excess of expectation Irom the stellar fux., Detection of a disk may be gained if infrared emission from an object is in excess of expectation from the stellar flux.
 This is determined numerically by fitting a model spectrum and extrapolating out to the wavelength of observation - sce the Appendix of 7. for details of this. procedure., This is determined numerically by fitting a model spectrum and extrapolating out to the wavelength of observation - see the Appendix of \cite{bryden06} for details of this procedure.
 Tho excess ratio (Pis Pisa) is computed and detection of an infrared. excess is concluded when this value exceeds the 30 limit of the entire sample (2... 7)).," The excess ratio $F_{\text{obs}}/F_{\text{pred}}$ ) is computed and detection of an infrared excess is concluded when this value exceeds the $\sigma$ limit of the entire sample \cite{aumann_discovery_1984}, \cite{beichman_planets_2005}) )."
 In. almost all past and present studies of debris disks. ΑΟ) models (7)) are used for this procedure.," In almost all past and present studies of debris disks, ) models \cite{castelli2004}) ) are used for this procedure."
 Other model groups like (7)) and. (PIOENIX) (7)) remain largely unused and have only recently. become readily. available., Other model groups like \cite{gustaf08}) ) and ) \cite{haus1999}) ) remain largely unused and have only recently become readily available.
 Standard.. 1D model spectra are the result of⋅ a solution. of the equation of radiative transfer. the hydrostatie equation. müxing-length theory and the equations describing the gas phase in chemical equilibrium.," Standard 1D model spectra are the result of a solution of the equation of radiative transfer, the hydrostatic equation, mixing-length theory and the equations describing the gas phase in chemical equilibrium."
 Ideally. every possible atomic and molecular transition should be included however this is computationally impractical and. simplifications are necessary.," Ideally, every possible atomic and molecular transition should be included however this is computationally impractical and simplifications are necessary."
 and model spectra are produced. using Opacity Sampling (7)) where opacities are distributed. statistically in wavelength. giving rise to high resolution model spectra., and model spectra are produced using Opacity Sampling \cite{os}) ) where opacities are distributed statistically in wavelength giving rise to high resolution model spectra.
KURUCZ model spectra however are produced using Opacity Distribution Functions (ODEs) where absorption cross-sections are binned producing a smoother. lower resolution model spectrum. (7... 7).," model spectra however are produced using Opacity Distribution Functions (ODFs) where absorption cross-sections are binned producing a smoother, lower resolution model spectrum \cite{odf1}, \cite{odf2}) )."
 Discrepancies in line features are expected. between mode eroups due to these dillerent opacity treatments though such differences are not. expected at the longer wavelengths a which disks are observed., Discrepancies in line features are expected between model groups due to these different opacity treatments though such differences are not expected at the longer wavelengths at which disks are observed.
 The different model groups include opacities drawn from different literature., The different model groups include opacities drawn from different literature.
 This is likely to introduce sizeable dillerences in the spectral appearance of the model groups: for a discussion. see 7...," This is likely to introduce sizeable differences in the spectral appearance of the model groups; for a discussion, see \cite{jorgensen03}."
 Model spectra with identical stellar parameters but of different mode [anmilies are compared in Section 2 of this paper to determine whether such differences exist., Model spectra with identical stellar parameters but of different model families are compared in Section 2 of this paper to determine whether such differences exist.
 Lhe model spectra often do not extend to longer wavelengths at which observations of disks are made therefore. fitting of a Ravleigh-Jeans tat and extrapolation to the required. wavelength is necessary (Section. 3).," The model spectra often do not extend to longer wavelengths at which observations of disks are made therefore, fitting of a Rayleigh-Jeans tail and extrapolation to the required wavelength is necessary (Section 3)."
 models co extend. bevond 100 ym however. the sparse wavelength distribution at. these wavelengthsB make| it. necessaryUt t0 fit a “eeeBavleigh-Jeans| tail," models do extend beyond $100$ $\mu$ m however, the sparse wavelength distribution at these wavelengths make it necessary to fit a Rayleigh-Jeans tail"
 models co extend. bevond 100 ym however. the sparse wavelength distribution at. these wavelengthsB make| it. necessaryUt t0 fit a “eeeBavleigh-Jeans| tail|," models do extend beyond $100$ $\mu$ m however, the sparse wavelength distribution at these wavelengths make it necessary to fit a Rayleigh-Jeans tail"
halo populations in the solar neighborhood.,halo populations in the solar neighborhood.
 Abundances were determined for a sample of 94 dwarf and subgiant stars with halo or thick-disk kinematics selected from Strómmgren photometry (Schuster et al. 2006)), Abundances were determined for a sample of 94 dwarf and subgiant stars with halo or thick-disk kinematics selected from Strömmgren photometry (Schuster et al. \cite{schuster06}) )
 to have metallicities -1.6<[Fe/H]<—0.4 and to lie in a fairly narrow temperature range. 5200<Tay6300 KK. A differential abundance analysis of high-resolution spectra enabled us to determine with a precision of about ddex.," to have metallicities $-1.6 < \feh < -0.4$ and to lie in a fairly narrow temperature range, $5200 < \teff < 6300$ K. A differential abundance analysis of high-resolution spectra enabled us to determine with a precision of about dex."
 As seen from Fig., As seen from Fig.
" 1 of NSIO. the halo stars are distributed into two groups clearly separated in ffor the metallicity range —1.4<[Fe/H]-0.7. re. “high-a° stars having eclose to 0.30 dex. similar to thick-disk stars. and ""owes stars with ddecreasing from ddex at [Fe/H]=—1.4 to 0.10ddex at |Fe/H]=-0.7."," 1 of NS10, the halo stars are distributed into two groups clearly separated in for the metallicity range $-1.4 < \feh < -0.7$, i.e. $\alpha$ ' stars having close to 0.30 dex, similar to thick-disk stars, and $\alpha$' stars with decreasing from dex at $\feh = -1.4$ to dex at $\feh = -0.7$."
 Differences in aand aare also present with a remarkably tight correlation between these two ratios.[," Differences in and are also present with a remarkably tight correlation between these two ratios.,"
Cr/Fe]. on the other hand. is the same for the two populations.," on the other hand, is the same for the two populations."
 The halo stars in the NSIO sample were selected to have Viu>180., The halo stars in the NS10 sample were selected to have $\Vtotal > 180$.
 Yet. the high- and low-a populations have different kinematical properties as seen from Fig.," Yet, the high- and $\alpha$ populations have different kinematical properties as seen from Fig."
 3 in NS10., 3 in NS10.
 The high-c stars tend to move on prograde Galactic orbits. whereas two- of the low-a stars have retrograde orbits with an average Galactic rotational velocity component. Vjsj=—260kmprs.. close to that of the «o CCen globular cluster (Dinescu et al. 1990).," The $\alpha$ stars tend to move on prograde Galactic orbits, whereas two-thirds of the $\alpha$ stars have retrograde orbits with an average Galactic rotational velocity component, $V_{\rm LSR} \simeq -260$, close to that of the $\omega$ Cen globular cluster (Dinescu et al. \cite{dinescu99}) )."
 Furthermore. the low-a stars have a very wide distribution of the radial velocity component. ÜU. but comparatively small values of the W component.," Furthermore, the $\alpha$ stars have a very wide distribution of the radial velocity component, $U$, but comparatively small values of the $W$ component."
 All together. this suggests that the low-a stars have been accreted from satellite galaxies including the progenitor galaxy of ω CCen (Bekki Freeman 2003)).," All together, this suggests that the $\alpha$ stars have been accreted from satellite galaxies including the progenitor galaxy of $\omega$ Cen (Bekki Freeman \cite{bekki03}) )."
" The high-c stars. on the other hand. are more likely to be formed in situ’ in the Galactic bulge or disk. and then ""heated? to halo kinematics by the merging satellites. as suggested from numerical simulations (Zolotov et al. 2009.. 2010::"," The $\alpha$ stars, on the other hand, are more likely to be formed `in situ' in the Galactic bulge or disk, and then `heated' to halo kinematics by the merging satellites, as suggested from numerical simulations (Zolotov et al. \cite{zolotov09}, \cite{zolotov10};"
 Purcell et al. 20101:, Purcell et al. \cite{purcell10};
 Qu et al. 201100., Qu et al. \cite{qu11}) ).
 In order to obtain new insight into this scenario and the possible connection between ω (ει and some of the a stars. our abundance analysis has been extended to include Mn. Cu. Zn. Y. and Ba.," In order to obtain new insight into this scenario and the possible connection between $\omega$ Cen and some of the $\alpha$ stars, our abundance analysis has been extended to include Mn, Cu, Zn, Y, and Ba."
 In the following Sect. 2..," In the following Sect. \ref{sect:abundances},"
 we describe how the abundances of these elements are determined and evaluate statistical and systematic errors., we describe how the abundances of these elements are determined and evaluate statistical and systematic errors.
 In Sect. 3.," In Sect. \ref{sect:discussion},"
 1t is discussed if the metallicity trends and differences in abundance ratios between the two halo populations can be explained from current knowledge of the nucleosynthesis of elements in stars and supernovae., it is discussed if the metallicity trends and differences in abundance ratios between the two halo populations can be explained from current knowledge of the nucleosynthesis of elements in stars and supernovae.
 Furthermore. the abundance ratios of the α stars are compared with the corresponding ratios in present-day Milky Way satellite galaxies and the wCCen globular cluster.," Furthermore, the abundance ratios of the $\alpha$ stars are compared with the corresponding ratios in present-day Milky Way satellite galaxies and the $\omega$ Cen globular cluster."
 Finally. some conclusions are given in Sect. 4..," Finally, some conclusions are given in Sect. \ref{sect:conclusions}."
 Elemental abundances are derived from a model-atmosphere analysis of equivalent widths (EWs) measured in high signal-to-noise (S/N). high-resolution spectra observed either with the ESO/VLT UVES spectrograph or the FIES spectrograph at the Nordic Optical Telescope (NOT).," Elemental abundances are derived from a model-atmosphere analysis of equivalent widths (EWs) measured in high signal-to-noise (S/N), high-resolution spectra observed either with the ESO/VLT UVES spectrograph or the FIES spectrograph at the Nordic Optical Telescope (NOT)."
 Details on wavelength range. resolution. and S/N are given in NSIO.," Details on wavelength range, resolution, and S/N are given in NS10."
 Figure | shows a comparison of the spectra of a high- and a low-« star in spectral bands containing some of the Mn. Cu. Zn. Υ. and Ba lines applied in this paper.," Figure \ref{fig:all-lines} shows a comparison of the spectra of a high- and a $\alpha$ star in spectral bands containing some of the Mn, Cu, Zn, Y, and Ba lines applied in this paper."
 As deseribed in NSIO. local thermodynamic equilibrium (LTE) is assumed and abundance ratios are determined relative to two bright thick-disk stars. and76932.," As described in NS10, local thermodynamic equilibrium (LTE) is assumed and abundance ratios are determined relative to two bright thick-disk stars, and."
. They are nearby stars for which ccan be determined from photometric colors. and vvia Hipparcos parallaxes.," They are nearby stars for which can be determined from photometric colors, and via Hipparcos parallaxes."
 Using a subset of the spectral lines for which the equivalent widths could be measured reliably i the solar flux spectrum (Kuruez et al. 1984)).," Using a subset of the spectral lines for which the equivalent widths could be measured reliably in the solar flux spectrum (Kurucz et al. \cite{kurucz84}) ),"
 an abundance analysis relative to the Sun is carried out., an abundance analysis relative to the Sun is carried out.
 Adopting these abundances. an ‘inverted’ analysis leads to the determinatior of e f-values for all lines.," Adopting these abundances, an `inverted' analysis leads to the determination of $gf$ -values for all lines."
 The derived gf-values for ane agree within £0.03 ddex for the large majority of lines: for each line. the mean value of ef ts therefore adopted and used for the analysis of all program stars.," The derived $gf$ -values for and agree within $\pm
0.03$ dex for the large majority of lines; for each line, the mean value of $gf$ is therefore adopted and used for the analysis of all program stars."
 This procedure ensures a high internal precision of the abundance ratios., This procedure ensures a high internal precision of the abundance ratios.
 Relative to the Sun. the abundances are. however. more uncertain; line blending and continuum setting are often problematic in the solar spectrum. and non-LTE effects may be different for the Sun and a typical program star with [Fe/H]~—1.," Relative to the Sun, the abundances are, however, more uncertain; line blending and continuum setting are often problematic in the solar spectrum, and non-LTE effects may be different for the Sun and a typical program star with $\feh \sim -1$."
 Thus. we find the average difference in Fe abundance derived from aand llines to be |Fe/H|j-|[Fe/H]|;=0.075 ddex for the two standard stars.," Thus, we find the average difference in Fe abundance derived from and lines to be $\fehII - \fehI = 0.075$ dex for the two standard stars."
 This difference is probably caused by an over-ionization of wwith respect to LTE in the metal-poor stars (Asplund 2005))., This difference is probably caused by an over-ionization of with respect to LTE in the metal-poor stars (Asplund \cite{asplund05}) ).
 iis adopted as the metallicity parameter. but abundance ratios are based on lines from the same ionization stage. e.g. iis derived from aand llines.," is adopted as the metallicity parameter, but abundance ratios are based on lines from the same ionization stage, e.g. is derived from and lines."
 Hence. if the non-LTE effects are the same for the two sets of lines. iis on the correct scale. but as discussed in Sect.," Hence, if the non-LTE effects are the same for the two sets of lines, is on the correct scale, but as discussed in Sect."
 3.2.2. this may not the case., \ref{sect:cr-mn} this may not the case.
 Aetallicity dependent departures from LTE could also affect the derived trends of abundance ratios as a function of|Fe/H]., Metallicity dependent departures from LTE could also affect the derived trends of abundance ratios as a function of.
 However. when comparing stars with the same metallicity and within narrow ranges of aand loge. non-LTE effects are expected to be the same: the LTE approximation ts therefore adequate in providing reliable differential abundances at a given metallicity.," However, when comparing stars with the same metallicity and within narrow ranges of and , non-LTE effects are expected to be the same; the LTE approximation is therefore adequate in providing reliable differential abundances at a given metallicity."
 For many program stars. colors are affected by interstellar reddening and the parallaxes are not accurate enough to determine reliable surface gravities.," For many program stars, colors are affected by interstellar reddening and the parallaxes are not accurate enough to determine reliable surface gravities."
 Thus.," Thus,"
period!).,period!).
 Llere we shall discuss the superorbital period. along with reasons why it might vary with time.," Here we shall discuss the superorbital period, along with reasons why it might vary with time."
 Four mechanisms have been put forward to explain the superorbital period. and. we shall discuss these for the case of SAIC X-1I.," Four mechanisms have been put forward to explain the superorbital period, and we shall discuss these for the case of SMC X-1."
 Broadly speaking. mechanisms for generating a superorbital »eriod. must. do so by a variation in the X-ray intensity of 10 source. by variation in uncovered. area as seen from the bserver. or by a varying absorption along the line of sight to 16 observer.," Broadly speaking, mechanisms for generating a superorbital period must do so by a variation in the X-ray intensity of the source, by variation in uncovered area as seen from the observer, or by a varying absorption along the line of sight to the observer."
 However we can immediately rule out. varving bsorption., However we can immediately rule out varying absorption.
 This is because the BATSE lighteurve varies in 1e same way as that of the ASAL. showing that the variation occurs across the entire energy range 1.3 - 100 keV. This rules out absorption by wind or accretion disk.," This is because the BATSE lightcurve varies in the same way as that of the ASM, showing that the variation occurs across the entire energy range 1.3 - 100 keV. This rules out absorption by wind or accretion disk."
 For example. although the wind density in SAIC X-1 can vary by factors of up to ~1000 (Bloncin and. Woo 1995. Wojdowski et a 100. see also Vetilek et al.," For example, although the wind density in SMC X-1 can vary by factors of up to $\sim$ 1000 (Blondin and Woo 1995, Wojdowski et al 2000, see also Vrtilek et al.,"
 2001). this would have no elfec on the BATSE lighteurve.," 2001), this would have no effect on the BATSE lightcurve."
 This also rules out in the case of SAIC X-1: the conventional explanation for superorbita variations. in which varving absorption by a clisk warp causes the pertoclicity.," This also rules out in the case of SMC X-1 the conventional explanation for superorbital variations, in which varying absorption by a disk warp causes the periodicity."
 In itself. the relation between ASAI intensity anc spectral hardness does not act as à good discriminan between mechanisms. as the lighteurve represents a multiple component. source (sce Tanaka LOOT for a review).," In itself, the relation between ASM intensity and spectral hardness does not act as a good discriminant between mechanisms, as the lightcurve represents a multiple component source (see Tanaka 1997 for a review)."
 In the case of neutron star NR. the soft. component (usually identified with energies κο 3 keV) is taken to originate in the accretion disk. through multi-temperature disk blackbocdvy emission (Mitsuda et al 1984) combined with a Comptonised component (White. Stella and Parmar 1988).," In the case of neutron star XRB, the soft component (usually identified with energies $<$ 3 keV) is taken to originate in the accretion disk, through multi-temperature disk blackbody emission (Mitsuda et al 1984) combined with a Comptonised component (White, Stella and Parmar 1988)."
 The wind is also represented at à low level in the soft component of the emission from SAIC X-1 (Wojdowskietal2000)., The wind is also represented at a low level in the soft component of the emission from SMC X-1 \cite{woj2k}.
. Llowever it should be noted that a small component of the soft emission may in some cases arise at the surface of the neutron star (Schulz1905)., However it should be noted that a small component of the soft emission may in some cases arise at the surface of the neutron star \cite{schul}.
 Emission at harder energies is identified. with the surface region of the neutron star itself. through shock or Coulomb heating or Compton scattering at the boundary laver (Frank.Wine&Raine1995).," Emission at harder energies is identified with the surface region of the neutron star itself, through shock or Coulomb heating or Compton scattering at the boundary layer \cite{fkr}."
. lt has recently been suggested that the similarity in A-rayv spectra. between low state neutron star svstems ancl black hole candidates identifies the dominant source of the hard component as the advection dominated. accretion. How (ADA) between the inner disk and compact object (Barretetal.2000)., It has recently been suggested that the similarity in X-ray spectra between low state neutron star systems and black hole candidates identifies the dominant source of the hard component as the advection dominated accretion flow (ADAF) between the inner disk and compact object \cite{bar}.
. Thus the majority of the soft component. of the ASAI lighteurve can be identified with the acerction disk or matter ejected from the companion. whereas the hard component corresponds to emission. at or. near the surface of the neutron star itself," Thus the majority of the soft component of the ASM lightcurve can be identified with the accretion disk or matter ejected from the companion, whereas the hard component corresponds to emission at or near the surface of the neutron star itself."
 In this way. the spectral dependence of the ASM lighteurve allows us to determine the region in which variation is taking place.," In this way, the spectral dependence of the ASM lightcurve allows us to determine the region in which variation is taking place."
 The superorbital variation is stronglv energy dependent: as can be seen from figures 4.. 5 6. the ASM variation is most pronounced at. energies - 5keV. while there is little or no variation at energies « 3 keV. This suggests we can identify the observed superorbital variation with changes at or near the surface of the neutron star.," The superorbital variation is strongly energy dependent: as can be seen from figures \ref{fig:fold_hrat}, \ref{fig:hrat_counts} \ref{fig:fold_channs}, , the ASM variation is most pronounced at energies $>$ 5keV, while there is little or no variation at energies $<$ 3 keV. This suggests we can identify the observed superorbital variation with changes at or near the surface of the neutron star."
 This interpretation is only valid if ejected material does not dominate the soft component of the RATE lightcurve., This interpretation is only valid if ejected material does not dominate the soft component of the RXTE lightcurve.
 ASCA spectra of the X-ray pulsar Con N-3 (Ebisawaeal.1996) do indeed contain little or no variation a energies «2 keV in eclipse. whereas the hard. componen varies significantly.," ASCA spectra of the X-ray pulsar Cen X-3 \cite{ebi} do indeed contain little or no variation at energies $<$ 2 keV in eclipse, whereas the hard component varies significantly."
 The authors model the spectrum with three components. (figure. 6 of Ehisawa ct al 1996): a hard component from the neutron star itself. significant a energies >Λο. à component due to electron scattering in cireumstellar matter. dominant at energies 1 - 5 keV. ane a soft component due to scattering from interstellar dust. dominant at energies <2 keV. by virtue of the location of Cen X-3 in the ealactic plane.," The authors model the spectrum with three components (figure 6 of Ebisawa et al 1996): a hard component from the neutron star itself, significant at energies $>3 keV$, a component due to electron scattering in circumstellar matter, dominant at energies 1 - 5 keV, and a soft component due to scattering from interstellar dust, dominant at energies $<2$ keV, by virtue of the location of Cen X-3 in the galactic plane."
 In SAIC X- however. the circtuustcllar componen cannot be an important contributor to the ASAI lightcurve.," In SMC X-1, however, the circumstellar component cannot be an important contributor to the ASM lightcurve."
 In the case of Cen X-3. the pulse period is strongly energy dependent. with little or no evidence for the pulsation at energies corresponding (o the circumstellar electron scattering or interstellar dust components. (I5bisawaeal. 1996).. which suggests circumstellar ancl interstellar scattering matter are responsible for the emission. of the bulk of the soft. N-ravs.," In the case of Cen X-3, the pulse period is strongly energy dependent, with little or no evidence for the pulsation at energies corresponding to the circumstellar electron scattering or interstellar dust components \cite{ebi}, which suggests circumstellar and interstellar scattering matter are responsible for the emission of the bulk of the soft X-rays."
 Phe position of SMC. N-1 out of the galactic plane rules out any significant interstellar clus component. and its pulse profile is clear and present. righ down to QO.1keV. (Wojdowskietal199s).," The position of SMC X-1 out of the galactic plane rules out any significant interstellar dust component, and its pulse profile is clear and present right down to 0.1keV \cite{woj}."
. This rules ou domination of the soft. X-ray. lighteurve from circumstellar or interstellar scattering regions., This rules out domination of the soft X-ray lightcurve from circumstellar or interstellar scattering regions.
 Thus the absence of superorbital variation in the soft energy. bands of the ASA lighteurve is truly showing that radiation from the disk is not an important participant in the variation., Thus the absence of superorbital variation in the soft energy bands of the ASM lightcurve is truly showing that radiation from the disk is not an important participant in the variation.
 Ppecession of the neutron stars magnetic axis was invoked o explain the ~35d periodicity in Her X-1 (bv changing the accretion geometry: see Trümampoer ct al., Precession of the neutron star's magnetic axis was invoked to explain the $\sim$ 35d periodicity in Her X-1 (by changing the accretion geometry; see Trümmper et al.
 1986)., 1986).
 However his mechanism can be ruled out in the case of SAIC X-1 or two reasons., However this mechanism can be ruled out in the case of SMC X-1 for two reasons.
 Firstly. the precession torque is expected o be of the same order as the spin-up torque. (Lambetal. 1975).," Firstly, the precession torque is expected to be of the same order as the spin-up torque \cite{lam}."
. jut pointed N-ray observations of SAIC X- show the spin period to be monotonically decreasing (Wojdowskietal1998).. suggesting that precession should either also monotonically decrease in period. or if the sign of he torque changes. produce more violent behaviour in the dvnamie power spectrum over time.," But pointed X-ray observations of SMC X-1 show the spin period to be monotonically decreasing \cite{woj}, suggesting that precession should either also monotonically decrease in period, or if the sign of the torque changes, produce more violent behaviour in the dynamic power spectrum over time."
 This is not observed., This is not observed.
 Secondly. the precession of the neutron star would result in a superorbital variation of the pulse profile. whereas pointec observations suggest that the pulse profile is in fact. stable (Wojdowskietal1998).," Secondly, the precession of the neutron star would result in a superorbital variation of the pulse profile, whereas pointed observations suggest that the pulse profile is in fact stable \cite{woj}."
. Forced precession has recently been ruled out in the case of Ller X-1 for exactly this reason (Scot 2000).., Forced precession has recently been ruled out in the case of Her X-1 for exactly this reason \cite{scott}.
 Another possibility is variation in mass transfer [rom donorto neutron star on the superorbital period., Another possibility is variation in mass transfer from donorto neutron star on the superorbital period.
 Cian pulsators. most notably RRO Lyrae variables. are. known to change their pulsation periods (see. e.g: Jurscik 2001).," Giant pulsators, most notably RR Lyrae variables, are known to change their pulsation periods (see, e.g: Jurscik 2001)."
window function. let us calculate 2717(13)Ck.]).,"window function, let us calculate $2P_A^{(13)}(k,\eta)$."
 where dy is the mitial fluctuation., where $\delta_0$ is the initial fluctuation.
 This can be further siuplified by performing the augular part of the iutegration., This can be further simplified by performing the angular part of the integration.
 Considering oulv the fastest-growing mode aud assuming spherical top-hat smoothing (sce AppendixA aud for more details). 193) 1)ÓCOLUCN where the variauce of the smoothed deusitv fluctuation oF;=|PaPrly.t)WSoum. and the derivative Tere we have used the geometric properties of the spherical top-hat window function derived by 1996)..," Considering only the fastest-growing mode and assuming spherical top-hat smoothing (see \ref{sec:1loop_spt} and \ref{sec:tophat}
 for more details), \ref{eqn:SPT_13}) ) becomes where the variance of the smoothed density fluctuation $\sigma^2_R = \int d^3\vq ~ P_L(q,\eta) ~ W^2_{\rm th}(qR),$ and the derivative Here we have used the geometric properties of the spherical top-hat window function derived by \cite{B94a,B94,B96}. ."
" Simce ""ds proportional to the lincar power spectrum Py. aud it is well-known that (pe|ΓΠΣ»0)>0. the power spectrump at large scales is biased with respect to the density power spect Without the soothius terii. the bias m ((22)) wo14 equal 1262195."," Since $P^{(13)}_A$ is proportional to the linear power spectrum $P_L$, and it is well-known that $(P^{(13)}_{\delta}/P_L) (k\to 0) \to 0$, the power spectrum at large scales is biased with respect to the density power spectrum Without the smoothing term, the bias in \ref{eqn:SPT_bias}) ) would equal $1-26/21~\sigma_R^2$."
 This expression accords with the first-order perturbative bias in ((À2) (after expanding he cuunulants) of Nevrincketal.(2000111 it is also plotted asa fiction of vedshitt iu LL of that paper., This expression accords with the first-order perturbative bias in (A2) (after expanding the cumulants) of \citet{NSS09}; it is also plotted as a function of redshift in 4 of that paper.
 The bias becomes unplivsically negative at ;X[E for ecells., The bias becomes unphysically negative at $z\lesssim 4$ for cells.
 This smoothing term increases the predicted bias. extending its range of validitv to significantly smaller smoothing scales.," This smoothing term increases the predicted bias, extending its range of validity to significantly smaller smoothing scales."
 Towever. we found that it still goes pathologically negative at small ioothiug scales.," However, we found that it still goes pathologically negative at small smoothing scales."
 This iue is resolved in a renormalized approach. which we describe in the next section.," This issue is resolved in a renormalized approach, which we describe in the next section."
 As the density fluctuations evolve into the non-linear reein1ο at late times. the validity of standard perturbation theory breaks down.," As the density fluctuations evolve into the non-linear regime at late times, the validity of standard perturbation theory breaks down."
 Loop contributions become ill-bebhaved aud he convergence of perturbation series ects out of control., Loop contributions become ill-behaved and the convergence of perturbation series gets out of control.
 Several different approaches bevoud havebeen proposed receutlhv. tthe renormalized perturbation theory (RPT) (Crocce&Scoccinirro2006a.b:TIDornardeauctal. 2008).. the Lagrangian resummation theory 2008).. the closure theory (Tariuva&Llivaunaten2008).. and the time renormalization eroup theory (TRG) 2008).," Several different approaches beyond SPT have been proposed recently, the renormalized perturbation theory (RPT) \citep{CS06a,CS06b,BCS08}, , the Lagrangian resummation theory \citep{M08}, the closure theory \citep{TH08}, and the time renormalization group theory (TRG) \citep{P08}."
. Iu this section. after a brief review of the RPT introduced by Crocce&Scoccimarro(20062.b).. we will show that it is possible to construct the perturbation series for the transformed density field based ou the non-linear propagator of the density.," In this section, after a brief review of the RPT introduced by \cite{CS06a,CS06b}, we will show that it is possible to construct the perturbation series for the transformed density field based on the non-linear propagator of the density."
 The crucial step of renormalized perturbation theory is to define the generalized growth. factor. known as the propagator κ) which effectively describes the time evolution of individual Fourier modes when uou-linear mode-coupling is iucluded.," The crucial step of renormalized perturbation theory is to define the generalized growth factor, known as the propagator $G_{ab}(k)$ which effectively describes the time evolution of individual Fourier modes when non-linear mode-coupling is included."
" The propagator measures the depeudence of a non-linearly evolved Fourier mode V,(Kk.5j) ou its initial state o,(k) ou average."," The propagator measures the dependence of a non-linearly evolved Fourier mode $\Psi_a(\vk, \eta)$ on its initial state $\phi_b(\vk)$ on average."
" Intuitively. one should expect Gy, to decay to zero at πα scales since non-uear mode-coupling las erased all the information from the initial state at that k."," Intuitively, one should expect $G_{ab}$ to decay to zero at small scales since non-linear mode-coupling has erased all the information from the initial state at that $\vk$."
" Iudeed. with the help of the Feviuman diagrams introduced in (2006a).. the dominant contribution cau be sumuned up explicitly in the larec-k lit. and eivius the Caussian decay Gy,(k)=exptM?o2(i1)? /2)."," Indeed, with the help of the Feynman diagrams introduced in \cite{CS06a}, the dominant contribution can be summed up explicitly in the large-k limit, and giving the Gaussian decay $G_{ab}(k) = \exp(-k^2 \sigma_v^2 (a-1)^2/2)$ ."
 Ta this paper. the non-linearpropagator is diagraunaatically presented asa erey circle with au incoming brauch.," In this paper, the non-linearpropagator is diagrammatically presented as a grey circle with an incoming branch."
 It is a πα1ο of infinite umberof loop contributions. as illustrated iu," It is a summation of infinite numberof loop contributions, as illustrated in\\ref{fig:P_psi}. ."
scattering halo surrounding the black hole and inner disk region.,scattering halo surrounding the black hole and inner disk region.
 We are not necessarily espousing the view that this is a viable model for the NIR/X-ray flare characteristics observed in recent years., We are not necessarily espousing the view that this is a viable model for the NIR/X-ray flare characteristics observed in recent years.
" If Comptonization is important in producing the X-rays, the incipient emission need not be that due to a synchrotron-cooled particle distribution."," If Comptonization is important in producing the X-rays, the incipient emission need not be that due to a synchrotron-cooled particle distribution."
 Our main goal here is to demonstrate how the images and polarization fraction (and possibly the position angle) are altered by scattering from those corresponding to thedirect images., Our main goal here is to demonstrate how the images and polarization fraction (and possibly the position angle) are altered by scattering from those corresponding to the images.
" Scattering clearly broadens the image, by a degree dependent on the density n,. of scattering particles."," Scattering clearly broadens the image, by a degree dependent on the density $n_{sc}$ of scattering particles."
" In this simulation, we have chosen a value n,.=1011 cm, corresponding to a photon"," In this simulation, we have chosen a value $n_{sc}=10^{11}$ $^{-3}$, corresponding to a photon"
Nearby edge-on galaxies provide an important local laboratory for understanding physical conditions in. the. disk-halo interface.,Nearby edge-on galaxies provide an important local laboratory for understanding physical conditions in the disk-halo interface.
 This critical region separates two very different environments over a relatively small vertical distance., This critical region separates two very different environments over a relatively small vertical distance.
" It plays an important role in a galaxy’s energy balance as well as its chemical evolution. since outflows from the disk. circulating ""fountains'. and possibly infalling external material may each be involved at some point in a galaxy's evolutio"," It plays an important role in a galaxy's energy balance as well as its chemical evolution, since outflows from the disk, circulating `fountains', and possibly infalling external material may each be involved at some point in a galaxy's evolution."
" Moreover. knowledge of the physical conditions of outflowing gas m nearby galaxies can provide important constraints oαυ] galaxy formation scenarios. since outflows (1e. some form of ""feedback) are crucial to these models (e.g. Marr White. 2003 and others)."," Moreover, knowledge of the physical conditions of outflowing gas in nearby galaxies can provide important constraints on galaxy formation scenarios, since outflows (i.e. some form of `feedback') are crucial to these models (e.g. Marri White, 2003 and others)."
 Nearby edge-on galaxies have now been studied in many wavebands showing a variety of interstellar medium (ISM) components above the galactic disk., Nearby edge-on galaxies have now been studied in many wavebands showing a variety of interstellar medium (ISM) components above the galactic disk.
 Indeed. in at least one galaxy that has a high star formation rate (SFR). every component of the interstellar medium. including dust. molecular gas. HI. extra-planar diffuse ionized gas (eDIG. detected by H« emission) and X-ray emitting hot gas has been detected above the galactic disk (see Lee et al.," Indeed, in at least one galaxy that has a high star formation rate (SFR), every component of the interstellar medium, including dust, molecular gas, HI, extra-planar diffuse ionized gas (eDIG, detected by $\alpha$ emission) and X-ray emitting hot gas has been detected above the galactic disk (see Lee et al."
 2001 and Brar et al., 2001 and Brar et al.
 2003)., 2003).
 In each component. discrete disk-halo vertical structures are seen. suggestive of outflow.," In each component, discrete disk-halo vertical structures are seen, suggestive of outflow."
 The ∣⋪⊖∁⊜∏↾↳∐⋋∁⋯⇁⊖∣⋪⋝⇁⋅↑↴⋯⋯∏⋯↿−≣∏↑↴∣⋪⋯⋪⊜∐ (MIR) observations. of PAH (polyeyelic aromatic hydrocarbon) in the halo of the low SFR galaxy. NGC 5907 Urwin Madden 2006). has prompted a search for other galaxies that may show similar emission.," The recent discovery, from mid-infrared (MIR) observations, of PAH (polycyclic aromatic hydrocarbon) in the halo of the low SFR galaxy, NGC 5907 (Irwin Madden 2006), has prompted a search for other galaxies that may show similar emission."
" In this paper. we present results for the edge-on galaxy. NGC 5529. using archival data from the Infrared Space Observatory (ISO) in the 16.75 uim waveband. which selects the ,16.2. 417.7. and part of the 418.6 ym PAH features (see Fig."," In this paper, we present results for the edge-on galaxy, NGC 5529, using archival data from the Infrared Space Observatory (ISO) in the $\lambda\,6.75$ $\mu$ m waveband, which selects the $\lambda\,6.2$, $\lambda\,7.7$ , and part of the $\lambda\,8.6$ $\mu$ m PAH features (see Fig."
 7 of Irwin Madden 2006)., 7 of Irwin Madden 2006).
 These were the only ISO observations taken of this galaxy., These were the only ISO observations taken of this galaxy.
 In Sect. 2..," In Sect. \ref{ngc5529},"
 we introduce the galaxy. Sect.," we introduce the galaxy, Sect."
 3 presents the observations and data reduction. Sect.," \ref{observations_reductions} presents the observations and data reduction, Sect."
 + lists the results. and the discussion and conclusion are presented in Sects.," \ref{results} lists the results, and the discussion and conclusion are presented in Sects."
 5 and 6.. respectively.," \ref{discussion} and \ref{conclusions}, respectively."
 NGC 5529. whose basic parameters are listed in Table 5.. is an edge-on galaxy with a prominent dust lane.," NGC 5529, whose basic parameters are listed in Table \ref{basic_parameters}, is an edge-on galaxy with a prominent dust lane."
 It is classified as an Se galaxy. but is likely barred (Kregel van der Kruit 2004) and a peanut-shaped bulge ts also apparent in the optical image shown in Fig. l..," It is classified as an Sc galaxy, but is likely barred (Kregel van der Kruit 2004) and a peanut-shaped bulge is also apparent in the optical image shown in Fig. \ref{n5529_colour}."
 At a distance of D=43.9 Mpe (1 aresec = 213 pe). this galaxy is physically very large (diameter. D=81 kpe).," At a distance of $D\,=\,43.9$ Mpc (1 arcsec = 213 pc), this galaxy is physically very large (diameter, $D\,=\,81$ kpc)."
 However. its properties related to star formation (SF) are similar to those of NGC 5907 (see Table 5)) which is half its size.," However, its properties related to star formation (SF) are similar to those of NGC 5907 (see Table \ref{basic_parameters}) ) which is half its size."
 The SFRs of both galaxies are modest. re. 1 to 3 M. yr: by comparison. the SFR of the better-known quiescent galaxy. NGC 891. is 3.8 Ma vr! (Popescu et al.," The SFRs of both galaxies are modest, i.e. 1 to 3 $_\odot$ $^{-1}$; by comparison, the SFR of the better-known quiescent galaxy, NGC 891, is 3.8 $_\odot$ $^{-1}$ (Popescu et al."
 2000)., 2000).
 With an inclination of ¢=90° and an SER that is similar to a galaxy in which PAH halo emission has already been detected. GC 5529 ts an ideal candidate for the study of extraplana emission i general. and the search for high latitude PAHs 1 particular.," With an inclination of $i\,=\,90^\circ$ and an SFR that is similar to a galaxy in which PAH halo emission has already been detected, NGC 5529 is an ideal candidate for the study of extraplanar emission in general, and the search for high latitude PAHs in particular."
 To our knowledge. only one recent search for extraplana emission in NGC 5529 has been carried out.," To our knowledge, only one recent search for extraplanar emission in NGC 5529 has been carried out."
 This involved a Ho/N|IT observation (hereafter referred to as Ha) by Miller Veilleux (2003). with a positive detection.," This involved an $\alpha$ /N[II] observation (hereafter referred to as $\,\alpha$ ) by Miller Veilleux (2003), with a positive detection."
" These authors report widespreac eDIG within the central 10 kpe of the disk anc find a weak. extended component with an average exponential scale height of z, = 4.5 kpe."," These authors report widespread eDIG within the central 10 kpc of the disk and find a weak, extended component with an average exponential scale height of $z_e$ = 4.5 kpc."
 The total eDIG mass. Mepjc;- was found to be Mp;=2.0x10° Μ..," The total eDIG mass, $_{eDIG}$, was found to be $_{eDIG}\,=\,2.0\times\,10^8$ $_\odot$."
 By contrast. ai early radio continuum observation did not detect extraplanar radio emission. but the observations were likely not sensitive enough to do so (Hummel et al.," By contrast, an early radio continuum observation did not detect extraplanar radio emission, but the observations were likely not sensitive enough to do so (Hummel et al."
 1991)., 1991).
 In the soft X-ray band. an upper limit of Ly=2.2x101 erg s! (scaled to our adopted distance. and so throughout) hàs been placed on the galaxy's total X-ray emission (Benson et al.," In the soft X-ray band, an upper limit of $L_X\,=\,2.2\times\,10^{41}$ erg $^{-1}$ (scaled to our adopted distance, and so throughout) has been placed on the galaxy's total X-ray emission (Benson et al."
 2000)., 2000).
 The stellar component of GC 5529 has been widely studied. however. and various authors have reported a thick stellar disk. although optical neasurements have been hindered. to some extent. by the prominent dust lane (Fig. 1)).," The stellar component of NGC 5529 has been widely studied, however, and various authors have reported a thick stellar disk, although optical measurements have been hindered, to some extent, by the prominent dust lane (Fig. \ref{n5529_colour}) )."
" For example. vertical fits to B. V. and I-band data give z, = 0.99. 0.92. and 0.86 kpe. respectively. where the best fit is to a sec/|z|/z.) function (de Grijs van der Kruit 1996)."," For example, vertical fits to B, V, and I-band data give $z_{s}$ = 0.99, 0.92, and 0.86 kpc, respectively, where the best fit is to a $sech(|z|/z_s)$ function (de Grijs van der Kruit 1996)."
 More recently. exponential fits to R-banddata result in à mean value of z;=8” (1.7 kpe). increasing with galactocentric distance (Schwarzkopf Dettmar 2001). an effect that may be related to the fact that," More recently, exponential fits to R-banddata result in a mean value of $z_e\,=\,8^{\prime\prime}$ (1.7 kpc), increasing with galactocentric distance (Schwarzkopf Dettmar 2001), an effect that may be related to the fact that"
is observed.,is observed.
" However, the data are useful to check that none of the extended emission found at longer wavelengths is due to a such as a low redshift galaxy."," However, the data are useful to check that none of the extended emission found at longer wavelengths is due to a contaminating object such as a low redshift galaxy."
 At z’ band the object contaminatingis objectclearly more resolved than stars in the field., At $z'$ band the object is clearly more resolved than stars in the field.
" G09 carried out point-spread function (PSF) subtraction revealing an elongated (roughly E-W) object extended over 4 arcsec with magnitude z’=23.5c0.3 in a 2.6"" radius aperture."," G09 carried out point-spread function (PSF) subtraction revealing an elongated (roughly E-W) object extended over 4 arcsec with magnitude $z'=23.5
\pm 0.3$ in a 2.6” radius aperture."
" Due to the lower sensitivity of Suprime-Cam at z, band, G09 found only tentative detection of extended flux in this filter (which probesa longward of the eemission line) with magnitude z,>24."," Due to the lower sensitivity of Suprime-Cam at $z_r$ band, G09 found only a tentative detection of extended flux in this filter (which probes longward of the emission line) with magnitude $z_r>24$."
 The z’ band contains the Iline and the brightness of the extended emission in this filter suggests that at least some fraction of the extended emission is fromα., The $z'$ band contains the line and the brightness of the extended emission in this filter suggests that at least some fraction of the extended emission is from.
". In order to compare these imaging observations with our spectroscopy to be presented in the following section, we have obtained the Suprime-Cam z/ band image from the authors of (09."," In order to compare these imaging observations with our spectroscopy to be presented in the following section, we have obtained the processed Suprime-Cam $z'$ band image from the authors of G09."
 processedWe performed an independent PSF subtraction using field stars of similar magnitude to the quasar., We performed an independent PSF subtraction using field stars of similar magnitude to the quasar.
" Six stars were used to construct a model PSF which has FWHM 2.4 or 0.48"".", Six stars were used to construct a model PSF which has FWHM 2.4 pixels or 0.48”.
" The PSF was scaled to match the total magnitude pixelsof the quasar in a circular aperture of radius 2"".", The PSF was scaled to match the total magnitude of the quasar in a circular aperture of radius 2”.
 This PSF was then subtracted from the quasar image., This PSF was then subtracted from the quasar image.
 By trial and error it was found that the model PSF should be scaled by a factor of 0.87 such that the integrated flux within the central few pixels was just above zero., By trial and error it was found that the model PSF should be scaled by a factor of 0.87 such that the integrated flux within the central few pixels was just above zero.
 Note that this method (maximal PSF-subtraction) therefore gives a lower limit on the extended flux., Note that this method (maximal PSF-subtraction) therefore gives a lower limit on the extended flux.
 We determined an upper limit to the PSF scaling by considering when the integrated flux within the central few pixels was more negative than the noise across clean parts of the image., We determined an upper limit to the PSF scaling by considering when the integrated flux within the central few pixels was more negative than the noise across clean parts of the image.
 It was found that the PSF had to be «0.89 at the c level., It was found that the PSF scaling had to be $< 0.89$ at the $\sigma$ level.
 The scalingPSF-subtracted residual image (center panel of Figure 1)) shows a similar structure to that in (09., The PSF-subtracted residual image (center panel of Figure \ref{fig:image}) ) shows a similar structure to that in G09.
 The ring-like nature of the residuals could be real or could be an artifact of the maximal PSF-subtraction., The ring-like nature of the residuals could be real or could be an artifact of the maximal PSF-subtraction.
" We measure a magnitude of z/=23.45+0.06 within a 2"" radius aperture, consistent with the magnitude measured by (09."," We measure a magnitude of $z'=23.45 \pm 0.06$ within a 2” radius aperture, consistent with the magnitude measured by G09."
 Note the uncertainty on the magnitude quoted above is only due to the measured pixel-to-pixel noise in the sky aperture., Note the uncertainty on the magnitude quoted above is only due to the measured pixel-to-pixel noise in the sky aperture.
 It does not include the uncertainty due to PSF scaling or the effect of correlated pixels., It does not include the uncertainty due to PSF scaling or the effect of correlated pixels.
 The 4c limit based on the PSF scaling discussed previously corresponds to an extended component magnitude limit of z’«23.62., The $\sigma$ limit based on the PSF scaling discussed previously corresponds to an extended component magnitude limit of $z'<23.62$.
 Inspection of the image substantial correlation between nearby pixels due tosuggests a spatiallyvariable background., Inspection of the image suggests substantial correlation between nearby pixels due to a spatiallyvariable background.
" To ascertain the significance of the extended emission in the z' band image, we placed 2"" radius apertures at random across the image (after masking of real sources) and measured the distribution of fluxes in these apertures."," To ascertain the significance of the extended emission in the $z'$ band image, we placed 2” radius apertures at random across the image (after masking of real sources) and measured the distribution of fluxes in these apertures."
 A flux equivalent to magnitude fainter than z'=25.2 is found in of the random apertures and magnitude fainter than z’=24.6 is found of the time., A flux equivalent to magnitude fainter than $z'=25.2$ is found in of the random apertures and magnitude fainter than $z'=24.6$ is found of the time.
" Assuming a Gaussian distribution, we extrapolate this to infer that the extended z emission has a significance of 5σ."," Assuming a Gaussian distribution, we extrapolate this to infer that the extended $z'$ emission has a significance of $5\,\sigma$."
 (09 stated that the extended z’ emission has significance 160.," G09 stated that the extended $z'$ emission has significance $16\,\sigma$."
" However, their random aperture results are similar to ours; they found a lo sky noise of z’=24.9, which would make the detection significant at only 3.66."," However, their random aperture results are similar to ours; they found a $\,\sigma$ sky noise of $z'=24.9$, which would make the detection significant at only $3.6\,\sigma$."
" In conclusion, the extended emission at z’ band is significant at at least the 5c level (recalling that we subtracted off a maximal PSF)."," In conclusion, the extended emission at $z'$ band is significant at at least the $5\,\sigma$ level (recalling that we subtracted off a maximal PSF)."
 With a magnitude of z'=23.5 the extended emission comprises at least of the total flux from this object in this filter passband., With a magnitude of $z'=23.5$ the extended emission comprises at least of the total flux from this object in this filter passband.
 The quasar JJ2329-0301 was observed with the Echelle Spectrograph and Imager (ESI) spectrograph (Sheinis et al., The quasar J2329-0301 was observed with the Echelle Spectrograph and Imager (ESI) spectrograph (Sheinis et al.
 2002) at the Keck-II telescope during the nights of 4+5 October 2007 and 24 September 2008., 2002) at the Keck-II telescope during the nights of 4+5 October 2007 and 24 September 2008.
 The atmospheric conditions were variable and only the data taken in reasonable conditions were included., The atmospheric conditions were variable and only the data taken in reasonable conditions were included.
" The total integration time was 8.5 hours, split into 30 minute exposures with the quasar position offset along the slit between each exposure."," The total integration time was 8.5 hours, split into 30 minute exposures with the quasar position offset along the slit between each exposure."
" The 1"" wide slit was positioned at an angle of 105? East of North, as indicated on Figure 1.."," The 1” wide slit was positioned at an angle of $^\circ$ East of North, as indicated on Figure \ref{fig:image}."
" Note that the spectroscopic observations were performed before publication of G09, so the good alignmentof the slit with the elongation angle of the extended emission is coincidental."," Note that the spectroscopic observations were performed before publication of G09, so the good alignmentof the slit with the elongation angle of the extended emission is coincidental."
" ffalls in order 7 of the echelle, where the spatial pixel scale is 0.163"" and the spectral scale is ! pixel!."," falls in order 7 of the echelle, where the spatial pixel scale is 0.163” $^{-1}$ and the spectral scale is $^{-1}$ $^{-1}$ ."
 The spectral pixel”!resolution with the 1” slit is 75 kmss~! or 4000., The spectral resolution with the 1” slit is 75 $^{-1}$ or $R=4000$ .
 Data reduction was performed using the ESIRedux, Data reduction was performed using the ESIRedux
error bars by 0.1 and 0.18 dex. respectively.,"error bars by $\pm$ 0.4 and $\pm$ 0.18 dex, respectively."
 Ou the other haud. if oue eniplovs the same ilker et al. (," On the other hand, if one employs the same Hilker et al. ("
1995) Strommeren plotometiy of NGC 1866 supereiauts. iu conjunction with the |Fe/H]-Stróunuereun photometry calibration bv. Arellano Ferro Mendoza (1993). au approximately solar |Fe/TII] is obtained.,"1995) Strömmgren photometry of NGC 1866 supergiants, in conjunction with the [Fe/H]-Strömmgren photometry calibration by Arellano Ferro Mendoza (1993), an approximately solar [Fe/H] is obtained."
 lu addition. Feast (1989) has used V7 photometry of Cepheids in the cluster aud deteriuned thei metallicity following the procedure by Caldwell Coulson (1985). obtaining a ean abundauce 0.1 0.2.," In addition, Feast (1989) has used $BVI$ photometry of Cepheids in the cluster and determined their metallicity following the procedure by Caldwell Coulson (1985), obtaining a mean abundance $-0.1\pm$ 0.2."
 Iun light of these existing uncertaimties about NCC 1866 metallicity. new spectroscopic observations of a wide sample of cluster stars (for example by using FLAMESVLT) are urgently needed.," In light of these existing uncertainties about NGC 1866 metallicity, new spectroscopic observations of a wide sample of cluster stars (for example by using FLAMESVLT) are urgently needed."
 We want also to mention the uncertainties due to the calibration of the photometric measurements., We want also to mention the uncertainties due to the calibration of the photometric measurements.
 We consider a possible svstematic zero-poiut error on the V aud { WEPC2 magnitudes of £0.02 mae (Dolphin 20005)., We consider a possible systematic zero-point error on the $V$ and $I$ WFPC2 magnitudes of $\pm$ 0.02 mag (Dolphin 2000b).
 Such a shift iu the V. iiaguitude would mareinally affect the distance values aud do uot help iu decreasing substantially the gap between the MS and RC distances., Such a shift in the $V$ magnitude would marginally affect the distance values and do not help in decreasing substantially the gap between the MS and RC distances.
 Ou the other liaud. a corresponding zero-point error by 0.03 mae on the (V.I) color is expected to produce sizable differences for the MS-fittiug distances due to the slope of the MS itself.," On the other hand, a corresponding zero-point error by 0.03 mag on the $(V-I)$ color is expected to produce sizable differences for the MS-fitting distances due to the slope of the MS itself."
 To quautitatively explore the effect of this source of uncertainty on our distance evaluations we derived again the distances with both the AIS of NGC 1866 aud the RC of the surrouudiug field: by considering point errors of £0.02 mae on both V aud £ magnitudes. the systematic error on the distance moduli cliffereuce A previously defined amounts to £0.19 mae.," To quantitatively explore the effect of this source of uncertainty on our distance evaluations we derived again the distances with both the MS of NGC 1866 and the RC of the surrounding field; by considering zero-point errors of $\pm$ 0.02 mag on both $V$ and $I$ magnitudes, the systematic error on the distance moduli difference $\Delta$ previously defined amounts to $\pm$ 0.19 mag."
 Clearly this may overcome the discrepancy between the two methods oulv in the case that the assumed photometric zero-poiut is in error bv 0.02 mae., Clearly this may overcome the discrepancy between the two methods only in the case that the assumed photometric zero-point is in error by $-$ 0.02 mag.
 The comparison with erouud based photometry (Walker 1995) as pertormed by WOl does uot provide any significant information. since the spread (a~0.06) is so laree that it prevents for any realistic conclusion on the mater.," The comparison with ground based photometry (Walker 1995) as performed by W01 does not provide any significant information, since the spread $\sigma \sim 0.06$ ) is so large that it prevents for any realistic conclusion on the matter."
 Asa last point. let us recall that the previous discussiou relies on the assumption that oth the cluster aud the RC field stars are located at the same distance from us.," As a last point, let us recall that the previous discussion relies on the assumption that both the cluster and the RC field stars are located at the same distance from us."
 If this is nof the case. the cdiscussec dichotomy would iudicate that the cluster is about 5 hpc closer to us than the nuderlving field population of he LMC.," If this is not the case, the discussed dichotomy would indicate that the cluster is about 5 Kpc closer to us than the underlying field population of the LMC."
 To investigate this possibility oue will need to decrumine MS-fittiug distances to a larger sample of LMC clusters. so that depth effects due to the spatial distributio1i of the clusters cancel out when a lca1 value of their distances is computed.," To investigate this possibility one will need to determine MS-fitting distances to a larger sample of LMC clusters, so that depth effects due to the spatial distribution of the clusters cancel out when a mean value of their distances is computed."
 Iu conclusion. we find that the dichotoiuy of the cüstauce of the LMC largely debated iu the literature arises also when the distance to one single cluster (NCC 15060) and the πππποπιο LAIC field are compared.," In conclusion, we find that the dichotomy of the distance of the LMC largely debated in the literature arises also when the distance to one single cluster (NGC 1866) and the surrounding LMC field are compared."
 This iav sugeest that elobal uncertainties iu the methods of determiuiug distances are unuderestimated., This may suggest that global uncertainties in the methods of determining distances are underestimated.
 NCC 1866 also contaius more than 20 Cepheids aud this make this cluster a critical benchmark to probe the distance of the LMC., NGC 1866 also contains more than 20 Cepheids and this make this cluster a critical benchmark to probe the distance of the LMC.
 Available works on Cepheids in NGC 1866 (Cioren et al., Available works on Cepheids in NGC 1866 (Gieren et al.
 2000 aud references thereiu) do not solve defiuitelv the question. eiven that the derived distances show a sizable nucertaity.," 2000 and references therein) do not solve definitely the question, given that the derived distances show a sizable uncertainty."
 Thus. specific high resolution spectroscopy and detailed ποτ curves of the clusters Cepheids are required.," Thus, specific high resolution spectroscopy and detailed light curves of the cluster's Cepheids are required."
contributed by the individual overlapping regionsτρ: To obtain the instantaneous \IDF after 0 generations. we stun the distributious .Vj;(i) for iudividual sets of regions occupied by jf objects. weighted by their P: ⊺∪≺∐⊔↖↸∖↑∐↸∖↑∪↑∡↧↕∖∐≻⋮∪↕≺↧∐∪∩↑∖↸∖↖↸⊓⊓∡↧↑⊓↧∪↖∪ the ∣∣∶↴∙↸∐∪≺↧⊓∪∐↴∖∙↖↖↸∖⋯∏↴∖↑↥⊓⊔↕∩∖ total gas nass after cach generation &À by a factor Dj=1.hd. to account for the couversion of gas into stars;,"contributed by the individual overlapping regions$z_i$: To obtain the instantaneous MDF after $n$ generations, we sum the distributions $N_j(z)$ for individual sets of regions occupied by $j$ objects, weighted by their $P_j$: To derive the total MDF of all objects ever created over $n$ generations, we must reduce the total gas mass after each generation $k$ by a factor $D_k = 1-k\delta$, to account for the conversion of gas into stars."
 We assume the reduction increment à to be constant. so that the preseut-dayv gas fraction ji=Lod.," We assume the reduction increment $\delta$ to be constant, so that the present-day gas fraction $\mu_1 = 1-n\delta$."
" The total. time-imteerated MDE is then: where Py, is P; for a given gcucration &."," The total, time-integrated MDF is then: where $P_{j,k}$ is $P_j$ for a given generation $k$."
" Note that this troatiueut for gas consunptionassunes that 274, roniain uuaffected by the reduction of gas."," Note that this treatment for gas consumptionassumes that $P_{j,k}$ remain unaffected by the reduction of gas."
 We now assume that f(z) rendus the same for all eeneratious., We now assume that $f(z)$ remains the same for all generations.
 For large j. the Ceutral Limit Theorem applies. which predicts that the κος. MDE approximates a uormal distribution with imean ancl variance even by j times the mean e aud variance 6? of f(z).," For large $j$, the Central Limit Theorem applies, which predicts that the summed MDF approximates a normal distribution with mean and variance given by $j$ times the mean $a$ and variance $\sigma^2$ of $f(z)$."
 Thus. Αντ) is elven simply by: where Pj is given by equatiou L..," Thus, $\nins$ is given simply by: where $P_j$ is given by equation \ref{binomial}."
 Likewise. to estiniate Ns) for huge &. the binomial distribution approximates anormal distribution with nean aud variance given by Thus equation 5 may be written: For sinall o. equations 6 and & break down unless f(: describes a normal distribution.," Likewise, to estimate $\ntot$ for large $k$, the binomial distribution approximates a normal distribution with mean and variance given by Thus equation \ref{ntot_gen} may be written: For small $n$, equations \ref{nz} and \ref{ntot} break down unless $f(z)$ describes a normal distribution."
 Several considerations sugecstee that Τί:) should be eiven by a power-law., Several considerations suggest that $f(z)$ should be given by a power-law.
 As a rough estimate. we consider the volume affectedby a star-forming region to be that of the superbubble generated by its IN. core-collapse superuovae (SNe).," As a rough estimate, we consider the volume affected by a star-forming region to be that of the superbubble generated by its $N_*$ core-collapse supernovae (SNe)."
 Following Oev Clarke (1997: hereafter: OC'97). the mechanical power £—N.Beyft. wheref.=10 Myr is the life expectancy of £L for OB associations. aud Lox=10° ere is the individual SN cucreyv.," Following Oey Clarke (1997; hereafter OC97), the mechanical power $L=N_* E_{\rm SN}/t_{\rm e}$, where $t_{\rm e}=40$ Myr is the life expectancy of $L$ for OB associations, and $E_{\rm SN}=10^{51}$ erg is the individual SN energy."
 The mechanical huninosity function of £ is assumed to be a power-law witli iudex 3 (OC'97. eq.," The mechanical luminosity function of $L$ is assumed to be a power-law with index $-\beta$ (OC97, eq."
 1)., 1).
 Π the superbubbles evow until they are coufined by the ambient pressure. the staucard. adiabatic shell evolution inuplies that their final radii will be related to Las RoxLU? (OCOT. eq.," If the superbubbles grow until they are confined by the ambient pressure, the standard, adiabatic shell evolution implies that their final radii will be related to $L$ as $R\propto L^{1/2}$ (OC97, eq."
 29)., 29).
 The total affected volume is therefore. Iu contrast. the total mass of metals produced by I. SNe is given by. where n is the mean vield of iietals per SN.," The total affected volume is therefore, In contrast, the total mass of metals produced by $N_*$ SNe is given by, where $m_y$ is the mean yield of metals per SN."
 Taking 2=ALVA. we find that ;x»L1/52 for the pernbble population. showing that the largest superbubbles generate the lowest imetallicities. owing to dilution iuto larger volume.," Taking $z = M_z/V_s$, we find that $z\propto L^{-1/2}$ for the superbubble population, showing that the largest superbubbles generate the lowest metallicities, owing to dilution into larger volume."
" The MDE for superbubbles is. where N,CR)xRH?) is the size distribution for objects at their ud radi (OCO7. eq."," The MDF for superbubbles is, where $N_s(R)\propto R^{1-2\beta}$ is the size distribution for objects at their final radii (OC97, eq."
 LL)., 41).
 The final relation is obtained for Jj=2. a value typical. perliaps. for all ealaxies (Οσον Clarke 1998).," The final relation is obtained for $\beta=2$, a value typical, perhaps, for all galaxies (Oey Clarke 1998)."
 Hore we see that there are larecr ΗΠΙΟΥΣ of metal-richlies. objects owing to the larger απος of smaller superbub, Here we see that there are larger numbers of metal-rich objects owing to the larger numbers of smaller superbubbles.
 Finally. the probability deusitv for; at anygiven poiut in the ISM is: For 3)=2. we obtain. Since (2) is a probability density function. its iuteeral must be unity. aud C eives the appropriate normalization.," Finally, the probability density for $z$ at anygiven point in the ISM is: For $\beta=2$, we obtain, Since $f(z)$ is a probability density function, its integral must be unity, and $C$ gives the appropriate normalization."
 Thus. for a eiven generation of star formation. f(:) is weighted toward low-netallicitv regions because of their large sizes.," Thus, for a given generation of star formation, $f(z)$ is weighted toward low-metallicity regions because of their large sizes."
 The Πιο values tai. ad nga are determined by the corresponding AZ; aud A., The limiting values $z_{\rm min}$ and $z_{\rm max}$ are determined by the corresponding $R_{\rm max}$ and $R_{\rm min}$.
 For a Salpeter (1955) IME of slope 2.35 aud SN progenitor masses between 8 aud 120 M... we use y= M... based roughly on models bv Woosley Weaver (1995).," For a Salpeter (1955) IMF of slope –2.35 and SN progenitor masses between 8 and 120 $\msol$, we use $m_y=10\ \msol$ , based roughly on models by Woosley Weaver (1995)."
 We take Rivas1300 pc. a rather arbitrary value corresponding to the characteristic parameter 2. of OCO9T: and μμ=25 pe. correspoucding to individual SN remnants.," We take $R_{\rm max} = 1300$ pc, a rather arbitrary value corresponding to the characteristic parameter $R_{\rm e}$ of OC97; and $R_{\rm 
min} = 25$ pc, corresponding to individual SN remnants."
 These vield logtinLF and loguas 3.0. or. taking O as a tracerot primary elements. ninm aud maxima [O/T of 10M.3.0 aud L2> relative to solar.," These yield $\log z_{\rm min}=-4.7$ and $\log z_{\rm max}=-3.0$ , or, taking O as a tracerof primary elements, minimum and maximum [O/H] of –3.0 and –1.3 relative to solar."
" The adopted yield of 0,= may be somewhat lugh(e.g. Woosley Weaver 1995).but the mean f(: iav be somewhat modified depending on tain ALC uas "," The adopted yield of $m_y=10\msol$ may be somewhat high (e.g., Woosley Weaver 1995), but the mean $f(z)$ may be somewhat modified depending on $z_{\rm min}$ and $z_{\rm max}$ ."
We can now construct Monte Carlo models of the AIDF for znall η., We can now construct Monte Carlo models of the MDF for small $n$ .
 We generate the compoucuts ες) by drawiug from £(2) (equation 13)) 7 times and suming the drawu n (equation 3))., We generate the components $N_j(z)$ by drawing from $f(z)$ (equation \ref{fz_pwr}) ) $j$ times and summing the drawn $z_i$ (equation \ref{zsum}) ).
 Du is repeated 5000. times to obtain distributious for Αν)., This is repeated 5000 times to obtain distributions for $N_j(z)$ .
d Equations Lo aud 5. then provide the iustautaneous total AIDFs., Equations \ref{nz_gen} and \ref{ntot_gen} then provide the instantaneous and total MDFs.
 The histogranus in Figures laendeshowthesceinodelsof Nu) aud Nu) for n—2 uud Q—0.72. and pauels b aud d show the same for," The histograms in Figures \ref{montecarlo}$ $a$ and $c$ show these models of $\nins$ and $\ntot$ for $n=2$ and $Q=0.72$, and panels $b$ and $d$ show the same for"
perhaps the completeness thresholds for (he GOODS SN survev. the sensitivity as a function ol change in flux. were significantly overestimated.,"perhaps the completeness thresholds for the GOODS SN survey, the sensitivity as a function of change in flux, were significantly overestimated."
 A simple test of the S04 model would be to push the limits of survey. sensitivity. well bevond those achieved in the GOODS survey. and continue to look for high τους SNe Ia al z>1.4.," A simple test of the S04 model would be to push the limits of survey sensitivity well beyond those achieved in the GOODS survey, and continue to look for high redshift SNe Ia at $z > 1.4$."
 One can (then compare observed vields to those expected in this recdshilt regime from the S04 and other viable models of the SN Ia rate history., One can then compare observed yields to those expected in this redshift regime from the S04 and other viable models of the SN Ia rate history.
" I£729,4,(z) is indeed delaved by ~4Gyr from the SFRy:(2). then it is expected (that the rate at z>1.4 will be very low. and (hus essentially no («& 1) SNe Ia in this range could be found in this deep survey."," If $\mathcal R_{Ia}(z)$ is indeed delayed by $\sim 4$Gyr from the $_U$ $z$ ), then it is expected that the rate at $z> 1.4$ will be very low, and thus essentially no $\ll 1$ ) SNe Ia in this range could be found in this deep survey."
" Even if the πι) is closely tied to the SETo (2). 1, e. with mean delays of less (han approximately one billion vears from star formation. then there could be on the order of only one SN Ia discovery al 2>1.4."," Even if the $\mathcal R_{Ia}(z)$ is closely tied to the $_U$ $z$ ), i. e. with mean delays of less than approximately one billion years from star formation, then there could be on the order of only one SN Ia discovery at $z> 1.4$."
" However. a very meaningful test (his survey could. provide is if the Rig(z) greatly increases with redshift. as it would if it were tied to the SFR,(2). where one could expect at least few SNe In at z21.4 in this deep field survey."," However, a very meaningful test this survey could provide is if the $\mathcal R_{Ia}(z)$ greatly increases with redshift, as it would if it were tied to the $_{IR}$ $z$ ), where one could expect at least few SNe Ia at $z > 1.4$ in this deep field survey."
 We have searched spans of images of the UDF. UDFP. and IRUDE observations [ου high redshift SNe and have found. a total of four events. none of which were al z>1.4.," We have searched spans of images of the UDF, UDFP, and IRUDF observations for high redshift SNe and have found a total of four events, none of which were at $z>1.4$."
 In 82 we describe the UDF+UDFP and IRUDE searches. show discovery ancl subsequent photometric data for the events found in the optical data. ancl present the evidence which show the relatively low redshift of these events.," In 2 we describe the UDF+UDFP and IRUDF searches, show discovery and subsequent photometric data for the events found in the optical data, and present the evidence which show the relatively low redshift of these events."
 In 82 we compare the observed low observed vield al z>1.4 t0 that which would be predicted from a lew SN rate mocels., In 3 we compare the observed low observed yield at $z> 1.4$ to that which would be predicted from a few SN rate models.
 The UDF and UDFP images were searched by differencing stacks of images in the F850LP-band. assembled as Chev were collected over the duration of the deep survey.," The UDF and UDFP images were searched by differencing stacks of images in the $F850LP$ -band, assembled as they were collected over the duration of the deep survey."
 A total of nine multi-orbit epoch image stacks were created for the UDF field. and (wo were made for the UDFP.," A total of nine multi-orbit epoch image stacks were created for the UDF field, and two were made for the UDFP."
 The first stack for each field was cillerenced with an overlapping area in either the ACS GOODS mosaic in the case of the UDF. or with ACS images obtained lor GO 9352 (A. Riess. PI) in the case of the UDEP.," The first stack for each field was differenced with an overlapping area in either the ACS GOODS mosaic in the case of the UDF, or with ACS images obtained for GO 9352 (A. Riess, PI) in the case of the UDFP."
 The subsequent stacks in both survevs were cillerencecl with their preceding stacks., The subsequent stacks in both surveys were differenced with their preceding stacks.
 Similarly. (he IRUDF images were searched by differencing two epochs of mosaic images. each 4 orbits in depth. in both the FLOW and F160M passbands.," Similarly, the IRUDF images were searched by differencing two epochs of mosaic images, each 4 orbits in depth, in both the $F110W$ and $F160W$ passbands."
 The two mosaic image stacks of the IRUDE were separated by approximately 77 davs., The two mosaic image stacks of the IRUDF were separated by approximately 77 days.
 In each the UDF. UDFP. and IRUDE searches. image stacks were produced using (IXoekemoeretal.2002) by using the median of several exposures.," In each the UDF, UDFP, and IRUDF searches, image stacks were produced using \citep{koekemoe2003} by using the median of several exposures."
 This successfully rejects cosmic ravs without rejecting transienis which, This successfully rejects cosmic rays without rejecting transients which
"current Type la supernova limit at ~1.8, therefore providing a reliable test of our proposal when compared to other models.","current Type Ia supernova limit at $\sim1.8$, therefore providing a reliable test of our proposal when compared to other models."
" Standard cosmology is based on the Friedmann-Robertson-Walker (FRW) metric for a spatially homogeneous and isotropic three-dimensional space, in which the coordinates expand or contract as a function of time: The coordinates for this metric have been chosen so that ¢ represents the time measured by a comoving observer (and is the same everywhere, so it functions as a ""community"" time), atf) is the expansion factor, and r is an appropriately scaled radial coordinate in the comoving frame."," Standard cosmology is based on the Friedmann-Robertson-Walker (FRW) metric for a spatially homogeneous and isotropic three-dimensional space, in which the coordinates expand or contract as a function of time: The coordinates for this metric have been chosen so that $t$ represents the time measured by a comoving observer (and is the same everywhere, so it functions as a “community"" time), $a(t)$ is the expansion factor, and $r$ is an appropriately scaled radial coordinate in the comoving frame."
" The geometric factor & is +1 for a closed universe, 0 for a flat, open universe, or —1 for an open universe."," The geometric factor $k$ is $+1$ for a closed universe, $0$ for a flat, open universe, or $-1$ for an open universe."
" Applving the FRW metric to Einstein's field equations of GR, one obtains the corresponding FRW differential equations of motion."," Applying the FRW metric to Einstein's field equations of GR, one obtains the corresponding FRW differential equations of motion."
" These are the Friedmann equation, and the ""acceleration"" equation, An overdot denotes a derivative with respect to cosmic time f£, and p and p represent the total energy density and total pressure, respectively."," These are the Friedmann equation, and the “acceleration"" equation, An overdot denotes a derivative with respect to cosmic time $t$, and $\rho$ and $p$ represent the total energy density and total pressure, respectively."
" A further application of the FRW metric to the energy conservation equation in GR vields the final equation, which, however, is not independent of Equations (3) and (4)."," A further application of the FRW metric to the energy conservation equation in GR yields the final equation, which, however, is not independent of Equations (3) and (4)."
" In comoving coordinates, the proper distance A(t) is measured at constant £ and one can easily see from Equation (2) that for purely radial paths in a flat cosmology, R(r)=a(f)r."," In comoving coordinates, the proper distance $R(t)$ is measured at constant $t$ and one can easily see from Equation (2) that for purely radial paths in a flat cosmology, $R(t)=a(t)r$."
" It is sometimes useful to recast Equation (2) in terms of A(t) (see Equation 9 below) which can reveal, c.9.. the dependence of the metric coefficients on the observer's gravitational horizon, which we now define."," It is sometimes useful to recast Equation (2) in terms of $R(t)$ (see Equation 9 below) which can reveal, e.g., the dependence of the metric coefficients on the observer's gravitational horizon, which we now define."
nuuber density at the disc iudlane. according to the ninimuiunnass solar nebula iodel (ITavashi:1981).. is ~Gs104+ 7.,"number density at the disc midplane, according to the minimum–mass solar nebula model \citep{H81}, is $\sim 6 \times 10^{14}$ $^{-3}$."
" This implies tiat. under the adopted assumptions. the gas at this loc‘ation would be iu the Olunic diffusivity regine for fiek streneths =2.5 cane in the Wall linüt for stronger fiecls,"," This implies that, under the adopted assumptions, the gas at this location would be in the Ohmic diffusivity regime for field strengths $ \lesssim 2.5$ G and in the Hall limit for stronger fields."
 More generally. the location οthe curves delincatiug the regions where cach of fje diffusivity regimes dominates is a function of the mass of the charged particles and the fractional ionisation of the fiuid.," More generally, the location of the curves delineating the regions where each of the diffusivity regimes dominates is a function of the mass of the charged particles and the fractional ionisation of the fluid."
 If the charged particles are more niassive. the Aabipolar-Hall and TallOli transitions will «)neur at stronger fields for a given value of the neutral deusity.," If the charged particles are more massive, the Ambipolar-Hall and Hall-Ohm transitions will occur at stronger fields for a given value of the neutral density."
 Similarly. if the jionisatiou fractiou dncreases (6.B.o the neutral deusitv is a lower ratio of the tota eusitv of the fluid}. the above transitions will take place at comparatively lower magnetic field strengths.," Similarly, if the ionisation fraction increases (e.g. the neutral density is a lower ratio of the total density of the fluid), the above transitions will take place at comparatively lower magnetic field strengths."
 À ]xuwtieularlv imaportaut scenario is when dust particles are wel mused with the eas., A particularly important scenario is when dust particles are well mixed with the gas.
 The deeree of fieldmatter coupling is strongly dependent on the abundance and size distribution of dust particles mixed with the gas., The degree of field–matter coupling is strongly dependent on the abundance and size distribution of dust particles mixed with the gas.
 Dust grains tend to dramatically reduce the fraclonal lonisation of the fluid. as charged. particlos can stick to thou iid recombine on their surfaces.," Dust grains tend to dramatically reduce the fractional ionisation of the fluid, as charged particles can stick to them and recombine on their surfaces."
 Additionally. grains can also become an imuportaut species in ligh-densityv regions (e.g.Uinebavashi&Nakano—1990:Nishietal— 1991).," Additionally, grains can also become an important species in high-density regions \citep[e.g.][]{UN90, NNU91}."
 As they have relatively large CLOSS sections. erains typically becon16 decoupled from the magnetic field at lower deusities than those for which other simaller species do.," As they have relatively large cross sections, grains typically become decoupled from the magnetic field at lower densities than those for which other –smaller– species do."
 As a result. the coucuctivity 6 the fuid diminishes when dust particles (particularly when these are siuall) are well mise: with the eas.," As a result, the conductivity of the fluid diminishes when dust particles (particularly when these are small) are well mixed with the gas."
" This is most relevant at carly stages of the accretion process. or When turbulent motions prevent the grains frou, settling towards the dise midplane."," This is most relevant at early stages of the accretion process, or when turbulent motions prevent the grains from settling towards the disc midplane."
 Iu fact. recent diffusivity calculations for ΕΠlaass solar uchula disc at 1 AU Wardle2007: WSII) have shown that the maeuetic coupling mas* be adequate over the entire disc cross-section. provicl‘cd that dust eraius are settled out of the gas phase.," In fact, recent diffusivity calculations for a minimum-mass solar nebula disc at 1 AU \citealt{W07}; WS11) have shown that the magnetic coupling may be adequate over the entire disc cross-section, provided that dust grains are settled out of the gas phase."
 A word of caution applies to the above stateinents., A word of caution applies to the above statements.
 The turbulent stimiug of dust erains is. evideutly. ependent on the dise beiug able ο sustain turbulence when dust eras are well mixed with the gas.," The turbulent stirring of dust grains is, evidently, dependent on the disc being able to sustain turbulence when dust grains are well mixed with the gas."
 However. as dise turbulence is thought to be a magneticallyriven process (soe Section 5). re reduction in the degree of feld-matter coupling nacelated by dust eraius nay. in turn. weaken the turhtuent imiotions. or even prevent their development iu dusty plasinas.," However, as disc turbulence is thought to be a magnetically-driven process (see Section \ref{sec:MRI}) ), the reduction in the degree of field-matter coupling mediated by dust grains may, in turn, weaken the turbulent motions, or even prevent their development in dusty plasmas."
 Disc vunanues is then the result of a colplex iuterplav: When dust erains are well ήχος with the eas. the isc nav not be susceptible to turbulence. aud the eralus lay ¢ficiently settle to the midplane.," Disc dynamics is then the result of a complex interplay: When dust grains are well mixed with the gas, the disc may not be susceptible to turbulence, and the grains may efficiently settle to the midplane."
 On the other haud. ouce the eraius are effectively removed from the gas. turbulence may be viable again. aud the turbulent uxtions mav stir the eraius back iuto the eas phase.," On the other hand, once the grains are effectively removed from the gas, turbulence may be viable again, and the turbulent motions may stir the grains back into the gas phase."
" The :osulting structure auc evolution of discs. incorporating this important feedback effect. have not vot beeu calc""lated selt-cousisteutlv."," The resulting structure and evolution of discs, incorporating this important feedback effect, have not yet been calculated self-consistently."
 Alost current models of accretion discs are viscous models that adopt. im some form. the o-prescription of Shalkura&Svuuvaev(1973).," Most current models of accretion discs are viscous models that adopt, in some form, the $\alpha$ -prescription of \citet{SS73}."
". In. this formulation. the racial-azniuthal component of the stress tensor is assuned to scale with the gas pressure. or (4,=al."," In this formulation, the radial-azimuthal component of the stress tensor is assumed to scale with the gas pressure, or $w_{r \phi} = \alpha P$."
 The associated turbulent viscosity is paramoeterised as OM;=αιShige. where ος is the isothermal sound speed and fig is the tidal densitv scale-height of the cise (6.8. resulting frou the vertical gravitational compression).," The associated turbulent viscosity is parameterised as $\nu_{t} = \alpha c_{\rm s} h_{\rm T}$, where $c_{\rm s}$ is the isothermal sound speed and $h_{\rm T}$ is the tidal density scale-height of the disc (e.g. resulting from the vertical gravitational compression)."
 Despite its usefuluess iu providing au operational yauework to support the modelling of accretion discs. this methodology does not offer any explanation or the physical nature of the accretion torque. Whose origi remains unspecified.," Despite its usefulness in providing an operational framework to support the modelling of accretion discs, this methodology does not offer any explanation for the physical nature of the accretion torque, whose origin remains unspecified."
 Iu fact. iu this forimlation. all the uukuowus aud uucertauties relating to he accretion stress are just effectively huuped iuto he couvenicut parameter a.," In fact, in this formulation, all the unknowns and uncertainties relating to the accretion stress are just effectively lumped into the convenient parameter $\alpha$."
 Iu seucral. one of the most natural wavs of generatiug turbulence is via hvdrodsaiuuiicaliustabilities.," In general, one of the most natural ways of generating turbulence is via hydrodynamical instabilities."
 Laboratory shear flows. for example. readily break iuto turbulence at sufficicutly μοι ummbers.," Laboratory shear flows, for example, readily break into turbulence at sufficiently high numbers."
" However. Keplerian satisfs* the Ravleigl’s bydrodvuamical stability criterion (angular momentum mereases with radius). so the οσοvation and sustamniuse of hydrodynanmüc turbulence 1n these systems is. at best. unproven (ο,OoseerevicavbyBalbus&Dawley1998)."," However, Keplerian satisfy the Rayleigh's hydrodynamical stability criterion (angular momentum increases with radius), so the generation and sustaining of hydrodynamic turbulence in these systems is, at best, unproven \citep[e.g.~see review by][]{BH98}."
. Convective turbulence has also been considered as au option (Lin&Papaloizou1980:Linetal. 1993).. but further studies appear fo 1idicate that this imechanisni transports angular wouieutunmi towards the central object (CabotStone&Babus 1996).," Convective turbulence has also been considered as an option \citep{LP80, LPK93}, but further studies appear to indicate that this mechanism transports angular momentum towards the central object \citep{CP92, RG92, C96, SB96}."
. More generally. the sign of the radial fix generated by convective turbulence iav depend on t1e ratio of the epicyclic frequency aud the," More generally, the sign of the radial flux generated by convective turbulence may depend on the ratio of the epicyclic frequency and the"
"We also consider the effects of changing the initial and boundary temperatures of the disces as well as simulating clilferent disc sizes (with outer disc radii. 425,4=25 and 300 AU).","We also consider the effects of changing the initial and boundary temperatures of the discs as well as simulating different disc sizes (with outer disc radii, $R_{\rm out} = 25$ and 300 AU)."
 We find that the disc opacity ds νον important in determining whether a cise is likely to fragment., We find that the disc opacity is very important in determining whether a disc is likely to fragment.
 In particular. we find that fragmentation is promoted in disces with opacity values lower than interstellar Iosseland. mean values since this allows radiation to leave the cise quickly.," In particular, we find that fragmentation is promoted in discs with opacity values lower than interstellar Rosseland mean values since this allows radiation to leave the disc quickly."
 Low opacities may exist in low metallicity clises or discs with larger grain sizes., Low opacities may exist in low metallicity discs or discs with larger grain sizes.
 Phis is à particularly important ancl timely result eiven the recent discoveries of wide orbit planets (22) and the future. emphasis for survevs of planets on such wide orbits.," This is a particularly important and timely result given the recent discoveries of wide orbit planets \citep{Fomalhaut,HR8799} and the future emphasis for surveys of planets on such wide orbits."
 We show that it is possible for fragmentation to occur in gravitationally unstable disces even at radii where the innermost planet of the LR 8799. svstem is located (zT24 AU)., We show that it is possible for fragmentation to occur in gravitationally unstable discs even at radii where the innermost planet of the HR 8799 system is located $R \gtrsim 24$ AU).
 Furthermore. LR 8799. is known to be a metal-poor. A Bootis star with metallicity AZ/44]=0.47 (7) so it is reasonable to assume that its dise was similarly metal-poor.," Furthermore, HR 8799 is known to be a metal-poor, $\lambda$ Bootis star with metallicity $[M/H] = -0.47$ \citep{HR8799_metallicity} so it is reasonable to assume that its disc was similarly metal-poor."
 We have shown that such a scenario. favours fragmentation anc therefore. our results indicate that all three planets of the LIB. 8799 svstem may well have formed via gravitational instability.," We have shown that such a scenario favours fragmentation and therefore, our results indicate that all three planets of the HR 8799 system may well have formed via gravitational instability."
 Though a hybrid core accretion and gravitational instability scenario for planet. formation niv also be a possibility for this system. our calculations show that such a hybrid scenario may not be necessary.," Though a hybrid core accretion and gravitational instability scenario for planet formation may also be a possibility for this system, our calculations show that such a hybrid scenario may not be necessary."
 We find that the. presence of a as a result of the stellar iraciation inhibits fragmentation since the disces are only able to cool to the boundary temperature., We find that the presence of a as a result of the stellar irradiation inhibits fragmentation since the discs are only able to cool to the boundary temperature.
 However. we also show that uncer certain circumstances. fragmentation mav occur.," However, we also show that under certain circumstances, fragmentation may occur."
 Our results demonstrate that. for. fragmentation. weak irradiation is required such that the boundary. temperature and. hence ‘Toomre stability parameter is low. in addition to low enough opacities (even in large. cool discs) since this allows more cHicient cooling so that the discs temperature does not increase significantly (due to internal dissipation) above the boundary temperature.," Our results demonstrate that for fragmentation, weak irradiation is required such that the boundary temperature and hence Toomre stability parameter is low, in addition to low enough opacities (even in large, cool discs) since this allows more efficient cooling so that the disc's temperature does not increase significantly (due to internal dissipation) above the boundary temperature."
 We thank the anonymous referee whose comments greatly improved the clarity of the paper., We thank the anonymous referee whose comments greatly improved the clarity of the paper.
 The calculations reported here were performed. using the University of Exeter's SL Altix LCE 8200. supercomputer., The calculations reported here were performed using the University of Exeter's SGI Altix ICE 8200 supercomputer.
 “The disc images were produced. using ΘΕΟΙΛΑΟΙ (2).., The disc images were produced using SPLASH \citep{SPLASH}.
 AIRB is. grateful for the support of a EURYL Award which also funded FAL, MRB is grateful for the support of a EURYI Award which also funded FM.
 This work. conducted as part of the award “Lhe formation of stars ancl planets: Racliation hvdrodynamical and maenetohyedrocdvonamical simulations’ mace under the European Leads of Research Councils anc European Science Foundation EURYL (IZuropean. Young Investigator) Awards scheme. was supported by funds from the participating organizations of EURYL and the EC Sixth Framework Programme.," This work, conducted as part of the award `The formation of stars and planets: Radiation hydrodynamical and magnetohydrodynamical simulations' made under the European Heads of Research Councils and European Science Foundation EURYI (European Young Investigator) Awards scheme, was supported by funds from the participating organizations of EURYI and the EC Sixth Framework Programme."
in Table 5.,in Table 5.
 For NGC 3783 and. ALR 2251-178 the accretion rates are based on the observed bolometric Iuminosities. and not on the values determined during our SED fitting.," For NGC 3783 and MR 2251-178 the accretion rates are based on the observed bolometric luminosities, and not on the values determined during our SED fitting."
 This is because our determination of Al only accounts for the direct. release of gravitational power into the disc. while the X-ray heating is determined. by the observed: value of Ly.," This is because our determination of $\dot{M}$ only accounts for the direct release of gravitational power into the disc, while the X-ray heating is determined by the observed value of $L_{\rm x}$."
 lMowever. the high. energy. emission (and any released mechanical energv) is also powered. by the eravitational potential energy. which ultimately takes the form of a hot corona (and a mechanical jet).," However, the high energy emission (and any released mechanical energy) is also powered by the gravitational potential energy which ultimately takes the form of a hot corona (and a mechanical jet)."
 Since we do not have à prescription to relate AL with Ly. Ly remains as an extra parameter in the SED mocels.," Since we do not have a prescription to relate $\dot{M}$ with $L_{\rm x}$, $L_{\rm x}$ remains as an extra parameter in the SED models."
 As before. measurements of AZ and AZ for NGC 3783 are fairly accurate. while clear uncertainties remain [for AIR. 2251-178.," As before, measurements of $M$ and $\dot{M}$ for NGC 3783 are fairly accurate, while clear uncertainties remain for MR 2251-178."
" Besides. the black hole estimates presented in Section 2. the determination of the bolometric luminosity for AIR. 2251-178 is rather vague: assuming the bolometric corrections computed by Marconi et ((2004) we find that the B-band luminosity predicts Lie~2107 ergs/s. while the 2-10 keV luminosity gives Ly.)~2107 eres/s. This shows that AIR 2251-178 is fairly ""X-ray loud”. with an ἂν, index of 1.2 (for a definition of 6, see Tananbaum ct al.."," Besides, the black hole estimates presented in Section 2, the determination of the bolometric luminosity for MR 2251-178 is rather vague: assuming the bolometric corrections computed by Marconi et (2004) we find that the B-band luminosity predicts $L_{bol} \sim 2\times10^{45}$ ergs/s, while the 2-10 keV luminosity gives $L_{bol} \sim
2\times10^{46}$ ergs/s. This shows that MR 2251-178 is fairly “X-ray loud”, with an $\alpha_{ox}$ index of 1.2 (for a definition of $\alpha_{ox}$ see Tananbaum et al.,"
 1979). close to the edge of the distribution for quasars and local Sevferts (Zamorani et al..," 1979), close to the edge of the distribution for quasars and local Seyferts (Zamorani et al.,"
. 1981: Cirupe et. al..," 1981; Grupe et al.,"
. 2010)., 2010).
 Llence. there is about one order of magnitude uncertainty in the bolometric luminosity of ALR 2251-178. as well as in its accretion rate.," Hence, there is about one order of magnitude uncertainty in the bolometric luminosity of MR 2251-178, as well as in its accretion rate."
 ‘Table 5 shows that MIU 2251-178 is located in the lower part of the distribution of (AL/AL)iljli values when compared with local NGC Sevferts and. if its lower accretion rate is confirmed. it fits well within the range of GMAMybU values shown by Ixishimoto's sample of Quasars.," Table 5 shows that MR 2251-178 is located in the lower part of the distribution of $(\dot{M}/M)^{(1/4)}$ values when compared with local NGC Seyferts and, if its lower accretion rate is confirmed, it fits well within the range of $(\dot{M}/M)^{(1/4)}$ values shown by Kishimoto's sample of Quasars."
 However. the power to 1/4 clearly introduces a weak dependency on AL/AL. and eiven the uncertainties involved in these measurements it is still premature to draw firm conclusions.," However, the power to $1/4$ clearly introduces a weak dependency on $\dot{M}/M$, and given the uncertainties involved in these measurements it is still premature to draw firm conclusions."
 In fact. NGC 4395 (for which the detected intra-day near-L1i variations also argue in Favor of J and L-band disc emission: Alinezaki et abl.," In fact, NGC 4395 (for which the detected intra-day near-IR variations also argue in favor of J and H-band disc emission; Minezaki et al.,"
 2006). has a value of GMALL well consistent with the other NGC ealasxies.," 2006), has a value of $(\dot{M}/M)^{(1/4)}$ well consistent with the other NGC galaxies."
 Hence. even though ‘Table 5 could. be telling us that MIU 2251-178 hosts the coolest accretion disc it remains to be tested further whether the combination of these two fundamental parameters. Al and AZ. is sullicient to fully determine the temperature clistribution in thin accretion disces.," Hence, even though Table 5 could be telling us that MR 2251-178 hosts the coolest accretion disc it remains to be tested further whether the combination of these two fundamental parameters, $M$ and $\dot{M}$, is sufficient to fully determine the temperature distribution in thin accretion discs."
 As already shown in this work the geometry of the disc (e.g. the presence of IEaring. humps. or tapering of the disc) might also play a significant role.," As already shown in this work the geometry of the disc (e.g., the presence of flaring, humps, or tapering of the disc) might also play a significant role."
" This finding could have important implications for the analysis of other sources with small (AL/AL)'4°""! patios where the near-LRo emission. is normally interpreted: as a dust emission. while it might correspond. to emission. [rom the accretion disc.", This finding could have important implications for the analysis of other sources with small $(\dot{M}/M)^{(1/4)}$ ratios where the near-IR emission is normally interpreted as a dust emission while it might correspond to emission from the accretion disc.
 We have obtained high quality near-Ht light. curves for two nearby AGN: NGC 3783 and. ME 2251-178. which are characterised. by very dillerent values of their black hole mass.," We have obtained high quality near-IR light curves for two nearby AGN: NGC 3783 and MR 2251-178, which are characterised by very different values of their black hole mass."
 Phe light curves track variations on time-scales [rom a few days up to three vears for NGC 3783 and four vears for ALR. 2251-178., The light curves track variations on time-scales from a few days up to three years for NGC 3783 and four years for MR 2251-178.
 From the analysis of these data. together with the results already. presented in Aréyvvalo et ((2008. 2009) we find the following: We thank the anonymous reforee. for. helpful. comments.," From the analysis of these data, together with the results already presented in Arévvalo et (2008, 2009) we find the following: We thank the anonymous referee for helpful comments."
 PL gratefully acknowledges support bv. Fondecyt Provect 1080603 and Fondap Project 1501003., PL gratefully acknowledges support by Fondecyt Proyect 1080603 and Fondap Project 1501003.
more complicated dependencies obviously cannot be ruled out. a general decline of the Ne/O abundance with activity 15 present in all three regression curves.,"more complicated dependencies obviously cannot be ruled out, a general decline of the Ne/O abundance with activity is present in all three regression curves."
 Admittedly some of the results for individual stars have large statistical errors and our stellar sample ts quite limited. however our results clearly suggest low Ne/O ratios for the weakly active stars.," Admittedly some of the results for individual stars have large statistical errors and our stellar sample is quite limited, however our results clearly suggest low Ne/O ratios for the weakly active stars."
 Also. the overplotted range of ‘classical’ solar values for activity and Ne/O ratio fits well into the respective parameter space derived from our low-activity stars.," Also, the overplotted range of 'classical' solar values for activity and Ne/O ratio fits well into the respective parameter space derived from our low-activity stars."
 Our results also agree reasonably well with available literature values (transformed to ? abundance scale) derived from various. multi-temperature EMD modeling. despite different applied methods in data analysis. underlying atomic data. data sets or selection criteria.," Our results also agree reasonably well with available literature values (transformed to \cite{grsa} abundance scale) derived from various multi-temperature EMD modeling, despite different applied methods in data analysis, underlying atomic data, data sets or selection criteria."
 Concerning the individual targets. the most active star in our sample. the K-dwarf e EEri. shows the highest ratio with 00.4.," Concerning the individual targets, the most active star in our sample, the K-dwarf $\epsilon$ Eri, shows the highest ratio with 0.4."
 A similar range of values is present in the literature. e.g. Ne/O=0.36 (2?) or Ne/O=0.4 (?)..," A similar range of values is present in the literature, e.g. Ne/O=0.36 \citep{wood06} or Ne/O=0.4 \citep{san04}."
 For the group of slightly less active stars we measure = 00.35., For the group of slightly less active stars we measure $\approx$ 0.35.
 The stars CCyg. BCCom and 881809 have very similar activity levels and there might to be a trend towards lower Ne/O ratios for stars of earlier spectral type at a given activity level.," The stars Cyg, $\beta$ Com and 81809 have very similar activity levels and there might to be a trend towards lower Ne/O ratios for stars of earlier spectral type at a given activity level."
 Especially the GO- B starCCom ts better described by a ratio around Ne/O=0.25. but given the errors and the small sample size of only three stars. this has to be checked with larger and better observed stellar samples.," Especially the G0-star $\beta$ Com is better described by a ratio around Ne/O=0.25, but given the errors and the small sample size of only three stars, this has to be checked with larger and better observed stellar samples."
 For the less active a CCen system we find values around Ne/Oz0.25., For the less active $\alpha$ Cen system we find values around Ne/O=0.25.
 Additionally. in an analysis. of selected phases of the « CCen data where the contribution of the lesser active component aCCen A to the RGS spectra is maximal. indications of a decrease in the Ne/O ratio of the « CCen system are present.," Additionally, in an analysis of selected phases of the $\alpha$ Cen data where the contribution of the lesser active component $\alpha$ Cen A to the RGS spectra is maximal, indications of a decrease in the Ne/O ratio of the $\alpha$ Cen system are present."
 This points to a lower Ne/O ratio in a CCen A compared to aCCen B. ?. derived values of Ne/O=0.18+0.07 (A) and 0.240.069 (B) for the individual components of the a CCen system from the LETGS data. slightly lower than our values. but showing the same trend in the Ne/O ratio of the components as our results for the phase selected data.," This points to a lower Ne/O ratio in $\alpha$ Cen A compared to $\alpha$ Cen B. \cite{raa03} derived values of $\pm 0.07$ (A) and $\pm 0.09$ (B) for the individual components of the $\alpha$ Cen system from the LETGS data, slightly lower than our values, but showing the same trend in the Ne/O ratio of the components as our results for the phase selected data."
 Also. the Ne/O=0.27 result from their line based analysis in 2. fits into this picture. when keeping in mind that they used the RGS data of the « CCen system from 22005. ie. intermediate in terms of activity and contributing components between our 2003/04 dataset and the full data selection.," Also, the Ne/O=0.27 result from their line based analysis in \cite{lie06} fits into this picture, when keeping in mind that they used the RGS data of the $\alpha$ Cen system from 2005, i.e. intermediate in terms of activity and contributing components between our 2003/04 dataset and the full data selection."
 However. given the measurement errors this has to be confirmed by further observations.," However, given the measurement errors this has to be confirmed by further observations."
 Procyon. the least active star in our sample beside a CCen A. again shows a very low Ne/O ratio for most of the data used.," Procyon, the least active star in our sample beside $\alpha$ Cen A, again shows a very low Ne/O ratio for most of the data used."
 While ? derived Ne/O=0.4 from the LETGS data (AQ0+AO1). ? derived a value of Ne/O=0.22 for the RGS as well as the LETGS data (AOI).," While \cite{san04} derived Ne/O=0.4 from the LETGS data (AO0+AO1), \cite{raa02} derived a value of Ne/O=0.22 for the RGS as well as the LETGS data (AO1)."
 Some of these discrepancies are possibly due to the different analyzed datasets as discussed above., Some of these discrepancies are possibly due to the different analyzed datasets as discussed above.
 Although the even lower Ne/O ratios derived from the linear combination of D&TT probably have to be at least partially attributed to systematic errors. values around Ne/O=0.2 are found repeatedly in various global analysis and linear combination models for Procyon.," Although the even lower Ne/O ratios derived from the linear combination of T probably have to be at least partially attributed to systematic errors, values around Ne/O=0.2 are found repeatedly in various global analysis and linear combination models for Procyon."
 A value of Ne/O=0.2 also corresponds to the mean from all used methods and datasets. whereas our averaged global analysis yields Ne/O=0.22 for Procyon.," A value of Ne/O=0.2 also corresponds to the mean from all used methods and datasets, whereas our averaged global analysis yields Ne/O=0.22 for Procyon."
 In summary. for moderately active G- and K-dwarfs with activity levels of LLy/Ly =—5..—5.5 we measure abundance ratios of Ne/O=0.300.4.," In summary, for moderately active G- and K-dwarfs with activity levels of $_{\rm X}$ $_{\rm bol}\approx -5\,...-5.5$ we measure abundance ratios of 0.4."
 While the Ne/O ratio further decreases with decreasing stellar activity. values around Ne/Oz0.200.25 are found exclusively in the least active stars with activity levels around and below logLLx/Li =—6.5.," While the Ne/O ratio further decreases with decreasing stellar activity, values around 0.25 are found exclusively in the least active stars with activity levels around and below $_{\rm X}$ $_{\rm bol}\approx -6.5$ ."
 Additionally. none of these low activity stars shows a significantly higher ratio. therefore low Ne/O ratios seem to be common for stars with activity levels comparable to the Sun.," Additionally, none of these low activity stars shows a significantly higher ratio, therefore low Ne/O ratios seem to be common for stars with activity levels comparable to the Sun."
 This finding is also supported by phase selected and spatially resolved analysis of the a@CCen system. however measurement errors become large for these data.," This finding is also supported by phase selected and spatially resolved analysis of the $\alpha$ Cen system, however measurement errors become large for these data."
 A possible further dependence of spectral type. activity phase or other stellar parameters like lummosity class. binarity ete.," A possible further dependence of spectral type, activity phase or other stellar parameters like luminosity class, binarity etc."
 might be present. but cannot be investigated with the present data.," might be present, but cannot be investigated with the present data."
 We finally discussthe observed change of the Ne/O ratio in the context of more active stars., We finally discussthe observed change of the Ne/O ratio in the context of more active stars.
 The ratio increases from Ne/Ox0.2 up to Ne/O=0.4 when going from low to moderately activity stars. re. from LLy/Lyo =-7 to LLs/Lpo—5.," The ratio increases from $\approx$ 0.2 up to $\approx$ 0.4 when going from low to moderately activity stars, i.e. from $_{\rm X}$ $_{\rm bol}\approx -7$ to $_{\rm X}$ $_{\rm bol}\approx -5$ ."
 Given these finding it is not surprising that ? found an average ratio of Ne/O=0.41 for a stellar sample dominated by stars or binary systems that are mostly ten to hundred times more active than the stars in our sample., Given these finding it is not surprising that \cite{dra05} found an average ratio of Ne/O=0.41 for a stellar sample dominated by stars or binary systems that are mostly ten to hundred times more active than the stars in our sample.
 If anything. in the light of the inferred trend for the Ne/O ratio. their value appears rather low than high.," If anything, in the light of the inferred trend for the Ne/O ratio, their value appears rather low than high."
 This further points to a different behavior when going to more extreme activity levels., This further points to a different behavior when going to more extreme activity levels.
 While a linear dependence between Ne/O ratio and LLs/Lypo is a fair description for our stellar sample. we already noted that other models would also fit the data.," While a linear dependence between Ne/O ratio and $_{\rm X}$ $_{\rm bol}$ is a fair description for our stellar sample, we already noted that other models would also fit the data."
 Taking their active main sequence stars and binary systems into account. we suspect the Ne/O ratio to saturate already at values around = 00.5.," Taking their active main sequence stars and binary systems into account, we suspect the Ne/O ratio to saturate already at values around $\approx$ 0.5."
 However. the detailed investigation of such a flattening or saturation of the Ne/O ratio towards very active stars is beyond the scope of this study. but deserves further attention.," However, the detailed investigation of such a flattening or saturation of the Ne/O ratio towards very active stars is beyond the scope of this study, but deserves further attention."
 When putting the derived stellar Ne/O ratios in the context with the Sun. we first have to consider the solar activity level and the well known variations of the Sun's X-ray emission over the course of the activity cycle.," When putting the derived stellar Ne/O ratios in the context with the Sun, we first have to consider the solar activity level and the well known variations of the Sun's X-ray emission over the course of the activity cycle."
 Further. most of the solar data refer to individual regions and the usually spatially resolved solar data has to be transformed into some kind of point-source like data that is available for stars.," Further, most of the solar data refer to individual regions and the usually spatially resolved solar data has to be transformed into some kind of point-source like data that is available for stars."
 Using Yohkoh/SXT measurement take1 from 11995 and covering roughly a solar half cycle from maximum to minimum. ? derived a conversion factor between the observed X-ray index and the solar energy flux in the energy range kkeV. leading to a solar activity range of LLx/Li4=—7.9...—6.5 and average coronal temperatures between 2—3MMRK.," Using /SXT measurement taken from 1995 and covering roughly a solar half cycle from maximum to minimum, \cite{act96} derived a conversion factor between the observed X-ray index and the solar energy flux in the energy range keV, leading to a solar activity range of $_{\rm X}$ $_{\rm bol}\approx -7.9\,...-6.5$ and average coronal temperatures between MK."
 Another approach based on Yohkoh/SXT data made by ? resulted in à solar activity range over the cycle of LLy/Lyo= 6.0. while ? derived a slightly more active Sur with an activity range of LLy/Ly=—6.9...—5.8 from SNOE/SXP data.," Another approach based on /SXT data made by \cite{per00} resulted in a solar activity range over the cycle of $_{\rm X}$ $_{\rm bol} \approx -7.2\,...-6.0$ , while \cite{jud03} derived a slightly more active Sun with an activity range of $_{\rm X}$ $_{\rm bol} \approx -6.9\,...-5.8$ from /SXP data."
 These literature values were transformed to our 33.0kkeV band by a reduction of the giver, These literature values were transformed to our keV band by a reduction of the given
which illustrates how hierarchical formation naturally arises In the canonical approach.,which illustrates how hierarchical formation naturally arises in the canonical approach.
 Tn δ we modify our model to Ἱποπιάς ACN feedback. aud show how this leads to a radieallv different. auti-luerarchical history.," In 3 we modify our model to include AGN feedback, and show how this leads to a radically different, anti-hierarchical history."
" We compare these results with recent observations aud conclude with a short discussion iu {,", We compare these results with recent observations and conclude with a short discussion in 4.
" Throughout this study. we assume a ACDAL model with parameters 5=0.65. O4 = 0.3. O4 = 0.7. QO,=0.05. σς=OST. and »=lL. where 7 is the IIubble coustaut iu units of 100 kins t Mpe.ο, Qu. Q4. and Ον are the total latter. vacunm energv. aud barvouic deusities iu units of the critical density. oz is the variance of lear fluctuations ou the SA.1Mpe scale. aud a is tho tilt” of the primordial power spectrum Spergel 2003)."," 		 Throughout this study, we assume a $\Lambda$ CDM model with parameters $h=0.65$, $\Omega_0$ = 0.3, $\Omega_\Lambda$ = 0.7, $\Omega_b = 0.05$ , $\sigma_8
= 0.87$, and $n=1$, where $h$ is the Hubble constant in units of 100 km $^{-1}$ $^{-1}$, $\Omega_0$, $\Omega_\Lambda$, and $\Omega_b$ are the total matter, vacuum energy, and baryonic densities in units of the critical density, $\sigma_8^2$ is the variance of linear fluctuations on the $8 h^{-1}{\rm Mpc}$ scale, and $n$ is the “tilt” of the primordial power spectrum Spergel 2003)."
 The Eiseusteiu IIu (1999) transfer function is adopted iu all cases., The Eisenstein Hu (1999) transfer function is adopted in all cases.
 We begin with a model of coolne-reeulated galaxy ormation aud adopt as simple an approach as possible. to Hehbeht the key pliysical issues.," We begin with a model of cooling-regulated galaxy formation and adopt as simple an approach as possible, to highlight the key physical issues."
 Unlike more sophisticated approaches Wauffinann 1993: Somerville Primack 1999: Beuson 2000). we do not attempt to race the detailed history of a sauple of DM halos through he use of statistical 12erger trees.”," Unlike more sophisticated approaches Kauffmann 1993; Somerville Primack 1999; Benson 2000), we do not attempt to trace the detailed history of a sample of DM halos through the use of statistical “merger trees.”"
 Dustead. we use the fit described in vau den Bosch (2002). which gives theeveerageustory of a DAL halo with a preseut total mass of Aj as - ∙↽ where ov=131.|Ls6los(lter) 10323A£.).. MU) ds the amass of this svsteni at some redshift z. and zie is defined implicitly through he velation Diva)1=1|0.10/57(0.25IM)|e(My). with D being the liucar groxeth factor aud 6?(M) being the variance of liuear fluctuations within a sphere euclosiug a otal mass M. Following the canonical picture. we assume that at cach vedshitt eas flows onto the erowine halo along with he dark inatter. is shocked to its virial density aid eiiperature at that redshift {LS0 tinies the mean density and 7.2<10?(ALC)/1012ALPPP| K. respectively]. and is finally added to the pool of cold gas that forms stars after a delay of = 13S n)bis (to)A," Instead, we use the fit described in van den Bosch (2002), which gives thehistory of a DM halo with a present total mass of $M_0$ as ]= -0.301, where $\nu \equiv 1.34 + 1.86 \, {\rm log}(1+z_{\rm fit}) - 0.03 \,
{\rm log}(M_0/10^{12} \msun)$ , $M(z)$ is the mass of this system at some redshift $z$, and $z_{\rm fit}$ is defined implicitly through the relation $D(z_{\rm fit})^{-1} = 1 + 0.40 \sqrt{\sigma^2 (0.254 M_0) -
\sigma^2(M_0)},$ with $D$ being the linear growth factor and $\sigma^2(M)$ being the variance of linear fluctuations within a sphere enclosing a total mass $M.$ Following the canonical picture, we assume that at each redshift gas flows onto the growing halo along with the dark matter, is shocked to its virial density and temperature at that redshift $180$ times the mean density and $7.2 \times 10^5 \, (M(z)/10^{12} M_\odot)^{2/3} \, (1+z)$ K, respectively], and is finally added to the pool of cold gas that forms stars after a delay of = 1.3 T_6 <n_e (T_6)."
"\Iyrs.. Here 6023 is the average ummber of electrous per cu? iu the shocked gas. C—Gi?)da)2 isa “clumping factor. which accounts for inhomogeneities. Z5 is the teniperature in units of 105 Ik. and Aον is the radiative cooling rate of the eas iu units of 107° cres ceni? st, Nou 1 keV zmloo 1. Aο is roughly a constant but in general it is à function of temperature aud metallicity."," Here $\left< n_e \right>$ is the average number of electrons per $^{3}$ in the shocked gas, $C \equiv \left< n_e^2 \right> \left<n_e \right>^{-2}$ is a “clumping factor,” which accounts for inhomogeneities, $T_6$ is the temperature in units of $10^6$ K, and $\Lambda_{-23}$ is the radiative cooling rate of the gas in units of $10^{-23}$ ergs $^3$ $^{-1}.$ Near 1 keV $ \approx 10^7$ K, $\Lambda_{-23}$ is roughly a constant but in general it is a function of temperature and metallicity."
 For simplicity we adopt the Sutherland Dopita (1993) equilibriun values for this rate at a fixed. metallicity of 0.1... the level of pre-curichinent of Co dwarf stars in the solar neighborhood (Ostriker Thuan 1975) and in large neighborie galaxies (Thomas 1999).," For simplicity we adopt the Sutherland Dopita (1993) equilibrium values for this rate at a fixed metallicity of $0.1 Z_\odot,$ the level of pre-enrichment of G dwarf stars in the solar neighborhood (Ostriker Thuan 1975) and in large neighboring galaxies (Thomas 1999)."
 Finally in keeping with previous senianalvtical models. we ignore eas inhomoegencitics. taking C=1.," Finally in keeping with previous semianalytical models, we ignore gas inhomogeneities, taking $C=1$."
 Adopting a larecr value. consisteut with simulations. would result in the same overall trends. but with galaxies that were much too laree.," Adopting a larger value, consistent with simulations, would result in the same overall trends, but with galaxies that were much too large."
 We then compute the mass of cooled gas as a function of the total mass aud redshift of observation. as illustrated in Figure 1..," We then compute the mass of cooled gas as a function of the total mass and redshift of observation, as illustrated in Figure \ref{fig:cooled}."
 Tere we see the classic behavior first described in Rees Ostriker (1977) aud Silk (1977)., Here we see the classic behavior first described in Rees Ostriker (1977) and Silk (1977).
 As their cooling times are short. MecoΛέτο2On/Qy im small objects. while for larger masses with louger cooling times MisaMisOr/Og. Iu this picture the mania Mego(+) is the largest eas mass that has time to virialize and cool. a value that increases with time.," As their cooling times are short, $M_{\rm cool}/M_{\rm tot} \approx
\Omega_b/\Omega_0$ in small objects, while for larger masses with longer cooling times $M_{\rm cool}/M_{\rm tot} \ll \Omega_b/\Omega_0.$ In this picture the maximum $M_{\rm cool}(z)$ is the largest gas mass that has time to virialize and cool, a value that increases with time."
" Roughly one cau estimate this scaling as fe;444=tyx(11τε)3/72 whore f, and zy are the final time and redshift at which the gas cools while foo, is computed at the initial redshift at which the eas falls on the halo."," Roughly one can estimate this scaling as $t_{\rm cool,i}= t_f \propto (1+z_f)^{-3/2}$ where $t_f$ and $z_f$ are the final time and redshift at which the gas cools while $t_{\rm
cool,i}$ is computed at the initial redshift at which the gas falls on the halo."
" For the accretion histories described by ((1). La, depends only weakly ou redshift. which allows us to rewrite this as Me;X(11τε)BLT Tn Figure d. we also compute the average redshift of formation of the resulting stars. adoptiug both a nass average. that is za;=(fay) where andli a B-baud dhuniuositv-weiehted average. that is τειται) whereE) and[rt Yptt) is pythe tasB-baud mass to dieht ratio of a population of stars with an age f. assume a Salpeter initial mass function as computed by Bruzual Charlot (2003)."," For the accretion histories described by (1), $T_{\rm vir}$ depends only weakly on redshift, which allows us to rewrite this as $M_{\rm cool} \simpropto (1+z_f)^{-3/4} T_{\rm
vir}.$ In Figure \ref{fig:cooled} we also compute the average redshift of formation of the resulting stars, adopting both a mass average, that is ${\bar z}_M = z({\bar t}_M)$ where _M, and a $B$ -band luminosity-weighted average, that is ${\bar z}_{L_B} = z({\bar t}_{L_B})$ where, and $\Upsilon_B(t)$ is the $B$ -band mass to light ratio of a population of stars with an age $t,$ assuming a Salpeter initial mass function as computed by Bruzual Charlot (2003)."
" Di all cases tay and tp, are very close to the observed redshifts, iieaniung that iu objects of all masses aud redshifts. most of the stellar mass and hunuimositv cones from recently-formed stars."," In all cases ${\bar z}_M$ and ${\bar z}_{L_B}$ are very close to the observed redshifts, meaning that in objects of all masses and redshifts, most of the stellar mass and luminosity comes from recently-formed stars."
 We now inodifv our model to include. ACN feedback. adopting an approach that parallels Scannapieco Oh (2001). itself an extension of Writhe Loeb (2003).," We now modify our model to include AGN feedback, adopting an approach that parallels Scannapieco Oh (2004), itself an extension of Wyithe Loeb (2003)."
 These papers showed that nuposug an Mpgx(07 relationship between black-hole massaud halo circular velocity (Ferrarese 2002) aud assuming that these objects shine at their Eddinetouhuninosity for a fraction of the dynamical tine after cach major nierger gives aeood fit to the observed AGN luninosity function.," These papers showed that imposing an $M_{\rm BH} \propto
v_c^5$ relationship between black-hole massand halo circular velocity (Ferrarese 2002) and assuming that these objects shine at their Eddingtonluminosity for a fraction of the dynamical time after each major merger gives agood fit to the observed AGN luminosity function."
 Iu. particular. cach inerecr with a dass ratio less than bl was associated with an AGN with a total bolometric cnerey (in units of 1099 eres) of = saa 234057401Q4a," In particular, each merger with a mass ratio less than 4:1 was associated with an AGN with a total bolometric energy (in units of $10^{60}$ ergs) of = 23 (1+z)"
 Iu. particular. cach inerecr with a dass ratio less than bl was associated with an AGN with a total bolometric cnerey (in units of 1099 eres) of = saa 234057401Q4a|," In particular, each merger with a mass ratio less than 4:1 was associated with an AGN with a total bolometric energy (in units of $10^{60}$ ergs) of = 23 (1+z)"
 Iu. particular. cach inerecr with a dass ratio less than bl was associated with an AGN with a total bolometric cnerey (in units of 1099 eres) of = saa 234057401Q4a|u," In particular, each merger with a mass ratio less than 4:1 was associated with an AGN with a total bolometric energy (in units of $10^{60}$ ergs) of = 23 (1+z)"
 Iu. particular. cach inerecr with a dass ratio less than bl was associated with an AGN with a total bolometric cnerey (in units of 1099 eres) of = saa 234057401Q4a|uL," In particular, each merger with a mass ratio less than 4:1 was associated with an AGN with a total bolometric energy (in units of $10^{60}$ ergs) of = 23 (1+z)"
 Iu. particular. cach inerecr with a dass ratio less than bl was associated with an AGN with a total bolometric cnerey (in units of 1099 eres) of = saa 234057401Q4a|uLy," In particular, each merger with a mass ratio less than 4:1 was associated with an AGN with a total bolometric energy (in units of $10^{60}$ ergs) of = 23 (1+z)"
 Iu. particular. cach inerecr with a dass ratio less than bl was associated with an AGN with a total bolometric cnerey (in units of 1099 eres) of = saa 234057401Q4a|uLye," In particular, each merger with a mass ratio less than 4:1 was associated with an AGN with a total bolometric energy (in units of $10^{60}$ ergs) of = 23 (1+z)"
 Iu. particular. cach inerecr with a dass ratio less than bl was associated with an AGN with a total bolometric cnerey (in units of 1099 eres) of = saa 234057401Q4a|uLyes," In particular, each merger with a mass ratio less than 4:1 was associated with an AGN with a total bolometric energy (in units of $10^{60}$ ergs) of = 23 (1+z)"
that was evolving (changing shape. size. or location on the stellar surface).,"that was evolving (changing shape, size, or location on the stellar surface)."
 Llowever. the fact that so many |x. giants show lone-periocd variability at about the same timescales makes us reluctant to embrace the low-mass companion hypothesis.," However, the fact that so many K giants show long-period variability at about the same timescales makes us reluctant to embrace the low-mass companion hypothesis."
 1n a recent study of Xldebaran. Larson (1996) found no evidence of the 643-day. period in the equivalent: width measurements of the Ca LLA.. which is an indicator of chromospheric activity.," 	In a recent study of Aldebaran, Larson (1996) found no evidence of the 643-day period in the equivalent width measurements of the Ca II, which is an indicator of chromospheric activity."
 Εις casts some doubt on rotational moculation as a possible explanation for the RY variability., This casts some doubt on rotational modulation as a possible explanation for the RV variability.
 Llowever. the surface features on Aldebaran may be cilferent rom those which normally constitute stellar active regions (spots. plage. faculae) and as such these may show no modulation in chromospheric lines.," However, the surface features on Aldebaran may be different from those which normally constitute stellar active regions (spots, plage, faculae) and as such these may show no modulation in chromospheric lines."
 One thing is certain hough. for a surface feature. regardless of its nature. to ooduce RV. shifts in the stellar lines it would have to alter he line shapes.," One thing is certain though, for a surface feature, regardless of its nature, to produce RV shifts in the stellar lines it would have to alter the line shapes."
 Consequently the RY variations should be accompanied by spectral variability ancl the detection. of hese would. provide strong confirmation of the rotational modulation hypothesis., Consequently the RV variations should be accompanied by spectral variability and the detection of these would provide strong confirmation of the rotational modulation hypothesis.
 Of the current explanations for the wW variability of Aldebaran. a low-mass companion seems o be the only one which should. not. produce changes in he spectral line shapes.," Of the current explanations for the RV variability of Aldebaran, a low-mass companion seems to be the only one which should not produce changes in the spectral line shapes."
 Both surface features (e.g. Toner Gray 1988: Dempsey ct al., Both surface features (e.g. Toner Gray 1988; Dempsey et al.
 1992) ancl nonraclial pulsations (llatzes1900) should. produce spectral line variations., 1992) and nonradial pulsations \cite{hat96} should produce spectral line variations.
 A ack of such variability would provide more support for the planetary companion hypothesis., A lack of such variability would provide more support for the planetary companion hypothesis.
 Aldebaran is a slowly rotating star so the best method for cliscernine changes in the spectral line shapes is via line. bisectors., 	 	Aldebaran is a slowly rotating star so the best method for discerning changes in the spectral line shapes is via line bisectors.
 Disectors (the locus of the midpoints of horizontal line segments spanning the line width) have proved to be a powerful tool or studying subtle changes in the line shape due to the convection pattern on the star(Cirav 1982: Dravins LOST) or cool surface features (Loner Cray 1988: Dempsey et al., Bisectors (the locus of the midpoints of horizontal line segments spanning the line width) have proved to be a powerful tool for studying subtle changes in the line shape due to the convection pattern on the star(Gray 1982; Dravins 1987) or cool surface features (Toner Gray 1988; Dempsey et al.
 1992)., 1992).
 Recenthy. a study of the spectral line bisectors of 51 Pee was recently used to cast serious doubt on the validity of the planet hypothesis for this star (Gray 1997).," Recently, a study of the spectral line bisectors of 51 Peg was recently used to cast serious doubt on the validity of the planet hypothesis for this star (Gray 1997)."
 The objectives ofthis paper are simply to search for line bisector variability in Aldbaran., 	 	The objectives of this paper are simply to search for line bisector variability in Aldbaran.
 Such variability with the RY period would. immediately eliminate the planet hypothesis as a cause of this period., Such variability with the RV period would immediately eliminate the planet hypothesis as a cause of this period.
 A lack of bisector variability. would exclude rotational modulation ancl possibly pulsations as Viable hypotheses., A lack of bisector variability would exclude rotational modulation and possibly pulsations as viable hypotheses.
 The organisation of this paper is as follows., The organisation of this paper is as follows.
 The data are described in Section 822 and the analysis of the spectral line bisectors is presented in Section 833 and it is demonstrated that bisector variability is indeed present. but at a period unrelated to the RV. period.," The data are described in Section 2 and the analysis of the spectral line bisectors is presented in Section 3 and it is demonstrated that bisector variability is indeed present, but at a period unrelated to the RV period."
 In Section £44 the nature ofthe bisector variability is examined., In Section 4 the nature of the bisector variability is examined.
 Finally. in. Section. 855 the implications. of the bisector variability on the interpretation of the 643-day. period is in examined.," Finally, in Section 5 the implications of the bisector variability on the interpretation of the 643-day period is in examined."
 For the past. several vears at AleDonald Observatory we have been using precise stellar racial velocity measurements (σ = 1020 ) to study a variety. of phenomena in stars. including Ix giant. variability.," 	For the past several years at McDonald Observatory we have been using precise stellar radial velocity measurements $\sigma$ = 10–20 ) to study a variety of phenomena in stars, including K giant variability."
 In the carly vears of this programme the telluric Ου lines were used as a wavelength reference for measuring stellar radial velocity shifts., In the early years of this programme the telluric $_2$ lines were used as a wavelength reference for measuring stellar radial velocity shifts.
 This technique was first proposed bv the Grillin Cirillin (1973) and we have confirmed their claim that it is capable of measuring stellar radial velocities to a precision of, This technique was first proposed by the Griffin Griffin (1973) and we have confirmed their claim that it is capable of measuring stellar radial velocities to a precision of.
 These data are an ideal set for searching [or spectra variability as they were taken at high spectral. resolution. with high signal-to-noise ratio. and only a few of the spectra lines are blended with telluric Oo.," These data are an ideal set for searching for spectral variability as they were taken at high spectral resolution, with high signal-to-noise ratio, and only a few of the spectral lines are blended with telluric $_2$."
 Phere is the ade advantage that the stellar radial velocity. ancl line shape measurements were mace from the same data set., There is the added advantage that the stellar radial velocity and line shape measurements were made from the same data set.
 The observations were acquired. at MeDonak Observatory using the coudé focus of the 2.7-m telescope., 	The observations were acquired at McDonald Observatory using the coudé focus of the 2.7-m telescope.
 An cchelle erating was used in single pass along with a ‘Texas Instruments 500 3-phase CCD., An echelle grating was used in single pass along with a Texas Instruments $\times$ 800 3-phase CCD.
 The wavelength coverage was centred on at a spectral resolution of (130 jan slit which subtends 2 pixels on the CCD)., The wavelength coverage was centred on at a spectral resolution of (130 $\mu$ m slit which subtends 2 pixels on the CCD).
 An interference. filter was used. to. isolate the appropriate order from the cchelle grating., An interference filter was used to isolate the appropriate order from the echelle grating.
 Figure 1 shows a typical spectrum of Aldebaran taken with this setup., Figure \ref{fig:spectrum} shows a typical spectrum of Aldebaran taken with this setup.
 The horizontal lines mark the location of the telluric Os lines. (, The horizontal lines mark the location of the telluric $_2$ lines. (
Lhe wavelength scale of this spectrum has been corrected for both the Earth's orbital motion and the overall racial velocity motion of the star.,The wavelength scale of this spectrum has been corrected for both the Earth's orbital motion and the overall radial velocity motion of the star.
 Although this places the stellar lines at the true rest wavelength. the telluric Os lines appear shifted from their proper wavelength position.)," Although this places the stellar lines at the true rest wavelength, the telluric $_2$ lines appear shifted from their proper wavelength position.)"
 Typical signal-to-noise ratio per resolution clement for a given exposure is about 450., Typical signal-to-noise ratio per resolution element for a given exposure is about 450.
 For line bisector work one would like to use a relatively deep. unblencec tine.," 	For line bisector work one would like to use a relatively deep, unblended line."
 Unfortunately. Aldebaran: is a cool," Unfortunately, Aldebaran is a cool"
"Assuming that the time of peak [lux density can be approximated by travel time down the jet. this occurs at a distance 2,—1ο downstream. at which point the jet radius is 2),=z,1anó (assuming a cone opening angle of €)àoy— iin all cases).","Assuming that the time of peak flux density can be approximated by travel time down the jet, this occurs at a distance $z_{\rm p}
=\gamma\beta\delta ct_{\rm p}$ downstream, at which point the jet radius is $R_{\rm p}=z_{\rm p}\tan\phi$ (assuming a cone opening angle of $\phi=2$ in all cases)."
 The source angular size is given by 2h fd. where d is the source distance (takentobe10kpefor(νοX-3:Dickey 1983)..," The source angular size is given by $\theta_{\rm src} = 2R_{\rm p}/d$ , where $d$ is the source distance \citep[taken to be 10\,kpc for Cyg X-3;][]{Dic83}. ."
" Given the angular size 8,4. turnover [requencey. ου. and Bux density. ον. and the electron index 5. with the assumption of homogeneous spherical source with a power-law electron. energy.a spectrum. then the magnetic Ποια D. and the normalisation of the electron. spectrum. αν can be calculated at the turnover using the equations of Marseher(1987)... as previously done for a sample of Galactic and AGN jets by TFürler&Lindfors(2007)."," Given the angular size $\theta_{\rm src}$, turnover frequency, $\nu_{\rm p}$, and flux density, $S_{\rm p}$, and the electron index $s$, with the assumption of a homogeneous spherical source with a power-law electron energy spectrum, then the magnetic field $B$, and the normalisation of the electron spectrum, $\kappa$, can be calculated at the turnover using the equations of \citet{Mar87}, as previously done for a sample of Galactic and AGN jets by \citet{Tur07}."
. While or a shock-in-jet model the source has a evlindrical slab-like geometry rather than being spherical. the depth scales with I during the self-similar adiabatic phase. so approximating he source as a sphere will be in error by the ratio of the slab thickness to the jetradius 2. a [actor £~O.005h7 (asderivedforthecaseofBC273:AMarscher&Cear1985). where f=Lfy)/(100kmsAlpe1p and dy is the Llubble constant.," While for a shock-in-jet model the source has a cylindrical slab-like geometry rather than being spherical, the depth scales with $R$ during the self-similar adiabatic phase, so approximating the source as a sphere will be in error by the ratio of the slab thickness $x$ to the jetradius $R$, a factor $\xi\sim0.005h^{-5}$ \citep[as derived for the case of 3C\,273;
][]{Mar85}, where $h=H_0/(100{\rm \,km\,s}^{-1}{\rm \,Mpc}^{-1})$ and $H_0$ is the Hubble constant."
 For the currenthy-acceptec value of ff=0.732 (Spergelctal.2007)... £~ 0.024. M," For the currently-accepted value of $h=0.732$ \citep{Spe07}, $\xi\sim0.024$ ."
arscher&Cear(1985) also approximated the density of the shocked: region as »ng Constant up to a distance .=£2 behind the shock ront. at which point it would abrupthy drop to zero.," \citet{Mar85} also approximated the density of the shocked region as being constant up to a distance $x=\xi R$ behind the shock front, at which point it would abruptly drop to zero."
 They ound that this should not introduce significant errors into he proportionalities for their equations., They found that this should not introduce significant errors into the proportionalities for their equations.
 The assumption of homogeneity of the emitting region was further justified » Bjdrnsson&Aslaksen(2000).. and. should. certainly be sullicient to derive the scalings between physical parameters which we now present.," The assumption of homogeneity of the emitting region was further justified by \citet{Bjo00}, and should certainly be sufficient to derive the scalings between physical parameters which we now present."
" Adapting the equations of Marscher(1987) to account for the non-spherical gcometry. ancl assuming we see the emitting region with the shock face-on (Lindforsetal. 2007).. the magnetic field is given by the electron normalisation is andthe optical depth to svnchrotron self-absorption at the turnover frequency is where the source size Gayo is given in mas. the peak emission frequency f, in. Cllz. the peak Dux density ορ in Jv and the source distance d in kpe."," Adapting the equations of \citet{Mar87} to account for the non-spherical geometry, and assuming we see the emitting region with the shock face-on \citep{Lin07}, , the magnetic field is given by the electron normalisation is andthe optical depth to synchrotron self-absorption at the turnover frequency is where the source size $\theta_{\rm src}$ is given in mas, the peak emission frequency $\nu_{\rm p}$ in GHz, the peak flux density $S_{\rm
p}$ in Jy and the source distance $d$ in kpc."
 a=(s1)/2 is the spectral index. and (ea) and es(o) ave functions depending ogarithmically on a. with values of 0.27 and 1.2«Lotrespectively for a=0.5.," $\alpha=(s-1)/2$ is the spectral index, and $n(\alpha)$ and $c_2(\alpha)$ are functions depending logarithmically on $\alpha$, with values of 0.27 and $1.2\times10^{17}$respectively for $\alpha=0.5$."
 ία) is a slowly-varving function of a with a value of 3.2 for a=0.5., $b(\alpha)$ is a slowly-varying function of $\alpha$ with a value of 3.2 for $\alpha=0.5$.
 Away [rom the peak. he optical depth scales withfrequency às 7x£tle|Dia," Away from the peak, the optical depth scales withfrequency as $\tau \propto \nu^{-(2\alpha+5)/2}$."
 To meaninefully compare the physical parameters. for cillerent outbursts. they should. be scaled. to the same »osition 2 along the jet tto a fixed jet radius 2).," To meaningfully compare the physical parameters for different outbursts, they should be scaled to the same position $z$ along the jet to a fixed jet radius $R$ )."
 The magnetic field. and. electron normalisation scale as power aws with jet radius. Box4? and &xP.," The magnetic field and electron normalisation scale as power laws with jet radius, $B\propto
R^{-b}$ and $\kappa\propto R^{-k}$."
 b is given in ‘Table 1. and &=2(s|2)/3 for an adiabatic jet low as assumed here.," $b$ is given in Table \ref{tab:model}, and $k=2(s+2)/3$ for an adiabatic jet flow as assumed here."
" From the peak Uuxes. frequencies and times of the individual outbursts.the magnetic field. electron. normalisation. and. optical depth to synchrotron sell-absorption at a frequency of GOLIz were caleulated. scaled to the same jet radius of LAA using the derived values of b and &. and plotted against peak lux density. ορ. [requency. £. and time. /, in Fig. 9.."," From the peak fluxes, frequencies and times of the individual outbursts,the magnetic field, electron normalisation, and optical depth to synchrotron self-absorption at a frequency of GHz were calculated, scaled to the same jet radius of AU using the derived values of $b$ and $k$, and plotted against peak flux density, $S_{\rm p}$, frequency, $\nu_{\rm p}$, and time, $t_{\rm p}$ in Fig. \ref{fig:physparms}."
 In all cases we assumed adistance of I0kkpe (Dickey 1983)... a jet speed 7=0.63. an inclination angle of tto the line of sight and a jet hall-opening angle of 2(Miller- 2004).," In all cases we assumed adistance of kpc \citep{Dic83}, , a jet speed $\beta=0.63$, an inclination angle of to the line of sight and a jet half-opening angle of \citep{Mil04}."
. While we note that slightly. cülferent parameters were found by Mioduszewskietal.(2001)... we use a single set of parameters here for consistency. between outbursts.," While we note that slightly different parameters were found by \citet{Mio01}, we use a single set of parameters here for consistency between outbursts."
 There is clearly significantscatter in the individual events (due to thestrong dependences of D. and & on Sy. pO and 8.4).," There is clearly significantscatter in the individual events (due to thestrong dependences of $B$ and $\kappa$ on$S_{\rm p}$ , $\nu_{\rm
p}$ and $\theta_{\rm src}$ )."
" lHlowever. we can use equations 4.. and 6 together with the derived. sealings of Si. 1%, and /, with/one another. given in Table 3. to show the expected trends in D. & and 7."," However, we can use equations \ref{eq:magfield}, , \ref{eq:enorm} and \ref{eq:selfabs} together with the derived scalings of $S_{\rm p}$ , $\nu_{\rm p}$ and $t_{\rm p}$ withone another, given in Table \ref{tab:scalings}, , to show the expected trends in $B$ , $\kappa$ and $\tau$ ."
 These are plotted as dashed lines in Fig. 9.., These are plotted as dashed lines in Fig. \ref{fig:physparms}. .
 Spearman rank correlations show that, Spearman rank correlations show that
outer part. of the nebula is reduced. in the outer parts of the nebula in the Ritzervelcl approximation. primarily as a result of the lower ionization in the core.,"outer part of the nebula is reduced in the outer parts of the nebula in the Ritzerveld approximation, primarily as a result of the lower ionization in the core."
 For the case of i7 the assumption that. =1 within the Strómmegren sphere means that this approximation also does not capture the broadening of the ionization front.," For the case of $r^{-2}$, the assumption that $x=1$ within the Strömmgren sphere means that this approximation also does not capture the broadening of the ionization front."
 In Figure 3.. we show the angular distribution of the diffuse radiation αἲ various positions within the nebula.," In Figure \ref{f:polar}, we show the angular distribution of the diffuse radiation at various positions within the nebula."
 Where the absorption coefficients. of direct. anc diffuse radiation are the same. the cliffuse radiation Lickel is also beamed strongly. forward. through much. of the nebula. because it is principally emitted in the dense gas in the core.," Where the absorption coefficients of direct and diffuse radiation are the same, the diffuse radiation field is also beamed strongly forward through much of the nebula, because it is principally emitted in the dense gas in the core."
 llence. although the radiation energv may be distributed with the form characteristic of recombination emission. elfects of a directed raciation field such as shadowing should still be in evidence.," Hence, although the radiation energy may be distributed with the form characteristic of recombination emission, effects of a directed radiation field such as shadowing should still be in evidence."
 For the uniform densitv case with diffuse absorption coelIicient six time larger than direct. the radiation Licld is most strongly beamed at intermediate raclii within the nebula: the directionality is small at r=0.1pc for both uniform density cases as this is close to the centre of svmmetry. where as for the higher absorption case the clilfuse field is also symmetric near to the edge of the nebula because the minI surface samples a relatively small degree of anisotropy in the direct. field.," For the uniform density case with diffuse absorption coefficient six time larger than direct, the radiation field is most strongly beamed at intermediate radii within the nebula: the directionality is small at $r=0.1{\rm\,pc}$ for both uniform density cases as this is close to the centre of symmetry, where as for the higher absorption case the diffuse field is also symmetric near to the edge of the nebula because the $\tau_{\rm dif} \sim 1$ surface samples a relatively small degree of anisotropy in the direct field."
 For the innermost radius of the r.7 distribution with the higher absorption coelIicient for the direct Field. the inner edge of the photoionized gas can be clearly seen in the diffuse raciation field.," For the innermost radius of the $r^{-2}$ distribution with the higher absorption coefficient for the direct field, the inner edge of the photoionized gas can be clearly seen in the diffuse radiation field."
 In Figure 4.. we plot the dilfuse radiation intensity as à fraction of the direct radiation intensity as a function. of radius for our standard set of parameters.," In Figure \ref{f:radang}, we plot the diffuse radiation intensity as a fraction of the direct radiation intensity as a function of radius for our standard set of parameters."
 These results complement the polar diagrams at specific radii shown in Figure 3..., These results complement the polar diagrams at specific radii shown in Figure \ref{f:polar}.
 As we have seen. in the case of equal absorption coellicients. the dilluse intensity can be higher than the direct intensity.," As we have seen, in the case of equal absorption coefficients, the diffuse intensity can be higher than the direct intensity."
 However the ratio of the cülfuse [ux to the direct flux. relevant to the illumination of the tails of cometary globules. reaches at most. GO per cent at the edge of ar7 density distribution. and through most of the nebula is closer to the 15 per cent figure derived by (1998).," However the ratio of the diffuse flux to the direct flux, relevant to the illumination of the tails of cometary globules, reaches at most 60 per cent at the edge of a $r^{-2}$ density distribution, and through most of the nebula is closer to the 15 per cent figure derived by \cite{canto98}."
. In the case where the direct [ux is less strongly absorbed. all the dilfuse [ux components are close to 2.5 per cent of the direct Dux through almost all of the region: the outward-going [ux shows the most structure. but only reaches at most around 5.5 per cent of the direct Dux at its maximuni.," In the case where the direct flux is less strongly absorbed, all the diffuse flux components are close to 2.5 per cent of the direct flux through almost all of the region; the outward-going flux shows the most structure, but only reaches at most around 5.5 per cent of the direct flux at its maximum."
 This plot also shows the ratio of the diffuse to direct field in the OTS anc Ritzerveld approximations., This plot also shows the ratio of the diffuse to direct field in the OTS and Ritzerveld approximations.
 lt is clear that the OTS approximation predicts a reasonable average value for this ratio. but overestimates the relative importance of the dilfuse field in the core of the nebula and underestimates it at the. cdee of the nebula.," It is clear that the OTS approximation predicts a reasonable average value for this ratio, but overestimates the relative importance of the diffuse field in the core of the nebula and underestimates it at the edge of the nebula."
 la contrast. Ritzerveld’s formalism. unclerestimates relative importance of the dilluse field in the core but overestimates it at the edge of the nebula. particularly for the case of equal absorption coellicients ancl steeply decreasing density distibutions (as we saw in Figure 2.. this is primarily due to uncerestimating the direct field here).," In contrast, Ritzerveld's formalism underestimates relative importance of the diffuse field in the core but overestimates it at the edge of the nebula, particularly for the case of equal absorption coefficients and steeply decreasing density distibutions (as we saw in Figure \ref{f:radiat}, this is primarily due to underestimating the direct field here)."
 In both cases. the first approximation would be that the lateral dilfuse field was one quarter the radiation intensity. which would be a significant overestimate for the case of equal absorption coellicients.," In both cases, the first approximation would be that the lateral diffuse field was one quarter the radiation intensity, which would be a significant overestimate for the case of equal absorption coefficients."
 In Figure 5.. we show the same results as previously for a thin shell of uniform density between 0.9755 and rs.," In Figure \ref{f:shell}, we show the same results as previously for a thin shell of uniform density between $0.9r_{\rm S}$ and $r_{\rm S}$."
 In this case. Ritzerveld’s approximation obviously underestimates re clilfuse field dramatically in the empty core of the nebula. as there are no recombinations in this region. while 16 OTS approximation overestimates the dilfuse Lux.," In this case, Ritzerveld's approximation obviously underestimates the diffuse field dramatically in the empty core of the nebula, as there are no recombinations in this region, while the OTS approximation overestimates the diffuse flux."
 The radiation fick is svmumetric in angle at all our sample »»ints. which are at radii less than the inner edge of the gaell: this is to be expected. as the angle to the surface rom oppositely directed beams is the same and the surface xiehtness is constant.," The radiation field is symmetric in angle at all our sample points, which are at radii less than the inner edge of the shell: this is to be expected, as the angle to the surface from oppositely directed beams is the same and the surface brightness is constant."
 Particularly in the case with higher irect absorption coellicient. limb-brightening is seen for the sample point at the inner surface of the shell. corresponding o the higher direct [ux (and hence source function) at the TagΞ1 surface in this direction.," Particularly in the case with higher direct absorption coefficient, limb-brightening is seen for the sample point at the inner surface of the shell, corresponding to the higher direct flux (and hence source function) at the $\tau_{\rm dif}=1$ surface in this direction."
 In Ligure 6.. we compare the accuracy of a dilfusion approximation solution to the dilfuse field intensity to that of the on-the-spot) approximation.," In Figure \ref{f:accuracy}, we compare the accuracy of a diffusion approximation solution to the diffuse field intensity to that of the on-the-spot approximation."
 The diffusion results are almost evervwhere more accurate than the OTS approximation. being in error by less than 30 per cent even in the most difficult. central region of the Low.," The diffusion results are almost everywhere more accurate than the OTS approximation, being in error by less than 30 per cent even in the most difficult, central region of the flow."
 Both schemes are accurate at the edge of the nebula. where the most interesting dynamical structures occur. although the diffusion results are still significantly more precise.," Both schemes are accurate at the edge of the nebula, where the most interesting dynamical structures occur, although the diffusion results are still significantly more precise."
 A diffusion approximation solver could. be embedded in à dynamical code by updating the ionization. solution consistently with the sources derived. from. the. cülfusion solver and iterating. in a similar fashion to that which we use for the full field transfer.," A diffusion approximation solver could be embedded in a dynamical code by updating the ionization solution consistently with the sources derived from the diffusion solver and iterating, in a similar fashion to that which we use for the full field transfer."
 Experimenting with updating he ionization structure after. both inward. ancl outward ransfer are complete suggests that the solution converges more slowly than if the ionization structure is updated on each sweep., Experimenting with updating the ionization structure after both inward and outward transfer are complete suggests that the solution converges more slowly than if the ionization structure is updated on each sweep.
 However. in the context of a dynamical scheme. he initial approximation for the ionization structure will jwe come from the previous timestep. and hence would be expected to not be far from the correct. solution.," However, in the context of a dynamical scheme, the initial approximation for the ionization structure will have come from the previous timestep, and hence would be expected to not be far from the correct solution."
 The dilfusion approximation allows processes such as he lateral illumination of the tails of photoionized globules o be modelled (Cantóctal.1998:Pavlakiset2001).," The diffusion approximation allows processes such as the lateral illumination of the tails of photoionized globules to be modelled \citep{canto98,pavla01}."
. However. the lateral illumination of shadowed tails may be overstimated in cases where the beaming of the dilfuse field is expected to be strong (Alellemaetal.2006):: Eddington tensor approaches (e.g.González&Audit2005)— may improve the results in such limits.," However, the lateral illumination of shadowed tails may be overstimated in cases where the beaming of the diffuse field is expected to be strong \citep{melle06}: Eddington tensor approaches \citep[e.g.\][]{gonza05} may improve the results in such limits."
 The accuracy of the OTS approximation may be evaluated by comparing the scale of variation of the diffuse field source function. 44. to the mean free path for diffuse photons. orjeay.," The accuracy of the OTS approximation may be evaluated by comparing the scale of variation of the diffuse field source function, $H$, to the mean free path for diffuse photons, $\lambda = 1/n(1-x)a_1$ ."
 Using equation (29)) for the source function. ancl the OTS approximation equation (7)) for the ionization balance Note that we can still use the OTS approximation to close our equations at this higher order.," Using equation \ref{e:sdif}) ) for the source function, and the OTS approximation equation \ref{e:casebbalance}) ) for the ionization balance Note that we can still use the OTS approximation to close our equations at this higher order."
 The source function, The source function
During the course of the observations. there were some data sets obtained. uuder couclitions.,"During the course of the observations, there were some data sets obtained under less-than-ideal conditions."
 Although VLA observaious at 8.2 GHz are relatively unallected by poor weather. extreme weather conditios such as thuderstorums or very high: winds can produce data which cau not be properly calibrated eiven the observing strategy that we hac used.," Although VLA observations at 8.5 GHz are relatively unaffected by poor weather, extreme weather conditions such as thunderstorms or very high winds can produce data which can not be properly calibrated given the observing strategy that we had used."
 Based on our evaluation of the success of the calib‘ation procedre. we chose to delete a sinall number of epochs [roin the light curves.," Based on our evaluation of the success of the calibration procedure, we chose to delete a small number of epochs from the light curves."
 Paper I describes tlie reasons or ageingee two bad epocis ln season 1., Paper I describes the reasons for flagging two bad epochs in season 1.
 In seasous 2 aud d3 we flaggedOUO a total of LeΡος» due to bac| weather conditions., In seasons 2 and 3 we flagged a total of 4 epochs due to bad weather conditions.
 We also ΠασσοςMD 3 points obtained at the very begiuniug of he seasons durNEMij which the VLA was iua mixed coufiguratiou while the antennas were beiug moved [rom their D-array positious to their A-array positions., We also flagged 3 points obtained at the very beginning of the seasons during which the VLA was in a mixed configuration while the antennas were being moved from their D-array positions to their A-array positions.
 For these epochs. the absolute flux deusity calibratio Lwas not correct.," For these epochs, the absolute flux density calibration was not correct."
 The ntiiber of epochs remaining after [IageingMD the data from each season is Liste in the [N44 column of Table 1.., The number of epochs remaining after flagging the data from each season is listed in the $N_{\rm good}$ column of Table \ref{tab_obs}.
 Paper Laud FTOL describe a method that uses tie light curves of steep-spectrum radio sources to correct for systematic errors in the absolute flux deusity calibration., Paper I and FT01 describe a method that uses the light curves of steep-spectrum radio sources to correct for systematic errors in the absolute flux density calibration.
 We have followed the same method in this paper to correct tie. light. curves of the leused images., We have followed the same method in this paper to correct the light curves of the lensed images.
 The results preseuted in FTOI showed that CSOs are excellent flux. densiy Ccσαibrators at the frequeucies used for our monitoring οservatiolis., The results presented in FT01 showed that CSOs are excellent flux density calibrators at the frequencies used for our monitoring observations.
" As discussed in ΕΤΟΙ he lieht curves of J18234-7938 aud J19125--70J55 sliowed evidence of additional systemaic ellects uo ""eel in the light curves of the other CSOs."," As discussed in FT01, the light curves of J1823+7938 and J1945+7055 showed evidence of additional systematic effects not seen in the light curves of the other CSOs."
 Thus. those wo sources were not incltded in Ιthe «'onstruction of the season 3 correction curve.," Thus, those two sources were not included in the construction of the season 3 correction curve."
 Acelitioually. the seasoi 2 data show that 1633+71) altiough relatively stable. has uoisier light curves than hose of tle CSOs.," Additionally, the season 2 data show that 1633+741, although relatively stable, has noisier light curves than those of the CSOs."
 This is probaM7 ue to the siguificant extended. emission aud complicated uorphology of 16334711., This is probably due to the significant extended emission and complicated morphology of 1633+741.
 Our suapshotp observatious did not cover the (rv) plane well enough to accuratev model the complex erMission of the source.," Our snapshot observations did not cover the $(u,v)$ plane well enough to accurately model the complex emission of the source."
 Thus. we dropped 16334711 as a secondary flux deusity calibrator.," Thus, we dropped 1633+741 as a secondary flux density calibrator."
 There were two clealls to cousider in consructing the correction curves., There were two details to consider in constructing the correction curves.
 The first was that the oximary flux calibrator 3313) has some exeuded emission.," The first was that the primary flux calibrator, 343 has some extended emission."
 The B-coufiguration—— observations are more sensitive to lover suTace brightuess emission than are the A-co[umeuration observations., The B-configuration observations are more sensitive to lower surface brightness emission than are the A-configuration observations.
 Thus. the total flux desity observed [or 331κ)+) will increase wlienever ie VLA enters the B confieuration.," Thus, the total flux density observed for 343 will increase whenever the VLA enters the B configuration."
" The obse""vatlols of the secondary. Mus calibrators were use( Oo examine the size of his elfect.", The observations of the secondary flux calibrators were used to examine the size of this effect.
 The effect only appeared to be significant during season 3. perhapsReS because the brielituess oftle central coniponen had «ecreased slightly aud the extended emission tius contributed a larger yercentage to the total B-couiguration [lux density.," The effect only appeared to be significant during season 3, perhaps because the brightness of the central component had decreased slightly and the extended emission thus contributed a larger percentage to the total B-configuration flux density."
 To eusure that the correction curves did uot include any step fuuctioi introduced by the changiig Παν deusity of 3313. we computed for each season the mean flux cdeusities of the CSOs iu tle A. Buna. and B conligurations separately aud sed those values to LOLualize the light curves.," To ensure that the correction curves did not include any step function introduced by the changing flux density of 343, we computed for each season the mean flux densities of the CSOs in the A, BnA, and B configurations separately and used those values to normalize the light curves."
 This procedure also corrected [or any step functions, This procedure also corrected for any step functions
galaxies since thev are likely the building blocks of more massive svstems like the Milky Wav.,galaxies since they are likely the building blocks of more massive systems like the Milky Way.
 Because the basic properties of dwarf galaxies (luminosity. stellar ancl gaseous mass. star formation rate (SER). and metallicity) are very different [rom those of massive galaxies. detailed studies of the former objects provide us with an important (ool to understaiud better how the astrophyvsics of star formation and galaxy evolution depend on (hese characteristics.," Because the basic properties of dwarf galaxies (luminosity, stellar and gaseous mass, star formation rate (SFR), and metallicity) are very different from those of massive galaxies, detailed studies of the former objects provide us with an important tool to understand better how the astrophysics of star formation and galaxy evolution depend on these characteristics."
 The dwarl galaxies of the Local Group. in particular. provide a wealth of information about their evolution: unlike more distant objects. we are able (ο resolve them into individual stars to study Che characteristics of their stellar populations and to create detailed star formation histories (SFIIs).," The dwarf galaxies of the Local Group, in particular, provide a wealth of information about their evolution; unlike more distant objects, we are able to resolve them into individual stars to study the characteristics of their stellar populations and to create detailed star formation histories (SFHs)."
 The Local Group dwarl irregular (dI) ealaxy WEM. first discovered bx. Wolf(19090). and independently rediscovered by Landmark aud Melotte (Melotte1926).. has been the subject of many investigations over (le past (vodecades (6L.vandenBereh2000).," The Local Group dwarf irregular (dI) galaxy WLM, first discovered by \citet{wol09} and independently rediscovered by Lundmark and Melotte \citep{mel26}, has been the subject of many investigations over the past twodecades \citep[cf.,][]{van00}."
. al.(1981). and IIuchimeier&Richter(1986). first addressed the pproperties of the galaxy. finding it to be gas-rich. in keeping with its dl morphology.," \citet{huc81} and \citet{huc86} first addressed the properties of the galaxy, finding it to be gas-rich, in keeping with its dI morphology."
 were the first to resolve the galaxy in. and found a double-peaked fIux distribution.," \citet{jac04} were the first to resolve the galaxy in, and found a double-peaked flux distribution."
 Because WLAI is relatively nearby 2000).. the stellar population is easily resolved. enabling high quality optical photometry.," Because WLM is relatively nearby \citep[m$-$M~=~24.88;][]{dol00}, the stellar population is easily resolved, enabling high quality optical photometry."
 Dolphin(2000) and Rejkubaetal.(2000) find WLMs SFI suggests the bulk of the stus formed over 9 Gyr ago., \citet{dol00} and \citet{rej00} find WLM's SFH suggests the bulk of the stars formed over 9 Gyr ago.
 This episode of star formation was followed by a gradual decrease in the star formation rate until the most recent event that began between 1 and 2.5 Gyr ago and continues today (ILodge&Miller1995:Skillmanetal.1989:IIunterοἱ1993:al. 2005).," This episode of star formation was followed by a gradual decrease in the star formation rate until the most recent event that began between 1 and 2.5 Gyr ago and continues today \citep{hod95,ski89,hun93,lee05}."
. Jacksonetal.(2006). studied the emission from hot dust aud polvevelic aromatic hvdroearbons (PAlIs) and found only very low surface brightness emission coincident with the highest surface brightness rreeions. as is expected Irom a galaxy. with relatively low metallicity (12+log(O/11) and current star formation rate (0.003M.vr.|:1998).," \citet{jac06}
 studied the emission from hot dust and polycyclic aromatic hydrocarbons (PAHs) and found only very low surface brightness emission coincident with the highest surface brightness regions, as is expected from a galaxy with relatively low metallicity \citep[12~+~log(O/H)~=~7.83;][]{lee05} and current star formation rate \citep[0.003 M$_\sun$ yr$^{-1}$."
 The evolved stellar population of. WLAL was examined by Cooketal.(1986) and Datünelli&Demers(2004).. who both found an extremely. high. carbon star to M-tvpe asvinplolic giant branch (AGB) star ratio in accordance with its low metallicity 1983).," The evolved stellar population of WLM was examined by \citet{coo86}
 and \citet{bat04}, who both found an extremely high carbon star to M-type asymptotic giant branch (AGB) star ratio in accordance with its low metallicity \citep{ibe83}."
. Other studies using Ποτονρα filters centered on the CN and TiO absorption features have enabled the study of AGBs in a limited sample of Local Group galaxies (Cook2001:Dattünelli&Demers 2005).," Other studies using narrow-band filters centered on the CN and TiO absorption features have enabled the study of AGBs in a limited sample of Local Group galaxies \citep{coo86,now01,bat05}."
. Because optical studies can sulLer significantly from extinction effects (due to both the host galaxy interstellar medium (ISM) and the winds of individual stars). many AGB studies have been performed in the infrared (I(e.g..Doveretal.2006:Zijlstraοἱ1996: 2005)..," Because optical studies can suffer significantly from extinction effects (due to both the host galaxy interstellar medium (ISM) and the winds of individual stars), many AGB studies have been performed in the infrared \citep[e.g.,][]{boy06,zij96,woo98,van99,van05}. ."
 The majority of this effort has been limited to the Galaxy and the Magellanie Clouds, The majority of this effort has been limited to the Galaxy and the Magellanic Clouds
"Gv. where v, is the outflow velocity and we calculate the recombination temperature {νι from the density and entropy. assuming nuclear statistical equilibrium) (eg. QW96: their eq. [",", where $v_{\rm r}$ is the outflow velocity and we calculate the recombination temperature $T_{\rm rec}$ from the density and entropy, assuming nuclear statistical equilibrium (e.g. QW96; their eq. ["
62]).,62]).
" We calculate the velocity using mass continuity 44.=py,A, where A=Azfosar/RUP is the areal function of the dipolar flux tube (recombination generally occurs interior to the light cylinder) and. fopen*AyΛΙ is the fraction of the magnetosphere at the surface open to outflows."," We calculate the velocity using mass continuity $\dot{M} = \rho v_{r}A$ , where $A = 4\pi R_{\rm ns}^{2}f_{\rm open}(r/R_{\rm ns})^{3}$ is the areal function of the dipolar flux tube (recombination generally occurs interior to the light cylinder) and $f_{\rm open} \approx R_{\rm ns}/2R_{\rm L}$ is the fraction of the magnetosphere at the surface open to outflows."
 The entropy S(/) and expansion timescale Τουρ) for the winc solution in Figure |. are shown with dashed and dot-dashed lines. respectively.," The entropy $S(t)$ and expansion timescale $\tau_{\rm exp}(t)$ for the wind solution in Figure \ref{fig:wind} are shown with dashed and dot-dashed lines, respectively."
 Note that S is approximately constant in time because he rising value of S(Q=0)«£L.'?e;Rs: (eq. [I]]), Note that $S$ is approximately constant in time because the rising value of $S(\Omega = 0) \propto L_{\bar{\nu}_{e}}^{-1/6}\epsilon_{\bar{\nu}_{e}}^{-1/3}R_{\rm ns}^{-2/3}$ (eq. \ref{eq:SNRNM}] ])
 is otfse by an increasein the exponential factor PoyfP (eg. [ο]., is offset by an increasein the exponential factor $P_{\rm cent}/P$ (eq. \ref{eq:Somega}] ]).
 Also note that at late times (>100 s) the expansion timescale Τομ becomes comparable to the timescale over which the properties of he wind are changing. suggesting that the steady-state assumption we have adopted may break down.," Also note that at late times $\gtrsim 100$ s) the expansion timescale $\tau_{\rm exp}$ becomes comparable to the timescale over which the properties of the wind are changing, suggesting that the steady-state assumption we have adopted may break down."
 Although this does not atfec he conclusions of this paper because we are focused on the je ?roperties on timescales of tens of seconds. future numerical work is required to more accurately address the properties of magnetizec oroto-NS winds at late times and low neutrino luminosities.," Although this does not affect the conclusions of this paper because we are focused on the jet properties on timescales of tens of seconds, future numerical work is required to more accurately address the properties of magnetized proto-NS winds at late times and low neutrino luminosities."
 The electron fraction Y. is important for two reasons., The electron fraction $Y_{e}$ is important for two reasons.
" First. he value of Y, determines the channel by which helium burns to ‘orm carbon."," First, the value of $Y_{e}$ determines the channel by which helium burns to form carbon."
" Since this is the slowest reaction. Y, impacts the total reavy element yield."," Since this is the slowest reaction, $Y_{e}$ impacts the total heavy element yield."
" Under proton-rich conditions (Y,> 0.5). C ‘forms via the standard triple-a reaction sequence *He(a.y) C For Y,«0.5. on the other hand. the neutron capture channel Hetan.yYBeta.n) instead dominates (Woosley&Hoffman 1992)."," Under proton-rich conditions $Y_{e} \gtrsim 0.5$ ), $^{12}$ C forms via the standard $\alpha$ reaction sequence $^{4}$ $\alpha$ $\gamma$ $^{12}$ C. For $Y_{e} < 0.5$, on the other hand, the neutron capture channel $^{4}$ $\alpha$ $\gamma$ $^{9}$ $\alpha$ $^{12}$ C instead dominates \citep{Woosley&Hoffman92}."
. The electron C.fraction also determines the distribution of heavy nuclei synthesized., The electron fraction also determines the distribution of heavy nuclei synthesized.
" When Y,«0.5. nuclei up to the N=50 neutron closed shell CA.~ 90) are created via alpha particle and neutron captures. depending on Y, and the final « fraction (e.g. Woosley&Hotfman1992: Robertsetal. 2010: Arcones&Montes 20105)."," When $Y_{e} \lesssim 0.5$, nuclei up to the N=50 neutron closed shell $A \sim 90$ ) are created via alpha particle and neutron captures, depending on $Y_{e}$ and the final $\alpha$ fraction (e.g. \citealt{Woosley&Hoffman92}; \citealt{Roberts+10}; ; \citealt{Arcones&Montes10}) )."
" If Y, is sutficiently low. even heavier r-process elements (with characteristic peaks at A=130 andA= 195) can be created by additional neutron captures (e.g. Seegeretal. 1965))."," If $Y_{e}$ is sufficiently low, even heavier $r$ -process elements (with characteristic peaks at $A \approx 130$ and$A \approx 195$ ) can be created by additional neutron captures (e.g. \citealt{Seeger+65}) )."
" Under proton-rich conditions (Y,> 0.5). by contrast. mainly Fe-group elements (4~40— 60) are created. although for Y.>0.55 elements up to A=64 may be created by proton captures before the flow reaches ‘bottleneck’ with long 6—decay timescales such as **Ge (e.g. Arcones&Montes2010:: Robertsetal. 2010))."," Under proton-rich conditions $Y_{e} \gtrsim 0.5$ ), by contrast, mainly Fe-group elements $A \sim 40-60$ ) are created, although for $Y_{e} \gtrsim 0.55$ elements up to $A \approx 64$ may be created by proton captures before the flow reaches `bottleneck' with long $\beta-$ decay timescales such as $^{64}$ Ge (e.g. \citealt{Arcones&Montes10}; \citealt{Roberts+10}) )."
" Although the nucleosynthesis is sensitive to Y,. its value in proto-NS winds is rather uncertain."," Although the nucleosynthesis is sensitive to $Y_{e}$, its value in proto-NS winds is rather uncertain."
 The surface of the proto-magnetar is neutronrich (Y... 0.1. but as nucleons are accelerated outwards in the wind they are irradiated by electron neutrinos and antineutrinos.," The surface of the proto-magnetar is neutronrich $Y_{\rm e} \lesssim 0.1$ ), but as nucleons are accelerated outwards in the wind they are irradiated by electron neutrinos and antineutrinos."
" This drives Y, to a value ~0.4—0.6 that depends on the precise v,/T, luminosities and spectra (Qianetal.1993).", This drives $Y_{e}$ to a value $\sim 0.4-0.6$ that depends on the precise $\nu_{e}/\bar{\nu}_{e}$ luminosities and spectra \citep{Qian+93}.
".* The value of Y, is thus sensitive to the details of neutrino transport and interactions (e.g. Rampp&Janka2000:: Mezzacappaetal.2001 :: Duaneal. 20100).", The value of $Y_{e}$ is thus sensitive to the details of neutrino transport and interactions (e.g. \citealt{Rampp&Janka00}; \citealt{Mezzacappa+01}; \citealt{Duan+10}) ).
 Early SN calculations found that Y. decreased from >0.5 at early times to <0.5 as the proto-NS cooled (Woosley1994)., Early SN calculations found that $Y_{e}$ decreased from $\gtrsim 0.5$ at early times to $\lesssim 0.5$ as the proto-NS cooled \citep{Woosley+94}.
". More recent caleulations including additional neutral- interactions. however. find that Y, from ~0.5 to ~0.6 onatimescale ~10 seconds (Hüdepohletal.2010)."," More recent calculations including additional neutral-current interactions, however, find that $Y_{e}$ from $\sim 0.5$ to $\sim 0.6$ on a timescale $\sim 10$ seconds \citep{Hudepohl+10}."
. Unfortunately. few calculations have yet been performed in the case of rapidly spinning proto-NSs.," Unfortunately, few calculations have yet been performed in the case of rapidly spinning proto-NSs."
" Thompsonetal.(2005). find that the v/v, luminosities and temperatures at /<0.6 s are appreciably altered in the case of core collapse with rapid rotation. such that Y, is driven to a higher value than in the non-rotating case."," \citet{Thompson+05} find that the $\nu_{e}$ $\bar{\nu}_{e}$ luminosities and temperatures at $t \lesssim 0.6$ s are appreciably altered in the case of core collapse with rapid rotation, such that $Y_{e}$ is driven to a higher value than in the non-rotating case."
 It is. however. ditficult to extrapolate their results to the much later times ~10—100 s of interest here.," It is, however, difficult to extrapolate their results to the much later times $\sim 10-100$ s of interest here."
" In our calculations below we consider both possibilities 0.5 and Y,>0.5. keeping in mind that the true electron fraction probably [ies in the range 0.4.<Y,0.6 and could vary with time and from event to event. depending on e.g. the NS mass. rotation rate. and obliquity."," In our calculations below we consider both possibilities $Y_{e} < 0.5$ and $Y_{e} \gtrsim 0.5$, keeping in mind that the true electron fraction probably lies in the range $0.4 \lesssim Y_{e} \lesssim 0.6$ and could vary with time and from event to event, depending on e.g. the NS mass, rotation rate, and obliquity."
 However. given the uncertainties. we cannot rule out the possibility that Y. is appreciably lower. in which case even heavier r-process elements with A>100 are produced.," However, given the uncertainties, we cannot rule out the possibility that $Y_{e}$ is appreciably lower, in which case even heavier $r$ -process elements with $A \gtrsim 100$ are produced."
 In particular. although present theory does not tind the necessary conditions for the second or third peak r-process in proto-NS winds. indirect evidence (from e.g. Galactic chemical evolution) suggests that the r-process indeed originates from core-collapse SNe (e.g. Mathewsetal.1992).," In particular, although present theory does not find the necessary conditions for the second or third peak $r$ -process in proto-NS winds, indirect evidence (from e.g. Galactic chemical evolution) suggests that the $r$ -process indeed originates from core-collapse SNe (e.g. \citealt{Mathews+92}) )."
" Given S. Το. and Y,. we estimate the total mass fraction Xi, synthesized in heavy nuclei (A>56) using the following analytic expressions from Robertsetal.(2010) (their eqs. ["," Given $S$, $\tau_{\rm exp}$, and $Y_{e}$ , we estimate the total mass fraction $X_{\rm h}$ synthesized in heavy nuclei $A \gtrsim 56$ ) using the following analytic expressions from \citet{Roberts+10} (their eqs. ["
B3] and eq. [,B3] and eq. [
B3] and [B11]. ef. Hotfmanetal. 1997):,"B3] and [B11], cf. \citealt{Hoffman+97}) ):"
" The critical quantities in square brackets are essentially the reaction rate for tHe — ""C integrated over the timescale available for burning Δ/Top.", The critical quantities in square brackets are essentially the reaction rate for $^{4}$ He $\rightarrow$ $^{12}$ C integrated over the timescale available for burning $\Delta t \sim \tau_{\rm exp}$.
" In the proton-rich case. the reaction rate depends on the entropy squared because triple-a is an ettective three-body reaction and p.«xή at fixed temperature. while in the neutron-rich case the reaction rate depends on entropy cubed because ""Hetan.yY Beta.) C is an etfective reaction."," In the proton-rich case, the reaction rate depends on the entropy squared because $-\alpha$ is an effective three-body reaction and $\rho \propto 1/S$ at fixed temperature, while in the neutron-rich case the reaction rate depends on entropy cubed because $^{4}$ $\alpha$ $\gamma$ $^{9}$ $\alpha$ $^{12}$ C is an effective reaction."
 Figure 2. shows the total mass fraction in heavy nuclei Xi (eq. [4]]), Figure \ref{fig:Xh} shows the total mass fraction in heavy nuclei $X_{\rm h}$ (eq. \ref{eq:Xh}] ])
" as a function of time after core bounce. calculated for a magnetar with By,=5x10% Gand Py=2 ms. as in the wind model shown in Figure |.."," as a function of time after core bounce, calculated for a magnetar with $B_{\rm dip} = 5\times 10^{15}$ G and $P_{0} = 2$ ms, as in the wind model shown in Figure \ref{fig:wind}."
" We show four models. corresponding to ditferent values of the magnetic obliquity y and the wind electron fraction Y,."," We show four models, corresponding to different values of the magnetic obliquity $\chi$ and the wind electron fraction $Y_{e}$ ."
" In the case of a neutron-rich outflow (Y,= 0.43. both aligned and oblique rotators have X,~| throughout the entire wind phase: heavy elements form etticiently because the y ""Hetan. Beta.n) C channel is fast and available (Woosley&Hoffman 1992).."," In the case of a neutron-rich outflow $Y_{e} = 0.4$ ), both aligned and oblique rotators have $X_{\rm h} \sim 1$ throughout the entire wind phase; heavy elements form efficiently because the $^{4}$ $\alpha$ $\gamma$$^{9}$ $\alpha$ $^{12}$ C channel is fast and available \citep{Woosley&Hoffman92}. ."
 By contrast. proton-rich outflows have Xi.<| at all times because carbon must form through the (slower) triple-a channel.," By contrast, proton-rich outflows have $X_{\rm h} \lesssim 1$ at all times because carbon must form through the (slower) $-\alpha$ channel."
" Note that at fixed Y,=0.6. the heavy fraction is larger for an oblique rotator because X, depends sensitively on the wind entropy. which is suppressed in equatorial outflows (eq. [2]."," Note that at fixed $Y_{e} = 0.6$, the heavy fraction is larger for an oblique rotator because $X_{\rm h}$ depends sensitively on the wind entropy, which is suppressed in equatorial outflows (eq. \ref{eq:Somega}] ])."
" Although the precise value of X, dependson the (uncertain) value of Y. we conclude that in all cases a significant fraction of the mass of the wind is locked into heavy nuclei during the epoch of UHECR acceleration."," Although the precise value of $X_{\rm h}$ dependson the (uncertain) value of $Y_{e}$ , we conclude that in all cases a significant fraction of the mass of the wind is locked into heavy nuclei during the epoch of UHECR acceleration."
stress.,stress.
 Orbits of disk particles dift acliabatically in response to absorbing angular momentum al rate inversely proportional to2Br=d(r?0)/dr. the radial gradient of specific angular momentum.," Orbits of disk particles drift adiabatically in response to absorbing angular momentum at rate inversely proportional to$2Br 
\equiv d(r^2\Omega)/dr$, the radial gradient of specific angular momentum."
 Ilowever. a rigorous derivation is more complicated since the Reynolds stress also (transports angular momentum and angular velocity perturbations affect (he angular momentum cdensitv.," However, a rigorous derivation is more complicated since the Reynold's stress also transports angular momentum and angular velocity perturbations affect the angular momentum density."
 These two effects cancel each other leaving equation intact., These two effects cancel each other leaving equation intact.
 Next we substitute (he expression for the mass flux given by equation into the equation of mass conservation and use equations and to eliminate the corotation torque density., Next we substitute the expression for the mass flux given by equation into the equation of mass conservation and use equations and to eliminate the corotation torque density.
 Under the assumption that X varies on a shorter spatialscale than ». ο. D. and dO/dr. we arrive al Specializing to near Ixeplerian rotation. we arrive al The large scale clensitv gradient in a gap is maintained by a balance between torques from principal Lindblad resonances and the viscous stress.," Under the assumption that $\Sigma$ varies on a shorter spatialscale than $\nu$, $\Omega$, $B$, and $d\Omega/dr$, we arrive at Specializing to near Keplerian rotation, we arrive at The large scale density gradient in a gap is maintained by a balance between torques from principal Lindblad resonances and the viscous stress."
 Understanding the action of the two terms on the right hand sides of equations ancl is straightforward., Understanding the action of the two terms on the right hand sides of equations and is straightforward.
 The mass fhix driven by the dillerential corotation torque tends to decrease the density gradient within a region of width ( around corotation., The mass flux driven by the differential corotation torque tends to decrease the density gradient within a region of width $\ell$ around corotation.
 Viscous diffusion acts to maintain the density gradient al (he value it has on larger scales., Viscous diffusion acts to maintain the density gradient at the value it has on larger scales.
 For parameters of interest to us. the densitv. gradient sulfers only a small fractional reduction.," For parameters of interest to us, the density gradient suffers only a small fractional reduction."
 By integrating the steady state version of equation once with respect to r. we find A more revealing form forfirst order corotation resonances is obtained using ICQ? me r/w. and substituting equation into equation (4): which implies," By integrating the steady state version of equation once with respect to $r$, we find A more revealing form forfirst order corotation resonances is obtained using $\phi_{m\pm 1,m}\sim
me(M_p/M_*)(r\Omega)^2$ , $m\approx r/\w$ , and substituting equation into equation ; which implies"
these 10 objects aver.=15.1542.71 kpe with a median of 14.95 kpe.,"these $10$ objects are $r_{e}=15.15\pm2.71$ kpc with a median of $14.
95$ kpc."
 Given the small number statistics that are available. and the inherent dillicul(y in constraining galaxy scalelengths. the median is the more robust measure of the tvpical scalelength.," Given the small number statistics that are available, and the inherent difficulty in constraining galaxy scalelengths, the median is the more robust measure of the typical scalelength."
 Using the median it can be seen from these figures that the one-dimensional analvsis technique employed. by DLIS has systematically overestimated the characteristic size of the 261 galaxies by =50%., Using the median it can be seen from these figures that the one-dimensional analysis technique employed by BLR has systematically overestimated the characteristic size of the 3CR galaxies by $\approx50\%$.
 1n order to check that the use of these two cut-down samples was not overthy biasing the results. a comparison was also performed. between the full. 16-objeet sample and the 12-objects from this sample for which there is also a BLRReclerivect scalelength figure available.," In order to check that the use of these two cut-down samples was not overtly biasing the results, a comparison was also performed between the full $16$ -object sample and the $12$ -objects from this sample for which there is also a BLR-derived scalelength figure available."
 Phe mean scalelength of the 16-0bject sample is r=11.48+1.59 kpe with a median of 9.20 kpe., The mean scalelength of the $16$ -object sample is $r_{e}=11.48\pm1.59$ kpc with a median of $9.20$ kpc.
 Phe corresponding BLIt-derived values for 12 of these objects are r.=15.6641.91 kpe with a median of 15.45 kpc., The corresponding BLR-derived values for $12$ of these objects are $r_{e}=15.66\pm1.91$ kpc with a median of $15.45$ kpc.
 lt can be seen from this that the use of the expanded samples strengthens the conclusion that the BLR. scalelengths are overestimated. with the median. BLR scalcloneth being =TOM greater than the two-dimensional modelling results presented here.," It can be seen from this that the use of the expanded samples strengthens the conclusion that the BLR scalelengths are overestimated, with the median BLR scalelength being $\approx70\%$ greater than the two-dimensional modelling results presented here."
 The absolute Cousins Z-band. magnitudes for the z2OS and z20.2 sub-samples are listed in Table 4 for the four cdillerent cosmologies., The absolute Cousins $I$ -band magnitudes for the $z\simeq0.8$ and $z\simeq0.2$ sub-samples are listed in Table \ref{mags} for the four different cosmologies.
 The values shown are caleulated from integrating the best-fit de Vaucouleurs profile. to infinite μασας in order to be consistent with the published results for the low-2 AGN hosts (AIeLureaf 1999. Dunlopaf 1999).," The values shown are calculated from integrating the best-fit de Vaucouleurs profile to infinite radius in order to be consistent with the published results for the $z$ AGN hosts (McLure 1999a, Dunlop 1999)."
 As was pointed out in Section 2. the objects comprising the sub-sample were imaged through two separate LIST filters: ESI4W and ET85LP., As was pointed out in Section \ref{smalltab} the objects comprising the sub-sample were imaged through two separate HST filters; F814W and F785LP.
 Ehe FSI4W filter is à member of the standard LST filter set and. closely mimics the Cousins Z-band filter., The F814W filter is a member of the standard HST filter set and closely mimics the Cousins $I$ -band filter.
 As a result. it was decided to convert all of the integrated apparent magnitudes to their equivalent Cousins {ρα value.," As a result, it was decided to convert all of the integrated apparent magnitudes to their equivalent Cousins $I$ -band value."
 The 5 objects from the sub-sample imaged with the FSI4W filter had no correction applied to them since the similarity of the FSI4W and Cousins Z-band filters is such that for elliptical galaxies the dilference in magnitude is expected to be less than 0.05 in all cases., The $5$ objects from the sub-sample imaged with the F814W filter had no correction applied to them since the similarity of the F814W and Cousins $I$ -band filters is such that for elliptical galaxies the difference in magnitude is expected to be less than $0.05$ in all cases.
 This level of photometric accuracy is substantially ereater than is possible in the face of the uncertainties in determining the host magnitucles., This level of photometric accuracy is substantially greater than is possible in the face of the uncertainties in determining the host magnitudes.
 The throughput for the F7S5LP filter is significantly different from that of the Cousins Z-band. filter., The throughput for the F785LP filter is significantly different from that of the Cousins $I$ -band filter.
 “Phe conversion from magnitudes obtained through this filter to Cousins {αμα magnitudes is therefore more complicated., The conversion from magnitudes obtained through this filter to Cousins $I$ -band magnitudes is therefore more complicated.
 The original strategy to overcome this problem was to make use of the one object from the 16-0bject. saniple (30239) or which there are exposures in both filters available with comparable signal-to-noise., The original strategy to overcome this problem was to make use of the one object from the $16$ -object sample (3C239) for which there are exposures in both filters available with comparable signal-to-noise.
 However. this object has the ughest redshift. 2=L7S1. of all of the 28 objects in the BL sample and subsequently the two images were too faint o get a reliable conversion.," However, this object has the highest redshift, $z=1.781$, of all of the $28$ objects in the BLR sample and subsequently the two images were too faint to get a reliable conversion."
 In addition to this. 3€239 has a iehlvy distorted morphology in the HIST images (BLE. 1997). making it unclear whether a conversion factor obtained from his object would be applicable to the mostly. undistorted objects in the sub-sample.," In addition to this, 3C239 has a highly distorted morphology in the HST images (BLR 1997), making it unclear whether a conversion factor obtained from this object would be applicable to the mostly undistorted objects in the sub-sample."
 In light of this. the method used," In light of this, the method used"
but has the advantage that it can be expressed in terms of effective quantities usually quoted in observational studies (e.g. MacArthur et al.,but has the advantage that it can be expressed in terms of effective quantities usually quoted in observational studies (e.g. MacArthur et al.
 2003)., 2003).
" We did additional tests replacing the exponential law used for the outer regions by either a power law or a Sérrsic profile, but we found that, in general, the fits were poorer."," We did additional tests replacing the exponential law used for the outer regions by either a power law or a Sérrsic profile, but we found that, in general, the fits were poorer."
" In order to properly model the break at which the transition between the exponential and the Sérrsic law takes place, which is clearly different for the eight simulations, we repeated the fits changing the break radius and estimated the goodness of the fit using the parameter Q, defined by: where n is the number of data points."," In order to properly model the break at which the transition between the exponential and the Sérrsic law takes place, which is clearly different for the eight simulations, we repeated the fits changing the break radius and estimated the goodness of the fit using the parameter $Q$, defined by: where $n$ is the number of data points."
" By minimizing Q we selected the best-fit parameters, which we show in Table 5.."," By minimizing $Q$ we selected the best-fit parameters, which we show in Table \ref{disk_bulge}."
" Furthermore, for each simulation we have repeated the fits twice, considering the profiles either to 1.5 or 2 times the corresponding optical radii, and found almost no difference in the best-fit parameters (in all cases lower than a fewp"," Furthermore, for each simulation we have repeated the fits twice, considering the profiles either to $1.5$ or $2$ times the corresponding optical radii, and found almost no difference in the best-fit parameters (in all cases lower than a few."
ercent)?.., In Fig.
 In Fig. 13 we show the best-fit model (solid line) as well as the relative contributions of the Sérrsic and exponential laws (dashed lines)., \ref{profiles} we show the best-fit model (solid line) as well as the relative contributions of the Sérrsic and exponential laws (dashed lines).
" The inner spheroids (“bulges” if bars were absent) are characterized by central surface mass densities between 10°° and 10? Μο kpc~?, and effective radii in the range 1.5—2.5 kpc."," The inner spheroids (“bulges” if bars were absent) are characterized by central surface mass densities between $10^{8.5}$ and $10^9$ $_\odot$ $^{-2}$, and effective radii in the range $1.5-2.5$ kpc."
" In terms of their shapes, the Sérrsic parameters n are typically around 1, i.e., similar to exponential profiles."," In terms of their shapes, the Sérrsic parameters $n$ are typically around $1$, i.e., similar to exponential profiles."
" Aq-A-5 and Aq-B-5 have the highest n parameters, a reflection of their very concentrated profiles; on the contrary, Aq-H-5 has the lowest n value indicating a shallower profile in the central regions."," Aq-A-5 and Aq-B-5 have the highest $n$ parameters, a reflection of their very concentrated profiles; on the contrary, Aq-H-5 has the lowest $n$ value indicating a shallower profile in the central regions."
" The outer regions of the spheroids have much lower central surface densities, with X~10'—10* Mo kpc?, and scale-lengths of the order of 4—10 kpc."," The outer regions of the spheroids have much lower central surface densities, with $\Sigma_{\rm o}\sim 10^{7}-10^{8}$ $_\odot$ $^{-2}$, and scale-lengths of the order of $4-10$ kpc."
" We find a wide range of relative contributions of the inner and outer components to the total spheroid mass, as can be clearly seen in Fig. 13.."," We find a wide range of relative contributions of the inner and outer components to the total spheroid mass, as can be clearly seen in Fig. \ref{profiles}."
" We note that, in the cases where bars are present, the profiles should perhaps include an additional component."," We note that, in the cases where bars are present, the profiles should perhaps include an additional component."
 This might explain the (small) differences between the results here and those presented in Scannapieco et al., This might explain the (small) differences between the results here and those presented in Scannapieco et al.
"(2010)9.. However, we do not find larger Q values for galaxies with bars, suggesting that the number of fitting parameters is large enough to fit the 1D profiles of Fig. 13,,"," However, we do not find larger $Q$ values for galaxies with bars, suggesting that the number of fitting parameters is large enough to fit the 1D profiles of Fig. \ref{profiles},"
 even if bars are not included as an extra component., even if bars are not included as an extra component.
 The structural properties of the inner and outer spheroids show very good agreement with the results for simulations with lower resolution (Table 5))., The structural properties of the inner and outer spheroids show very good agreement with the results for simulations with lower resolution (Table \ref{disk_bulge}) ).
" We find that all the fitting parameters change by less than 1096, with the exception of Aq-E-6 where larger differences (25— 4096) are"," We find that all the fitting parameters change by less than $10\%$, with the exception of Aq-E-6 where larger differences $25-40\%$ ) are"
 Staudard cosmology is based on the assuuptiou that the universe is spatially homogencous. at least on scales sutiicicnutly laree to justify its approximation by a Fricchnanu-Leiaittre-Robertson-Walker (FLRW) uoclel.," Standard cosmology is based on the assumption that the universe is spatially homogeneous, at least on scales sufficiently large to justify its approximation by a Friedmann-Lemaîttre-Robertson-Walker (FLRW) model."
 The Heh isotropy ineasured iu the cosnüc microwave backeround radiation (CAIBR) is usually taken as strong evidence in support of this hvpothesis., The high isotropy measured in the cosmic microwave background radiation (CMBR) is usually taken as strong evidence in support of this hypothesis.
 Coupled to a elobal. respectively local. Copernican principle. this observation has led to the forxiuulatiou of the Ehlers-CGeren-Sachs CEhlllers. Coren. Sachs. 1968). respectively almost Eblers-Coren-Sachs (Stoeger. Aaartens. Ellis. 1995). theorems.," Coupled to a global, respectively local, Copernican principle, this observation has led to the formulation of the Ehlers-Geren-Sachs (Ehlers, Geren, Sachs, 1968), respectively almost Ehlers-Geren-Sachs (Stoeger, Maartens, Ellis, 1995), theorems."
 Although the assuuptious undoerlviues these theorems have been discussed by some authors (Clarkson Barrett. 1999: Céll&xier. 2000a). they are generally considered asx robust support for the Cosmological Principle.," Although the assumptions underlying these theorems have been discussed by some authors (Clarkson Barrett, 1999; Céllérrier, 2000a), they are generally considered as robust support for the Cosmological Principle."
 The structures seeu in Oogalaxy cataloguesOo - C»groups. clusters ancl superclusters. distributed aloug voicls. filaments aud wa s- are not viewed as contradicting this Principle. as the scales on which the universe is assumed. to be homogeneous are much larger than those subtended bv these structures.," The structures seen in galaxy catalogues - groups, clusters and superclusters, distributed along voids, filaments and walls - are not viewed as contradicting this Principle, as the scales on which the universe is assumed to be homogeneous are much larger than those subtended by these structures."
 Furthermore. the most probable existence of dark matter muplies that the huuinous matter seen iu the galaxies docs not represent the entire source of eravitational energv.," Furthermore, the most probable existence of dark matter implies that the luminous matter seen in the galaxies does not represent the entire source of gravitational energy."
 Thus. the high isotropy of the CMDR could. iu principle. be reconciled with an Inhomogeneous buninous matter distribution. provided he nüssmegC» mass for an honiosgeneousD> pattern would be xovided by dark matter.," Thus, the high isotropy of the CMBR could, in principle, be reconciled with an inhomogeneous luminous matter distribution, provided the missing mass for an homogeneous pattern would be provided by dark matter."
" This iuportaut issue. which las oen formated as the problem of Πιοιο whether ""light races naass is still the subject of a uuuber of teresting works (sce e.g. Balicall et al."," This important issue, which has been formulated as the problem of finding whether “light traces mass” is still the subject of a number of interesting works (see e.g., Bahcall et al.,"
 2000). but appears far roni being settled.," 2000), but appears far from being settled."
 Towever. a majority of cosmologists seenus to feel more coufortable with the idea that the structure distribution should preseut a transition towards rolmogencity at a muore or less laree scale.," However, a majority of cosmologists seems to feel more confortable with the idea that the structure distribution should present a transition towards homogeneity at a more or less large scale."
 However. the consensus on a homogcucous feature of structures. even on very huge scales; has Hever been complete.," However, the consensus on a homogeneous feature of structures, even on very large scales, has never been complete."
 The debate exteuds back to Newton (pro-homogencous) VOLS haut and Lambert (pro-hierarchical)., The debate extends back to Newton (pro-homogeneous) versus Kant and Lambert (pro-hierarchical).
 The pro-hierarcliica approach was reconsidered in a wore moderl wav w Charlier (1908. 1922). and later developed by de Vancouleurs (1953. 1970). who exhibited a universal deusitv radius power law.," The pro-hierarchical approach was reconsidered in a more modern way by Charlier (1908, 1922), and later developed by de Vaucouleurs (1953, 1970), who exhibited a universal density radius power law."
 Recent ideas ou fractals (ALaudexot. 1982) led to the pioneering work bv DPietronero (LOST).," Recent ideas on fractals (Mandelbrot, 1982) led to the pioneering work by Pietronero (1987)."
 Since then. iurproveineuts in observational techniques have increased the number and depth of available three dimensional ealaxy catalogues.," Since then, improvements in observational techniques have increased the number and depth of available three dimensional galaxy catalogues."
 A program of systematic analyses of these catalogues has led to the claim that oue cau identify a fractal distribution up to 150 5.+ Mpe (Sylos, A program of systematic analyses of these catalogues has led to the claim that one can identify a fractal distribution up to 150 $h^{-1}$ Mpc (Sylos
Observations of normal spiral galaxies bv Schmidt (1959) first suggested that (heir star formation rates (SER) scales with the elobal properties.,Observations of normal spiral galaxies by Schmidt (1959) first suggested that their star formation rates (SFR) scales with the global properties.
 This conclusion was extended to other galaxies wilh higher SFR. such as the nuclear regions of spiral galaxies and Ultra Luminous Infralted Galaxies (ULIRGs) bv IxXennicutt (1993).," This conclusion was extended to other galaxies with higher SFR, such as the nuclear regions of spiral galaxies and Ultra Luminous InfraRed Galaxies (ULIRGs) by Kennicutt (1998)."
" These galactic scale observations have lead (o an empirical law for star formation called Ixennicutt-Schinidt (IAS) Law: where M4, and M, ave the gas surface density and SFR per unit area.", These galactic scale observations have lead to an empirical law for star formation called Kennicutt-Schmidt (KS) Law: where $\rm \Sigma_{gas}$ and $\rm \dot{\Sigma}_{star}$ are the gas surface density and SFR per unit area.
 since star formation is a local process that happens on subpirsece scale. (he correlation with global (galactic scale) quantities such as averaged M44. suggest the existence of a physical connection between galactic (> 1 kpe) and subparsec scales.," Since star formation is a local process that happens on subparsec scale, the correlation with global (galactic scale) quantities such as averaged $\rm \Sigma_{gas}$, suggest the existence of a physical connection between galactic $>$ 1 kpc) and subparsec scales."
 Motivated by the existence of the INS law. several authors have tried (o find (he link in which a elobal/laree," Motivated by the existence of the KS law, several authors have tried to find the link in which a global/large"
'TDM and other similar models by considering deviations [rom 2) Mass dependent textures which induce relations between mixine matrix elements aud mass eigenvalues.,TBM and other similar models by considering deviations from 2) Mass dependent textures which induce relations between mixing matrix elements and mass eigenvalues.
 Such textures naturally accommodate a non-zero 043., Such textures naturally accommodate a non-zero $\theta_{13}$.
" Some examples of these are zero textures |6].. vanishing minors [T. δι, hybrid textures |9]."," Some examples of these are zero textures \cite{6}, vanishing minors \cite{7,8}, , hybrid textures \cite{9}."
". Zero textures have been partüeularly successful in explaining both the quark ancl the lepton masses and In this work we identify a class of mass dependent textures which supplemented with the assumption of a large value of effective neutrino mass M, naturally predict near maximal 85; and non-zero £45.", Zero textures have been particularly successful in explaining both the quark and the lepton masses and In this work we identify a class of mass dependent textures which supplemented with the assumption of a large value of effective neutrino mass $M_{ee}$ naturally predict near maximal $\theta_{23}$ and non-zero $\theta_{13}$ .
 Recently. it was shown by Grimus [10]. that near maximal atmospheric mixing is predicted for class D3 aud B+ of two zero textures supplemented with the assumption of quasi degeneracy.," Recently, it was shown by Grimus \cite{10} that near maximal atmospheric mixing is predicted for class B3 and B4 of two zero textures supplemented with the assumption of quasi degeneracy."
 We consider (wo vanishing minors of (he neutrino mass matrix in the flavor basis together with the assumption of large M., We consider two vanishing minors of the neutrino mass matrix in the flavor basis together with the assumption of large $M_{ee}$.
 This assumption is well motivated by the extensive search for this parameter in the ongoing experiments., This assumption is well motivated by the extensive search for this parameter in the ongoing experiments.
 We found Chat the two cases of two vanishing minorsviz., We found that the two cases of two vanishing minorsviz.
" B; and By [Table 1] in the limit of large M, predict near maximal atmospheric mixing and (his property holds irrespective of the values of solar and reactor mixing angles.", $B_5$ and $B_6$ [Table 1] in the limit of large $M_{ee}$ predict near maximal atmospheric mixing and this property holds irrespective of the values of solar and reactor mixing angles.
 The seesaw mechanism |ll] is regarded as the prime candidate for understanding (the scale of neutrino masses., The seesaw mechanism \cite{11} is regarded as the prime candidate for understanding the scale of neutrino masses.
" In (the framework of tvpe-I seesaw mechanism. the effective Majorana mass matrix M, is given by where A is the Dirac neutrino mass matrix and Mg is the right-handecl Majorana mass malrix."," In the framework of type-I seesaw mechanism, the effective Majorana mass matrix $M_\nu$ is given by where $M_D$ is the Dirac neutrino mass matrix and $M_R$ is the right-handed Majorana mass matrix."
" In the framework of twpe-l seesaw mechanism Ad, is a quantity derived from Mp and Aly.", In the framework of type-I seesaw mechanism $M_\nu$ is a quantity derived from $M_D$ and $M_R$ .
" Therefore. zeros of Mj, and A), have a deeper theoretical meaning."," Therefore, zeros of $M_D$ and $M_R$ have a deeper theoretical meaning."
" In a basis where Mp is diagonal. the zeros of Mj propagate as zero minors in M, "," In a basis where $M_D$ is diagonal, the zeros of $M_R$ propagate as zero minors in $M_\nu$ ."
In the basis where the charged lepton mass matrix isdiagonal. there are fifteenpossibletwo," In the basis where the charged lepton mass matrix isdiagonal, there are fifteenpossibletwo"
The Luk between the SFR aud the FIR hunuinositv is nore indirect than for the UV Lunimositv since it depends of the dust heating which involves all types of stars.,The link between the SFR and the FIR luminosity is more indirect than for the UV luminosity since it depends of the dust heating which involves all types of stars.
 Nevertheless. in starbusting galaxies the situation ix expected to be less complex since the dust heating is dominated by the voung stars.," Nevertheless, in starbusting galaxies the situation is expected to be less complex since the dust heating is dominated by the young stars."
" Under such conditions Ixeuuicutt (1998)) has related the FIR luminosity to the star formation rate SFR=1.7110!""Lpgig(L.:) where Lgig is he tota FIR huuinositv."," Under such conditions Kennicutt \cite{kennicutt}) ) has related the FIR luminosity to the star formation rate $ \rm SFR=1.71~10^{-10}~ L_{FIR} (L\odot)$ where $ \rm
L_{FIR}$ is the total FIR luminosity."
 This conversiou factor SER/Lrx ds obtained using svuthesis population models and is also subject to uncertaiuties ou the stellar tracks. the IMFE or the star formation history.," This conversion factor $\rm SFR/L_{FIR}$ is obtained using synthesis population models and is also subject to uncertainties on the stellar tracks, the IMF or the star formation history."
 Frou. a comparison with the calculations of Lehuert Ueckman (1996)). Meurer et al. (19973) ," From a comparison with the calculations of Lehnert Heckman \cite{lehnert}) ), Meurer et al. \cite{meurheckI}) )"
and Buat Nu (1996)) we estimate the uncertaiutv of the order of ( for SERκ. Therefore we cau reasonably estimate that the SER deduced from the observed UV luminosity added to that deduced from the FIR one is also uncertain by a factor ~ 50%., and Buat Xu \cite{buxu}) ) we estimate the uncertainty of the order of $\%$ for $\rm SFR/L_{FIR}$ Therefore we can reasonably estimate that the SFR deduced from the observed UV luminosity added to that deduced from the FIR one is also uncertain by a factor $\sim 50\%$ .
 Nevertheless it 11st be noticed that the couversion formmlac only apply to galaxies which have experinieuted a continuous star formation for at least ~105 vears and will uot be valid for galaxies with more episodic star formation. especially post starbursting galaxies.," Nevertheless it must be noticed that the conversion formulae only apply to galaxies which have experimented a continuous star formation for at least $\sim 10^8$ years and will not be valid for galaxies with more episodic star formation, especially post starbursting galaxies."
 The IRAS/FOCA sample is FIR selected. thus it is vased against very blue chvarf galaxies which may exhibit episodic bursts of star formation as suggested w Fioc Rocca-Voluucranee (1999)).," The IRAS/FOCA sample is FIR selected, thus it is biased against very blue dwarf galaxies which may exhibit episodic bursts of star formation as suggested by Fioc Rocca-Volmerange \cite{fioc}) )."
 Therefore we cau expect that he derivation of a SFR from the FIR aud UV. eiiissions js valid or this sample of galaxies., Therefore we can expect that the derivation of a SFR from the FIR and UV emissions is valid for this sample of galaxies.
 The comparison of he star formation rates obtained by adding he FIR aud UV (observed) emissious to those deduced from the UV fluxes after a correction for extinction can be useful to est the consistency of both methods., The comparison of the star formation rates obtained by adding the FIR and UV (observed) emissions to those deduced from the UV fluxes after a correction for extinction can be useful to test the consistency of both methods.
 Therefore we have calculated he star formation rates for the IRAS/FOCA siuuple of galaxies adding the contribution of he FIR aud UV (not corrected for extinction) enissious and using the conversion- forimla of Trever et al. 19983) , Therefore we have calculated the star formation rates for the IRAS/FOCA sample of galaxies adding the contribution of the FIR and UV (not corrected for extinction) emissions and using the conversion formula of Treyer et al. \cite{treyer}) )
for the UV and I&enuicutt (1998)) for the FIR., for the UV and Kennicutt \cite{kennicutt}) ) for the FIR.
 Total FIR fluxes have been estimated using the relation found by Duat Burearclla (1998)) between the ratio of the total dust flux to the FIR (10-120/12) flux auc füo/figg.," Total FIR fluxes have been estimated using the relation found by Buat Burgarella \cite{bubu}) ) between the ratio of the total dust flux to the FIR $\mu$ m) flux and $\rm 
f_{60}/f_{100}$."
 The estimated SFRs can be compared to the ones obtained after correction of the UV fluxes youn extinction., The estimated SFRs can be compared to the ones obtained after correction of the UV fluxes from extinction.
 The correlation between the two estimates is very good but the SET deduced from the (FIR UV) emissious are higher than the SFR:4. deduced from the UV. corrected cussion by a factor 1., The correlation between the two estimates is very good but the SFRs deduced from the $\rm FIR+UV$ ) emissions are higher than the SFRs deduced from the UV corrected emission by a factor 1.4.
 Another interesting comparison is that of the relative contribution of the UW (not corrected for extinction) aud FIR eunussous to the total (FIR| UV) SER., Another interesting comparison is that of the relative contribution of the UV (not corrected for extinction) and FIR emissions to the total $\rm FIR+UV$ ) SFR.
" For our siuiple of IRAS/FOCA ealaxies the relative coutribution of the UV aud FIR euissious to the total SFR are 0.3 aud 0.7 respectively,", For our sample of IRAS/FOCA galaxies the relative contribution of the UV and FIR emissions to the total SFR are 0.3 and 0.7 respectively.
 However. these calculations assume that he FIR dus is exclusively due to the heating by vouug stars.," However, these calculations assume that the FIR flux is exclusively due to the heating by young stars."
 Since all our galaxics are certainly not starburstiug objects the contriion of the cussion of dust heated v old evolved stars must be deduced from the FIR flix fore translating 1 into star formation rate reduci he contribution of he FIR ciunission to the SER., Since all our galaxies are certainly not starbursting objects the contribution of the emission of dust heated by old evolved stars must be deduced from the FIR flux before translating it into star formation rate reducing the contribution of the FIR emission to the SFR.
 Let us assume that old stars coutribute for 30 54 of the dust jieating (Xu 1990... Buat Xu 19963). then the ratio of he SER deduced from the (FIR|UV) emissions aud the SER deduced TOlu ιο UV. corrected emission is reduced roni 1.{ to 1.1 aud the relative coutributions of the UV. aud FIR emissious to the total SFR are now 0.1 aud 0.6.," Let us assume that old stars contribute for 30 $\%$ of the dust heating (Xu \cite{xu}, Buat Xu \cite{buxu}) ), then the ratio of the SFR deduced from the $\rm 
FIR+UV$ ) emissions and the SFR deduced from the UV corrected emission is reduced from 1.4 to 1.1 and the relative contributions of the UV and FIR emissions to the total SFR are now 0.4 and 0.6."
 Cuven all the uncertainties inherent το these calculations we must be cautious in our conclusions., Given all the uncertainties inherent to these calculations we must be cautious in our conclusions.
" We can sav that the correctious for dust extinction deduced from the FIR/UV fiux ratio alu applied to the UV. observed cussions lead to a SER cousistent with that obtained by adding the SFRs deduced frou, both the FIR and observed UV. cussions.", We can say that the corrections for dust extinction deduced from the FIR/UV flux ratio and applied to the UV observed emissions lead to a SFR consistent with that obtained by adding the SFRs deduced from both the FIR and observed UV emissions.
 This nakes us confident in our estimate of the extinction., This makes us confident in our estimate of the extinction.
 Cuder the hypothesis of an average star formation of the local universe coutiumous over ~LOS years. we cau derive elobal star formation rates from the local FIR and UV Iuniunositv censitics: pl is calculated using the value of pgi (section 1.1) multiplied by 1.5 to account for the total dust luninosity deusitv (c.g. Nu Buat 1995.. Meurer et al. 1999)):," Under the hypothesis of an average star formation of the local universe continuous over $\sim~10^8$ years, we can derive global star formation rates from the local FIR and UV luminosity densities: $\rm \rho_{SFR}^{FIR}$ is calculated using the value of $\rm \rho_{FIR}$ (section 4.1) multiplied by 1.5 to account for the total dust luminosity density (e.g. Xu Buat \cite{xubu}, Meurer et al. \cite{meurheck}) );"
 Pen is caleulated from pos.," $\rm 
\rho_{SFR}^{UV}$ is calculated from $\rm\rho_{0.2}$."
 We find: This time the contributions of the FIR aud UV are very simular. this is due to the lower value of the FIR to UV density ratio as compared to the FIR to UV fiux ratio ofiuividual galaxies.," We find: This time the contributions of the FIR and UV are very similar, this is due to the lower value of the FIR to UV density ratio as compared to the FIR to UV flux ratio of individual galaxies."
 Then the volumc-average star formation rate deduced from the UN. hunuimositv. deusitv uot corrected for dust extinction mst be multiplied bx a factor —2 to account for the elobal extinction. this correspolnis fo a inecan extinction of 0.75 mag at 0.2 qnn. As diseussed. before this rather low value is due to the contribution of faint blue galaxies to the UV. luninosity Actually. using the pg/po.» ratio gives an extinction of 0.77 1nag using the model of Meurer et al.," Then the volume-average star formation rate deduced from the UV luminosity density not corrected for dust extinction must be multiplied by a factor $\sim 2$ to account for the global extinction, this corresponds to a mean extinction of 0.75 mag at 0.2 $\mu$ m. As discussed before this rather low value is due to the contribution of faint blue galaxies to the UV luminosity Actually, using the $\rm \rho_{FIR}/\rho_{0.2}$ ratio gives an extinction of 0.77 mag using the model of Meurer et al."
 aud 0.55 mae for our polvnoiial fit (section 3.1) As already uuderlined. the difference is likely to come from the coutribution of the old stars to the dust heating: let us assume that 30% of the FIR ciission comes from these old stars aucl is not related to the recent star formation then we find," and 0.55 mag for our polynomial fit (section 3.1) As already underlined, the difference is likely to come from the contribution of the old stars to the dust heating: let us assume that $\%$ of the FIR emission comes from these old stars and is not related to the recent star formation then we find"
reaches almost for the strongest shocks.,reaches almost for the strongest shocks.
 The average value of the acoustic flux at this height is1L. which is au order of magnitude below the required value to balance. the radiative losses.," The average value of the acoustic flux at this height is, which is an order of magnitude below the required value to balance the radiative losses."
 Ouly at the moments when shocks reach the chromosphere the energy supplied by the slow acoustic waves is similar to the chromospheric radiative losses., Only at the moments when shocks reach the chromosphere the energy supplied by the slow acoustic waves is similar to the chromospheric radiative losses.
" We have preseuted ποΊσα. simulations which closely reproduce the real wave propagation observed in the mubra of a sunspot from the photosphere to the chromosphere,", We have presented numerical simulations which closely reproduce the real wave propagation observed in the umbra of a sunspot from the photosphere to the chromosphere.
 T1ο analysis of these observations was previously report«d diu Felipeetal.(2010b)., The analysis of these observations was previously reported in \citet{Felipe+etal2010b}.
. As proved in Felipectal.(2Ha)... our code is able to manage the xopagationOo of waves with lougC» realistic periods. even iu the hard coucditions imposed by the hieh Alfvéóun speeds at the suuspot chromosphere.," As proved in \citet{Felipe+etal2010a}, our code is able to manage the propagation of waves with long realistic periods, even in the hard conditions imposed by the high Alfvénn speeds at the sunspot chromosphere."
 The similarity between the simulations aud the observations is achieved in two steps., The similarity between the simulations and the observations is achieved in two steps.
 Firsth. we had ο construct a MIIS model of a sunspot with properties as close as possible to the observed onc.," Firstly, we had to construct a MHS model of a sunspot with properties as close as possible to the observed one."
 We have ollowed the method of Ihomoeuko&Collaclos(2008).. oeitroduciug some iodificatious in order to take into ποσο thennodyvuanüc and magnetic properties of re particulary observed suuspot. as retrieved from the oewersion of the ine.," We have followed the method of \citet{Khomenko+Collados2008}, introducing some modifications in order to take into account thermodynamic and magnetic properties of the particular observed sunspot, as retrieved from the inversion of the line."
 Secondly. this static model was perturbed by a orce Which dives a vertical velocity simular to the LOS velocity 1icasured from the Doppler shift ofthe ine.," Secondly, this static model was perturbed by a force which drives a vertical velocity similar to the LOS velocity measured from the Doppler shift of the line."
 Once the observed. velocity is introduced in the MITIS atmosphere. it imaiulv drives slow maegneto-acoustic waves du the low-.) region.," Once the observed velocity is introduced in the MHS atmosphere, it mainly drives slow magneto-acoustic waves in the $\beta$ region."
 Most of their power ids concentrated in the 5 iminute band. with frequencies between 3 and £L ΙΤ. aud they form standing waves due to the higher cutoff frequency of the atmosphere.," Most of their power is concentrated in the 5 minute band, with frequencies between 3 and 4 mHz, and they form standing waves due to the higher cutoff frequency of the atmosphere."
 The driver also generates high frequency waves. which propagate upwards to the chromosphere.," The driver also generates high frequency waves, which propagate upwards to the chromosphere."
 Iun their travel through the suuspot atmosphere. they reach the formation height of several observed spectral lines.," In their travel through the sunspot atmosphere, they reach the formation height of several observed spectral lines."
 The siulated velocity maps and power spectra at their correspouding formation heights reproduce reasonably well those obtained for the chromospheric ccore aud line., The simulated velocity maps and power spectra at their corresponding formation heights reproduce reasonably well those obtained for the chromospheric core and line.
 In the case of the line. the comparison is hiudered due to the poor qualitv of the observed velocity ap.," In the case of the line, the comparison is hindered due to the poor quality of the observed velocity map."
 The comparison between the simulated and observed velocity iaps at chromospheric heights reveals some phase delay at the strougest wavefrouts. in the seuse that the observed wavefrouts lag the simulated. ones.," The comparison between the simulated and observed velocity maps at chromospheric heights reveals some phase delay at the strongest wavefronts, in the sense that the observed wavefronts lag the simulated ones."
 In the case of the ccore the phase lag is 0.£1z., In the case of the core the phase lag is $\pi$.
 The delay of the observed relative to the simulated yvelocities Is somewhat smaller., The delay of the observed relative to the simulated velocities is somewhat smaller.
1 This delay. together with the lareer amplitude of the observed strongest shocks. iav indicate that iu he dvuanuücal suuspot chromosphere. the noulincar waves shift the formation height of spectral lines to higher lavers.," This delay, together with the larger amplitude of the observed strongest shocks, may indicate that in the dynamical sunspot chromosphere, the nonlinear waves shift the formation height of spectral lines to higher layers."
 This is especially evident for the CCOLC., This is especially evident for the core.
 Tn the analysis of these simulations. we have chosen a certain height represeutative of the laver from. which the information about the velocity measured with the observed line comes.," In the analysis of these simulations, we have chosen a certain height representative of the layer from which the information about the velocity measured with the observed line comes."
 Obviously. this procedure is a simplification. since the contribution function from every spectral line spans over a thick laver.," Obviously, this procedure is a simplification, since the contribution function from every spectral line spans over a thick layer."
 The height that we are considering corresponds to the average height derived iu Felipeetal.(2010h)., The height that we are considering corresponds to the average height derived in \citet{Felipe+etal2010b}.
. This approximalon seenis to be eood for the photospheric lues. but uo for the highly 1ondinear chromospheric region.," This approximation seems to be good for the photospheric lines, but not for the highly nonlinear chromospheric region."
 Our sinnulatious reproduce the pliase aid ;uuplificatiou spectra between several pairs of lines wiha remarkable natch., Our simulations reproduce the phase and amplification spectra between several pairs of lines with a remarkable match.
 For those spectra between a pjotosphierie and a chromospheric signal. the phase «ifereuce shows stationary waves with Ao=0 below 1 ullz.," For those spectra between a photospheric and a chromospheric signal, the phase difference shows stationary waves with $\Delta \phi =0$ below 4 mHz."
 At lieher Yrequencies the waves progressively prooxiweate with Ao increasing with the frequency., At higher frequencies the waves progressively propagate with $\Delta \phi$ increasing with the frequency.
 The eimpincal cutoff of 1 ullz is assigned from the estimation of he frequency at which the phase differcuce starts to depart from Ao=0., The empirical cutoff of 4 mHz is assigned from the estimation of the frequency at which the phase difference starts to depart from $\Delta \phi =0$.
 Iu the solar atmosphere. the cutoff is a local parameter which depends on the temperature aud is stratificated with height.," In the solar atmosphere, the cutoff is a local parameter which depends on the temperature and is stratificated with height."
 In our AWTS model its πια value is 6 ullz. aud ouly waves with frequencies above this value can propagate all over the atmosphere.," In our MHS model its maximum value is 6 mHz, and only waves with frequencies above this value can propagate all over the atmosphere."
 However. iu most of the atmosphere the local cutoff is below 6 1lIIz. aud for his reasou the phase difference spectra show propagation or frequencies im the rauge between [aud 6 lI.," However, in most of the atmosphere the local cutoff is below 6 mHz, and for this reason the phase difference spectra show propagation for frequencies in the range between 4 and 6 mHz."
 The ouly strong discrepancy between observations aud siuulations is found iu the amplification spectra at requeucies above 8 τι] (Figures 15. aud 15))., The only strong discrepancy between observations and simulations is found in the amplification spectra at frequencies above 8 mHz (Figures \ref{fig:spiegel_sica} and \ref{fig:spiegel_fehe}) ).
 The auplification is cleary higher in the simulations than iu 16 observations. aud the observed power spectra of the Tiromiosphlieric lines show lower power than the simulated equivalent (Figure 11)).," The amplification is cleary higher in the simulations than in the observations, and the observed power spectra of the chromospheric lines show lower power than the simulated equivalent (Figure \ref{fig:he_spectra}) )."
 There are several causes that could be responsible for this discrepancy., There are several causes that could be responsible for this discrepancy.
 The comparison of the power spectra of the velocity retrieved from the liue aud the one obtained iu the simulation at its formation height (iiddle panel of Figure 6)) shows sole slight differences which are especially siguificaut above 8 iulIz., The comparison of the power spectra of the velocity retrieved from the line and the one obtained in the simulation at its formation height (middle panel of Figure \ref{fig:si_spectra}) ) shows some slight differences which are especially significant above 8 mHz.
" Since the amplitude of the waves increases exponentially with height. these small cdiscrepaucies could bethe origin of the larger ones that appear at larger weights,"," Since the amplitude of the waves increases exponentially with height, these small discrepancies could be the origin of the larger ones that appear at larger heights."
 For example. Figure 6 shows that the simulation las an excess of power at the height of the driver at 13 unllz and 17 mlz. and these frequencies also preseut ueher power than the observations at the chromosplere (Figure 111).," For example, Figure \ref{fig:si_spectra} shows that the simulation has an excess of power at the height of the driver at 13 mHz and 17 mHz, and these frequencies also present higher power than the observations at the chromosphere (Figure \ref{fig:he_spectra}) )."
 However. some frequencies where the power of the photospheric velocity in the simulation is not larger han the power of the line also present a significative difference at higher avers.," However, some frequencies where the power of the photospheric velocity in the simulation is not larger than the power of the line also present a significative difference at higher layers."
 It is possible that most of the power generated x the driver at high frequencies (vlich match the power, It is possible that most of the power generated by the driver at high frequencies (which match the power
"wwill become increasingly degenerate withΥ., since most of the signal will be coming from densities above the mean.","will become increasingly degenerate with, since most of the signal will be coming from densities above the mean."
 The degeneracy of aand iin setting the amount of small-scale structure in the fforest is demonstrated in Figure 5.., The degeneracy of and in setting the amount of small-scale structure in the forest is demonstrated in Figure \ref{fig:degeneracy}.
" Here we show the same simulated sections of fforest as in Figure 2,, but with simulation D15 replaced by D10, which has a nearly isothermal temperature-density relation (y~ 1)."," Here we show the same simulated sections of forest as in Figure \ref{fig:curv_example}, but with simulation D15 replaced by D10, which has a nearly isothermal temperature-density relation $\gamma \sim 1$ )."
 The gas at A> is colder in this simulation than it would be in the 1case of y~1.5., The gas at $\Delta > 1$ is colder in this simulation than it would be in the case of $\gamma \sim 1.5$.
" Consequently, at z=2.553 the fforest in B15 and D10 appears nearly identical, even though ddiffers by roughly a factor of two."," Consequently, at $z = 2.553$ the forest in B15 and D10 appears nearly identical, even though differs by roughly a factor of two."
 The similarity is also seen in the curvature values., The similarity is also seen in the curvature values.
" At z=4.586 the fluctuations in the forest are still significantly smoothed out in run D10 as compared to run B15, since the forest at this redshift is more sensitive to gas near the mean density than at lower redshifts."," At $z = 4.586$ the fluctuations in the forest are still significantly smoothed out in run D10 as compared to run B15, since the forest at this redshift is more sensitive to gas near the mean density than at lower redshifts."
" In principle one can take advantage of the fact that lower-column density absorbers tend to arise from density gas in order to constrain both aand citeple.g.,][|schaye2000,theuns2000b.."," In principle one can take advantage of the fact that lower-column density absorbers tend to arise from lower-density gas in order to constrain both and \\citep[e.g.,][]{schaye2000,theuns2000b}."
" In practice, however, this has proven to be extremely challenging, and the error bars on both quantities have remained large."," In practice, however, this has proven to be extremely challenging, and the error bars on both quantities have remained large."
" At z>4 the forest becomes increasingly difficult to separate into discrete absorbers, posing a further obstacle from specifying the full temperature-density relation from the fforest alone."," At $z > 4$ the forest becomes increasingly difficult to separate into discrete absorbers, posing a further obstacle from specifying the full temperature-density relation from the forest alone."
" Rather than attempt to measure both the normalization and the slope of the temperature-density relation, therefore, we use the curvature to measure the temperature at anoverdensity probed by the fforest."," Rather than attempt to measure both the normalization and the slope of the temperature-density relation, therefore, we use the curvature to measure the temperature at an probed by the forest."
" The optimal overdensity, A, which will change with redshift, is empirically determined to be the overdensity at which iis a simple function of(|&|),, regardless ofy."," The optimal overdensity, $\bar{\Delta}$, which will change with redshift, is empirically determined to be the overdensity at which is a simple function of, regardless of."
. Our approach is illustrated in Figure 6.., Our approach is illustrated in Figure \ref{fig:T_delta_vs_curvature}.
 In the left-hand panels we plot aas a function of ffor the various simulation runs., In the left-hand panels we plot as a function of for the various simulation runs.
 The simulations with 7~1.5 form a well-defined relationship between the curvature andTo., The simulations with $\gamma \sim 1.5$ form a well-defined relationship between the curvature and.
". However, runs with different y’ss can be well off of this relationship, particularly at lower redshifts."," However, runs with different s can be well off of this relationship, particularly at lower redshifts."
" For a given aat z~2, the difference in bbetween y~1.5 and y~1.0 can be =10000 K. This is clearly undesirable when iis poorly constrained."," For a given at $z \simeq 2$, the difference in between $\gamma \sim 1.5$ and $\gamma \sim 1.0$ can be $\gtrsim 10\,000$ K. This is clearly undesirable when is poorly constrained."
" In the right-hand panels of Figure 6,, in contrast, we plot the temperature at an optimal overdensity that has been empirically chosen to minimize the variations in the relationship between aandlog(|k|)."," In the right-hand panels of Figure \ref{fig:T_delta_vs_curvature}, in contrast, we plot the temperature at an optimal overdensity that has been empirically chosen to minimize the variations in the relationship between and."
". In this case, iis a nearly one-to-one function of rregardless ofy."," In this case, is a nearly one-to-one function of regardless of."
". These values can be interpreted as characteristic overdensities probed the fforest, though in practice the curvature measurement will be an average over a range of densities."," These values can be interpreted as characteristic overdensities probed the forest, though in practice the curvature measurement will be an average over a range of densities."
" The important point is that we can we can identify a value of A where iis a simple function of the curvature, which allows us to reliably measure the temperature at a density probed by the fforest without knowing the full form of the temperature- relation."," The important point is that we can we can identify a value of $\bar{\Delta}$ where is a simple function of the curvature, which allows us to reliably measure the temperature at a density probed by the forest without knowing the full form of the temperature-density relation."
" ? used a similar approach in their analysis, fitting T' and aat an overdensity where the errors were uncorrelated."," \citet{mcdonald2001b} used a similar approach in their line-fitting analysis, fitting $T$ and at an overdensity where the errors were uncorrelated."
"inversion controls A,,;,. too. aud now the nunber of iterations is indepeucent of resolution.","inversion controls $\lambda_{min}$, too, and now the number of iterations is independent of resolution."
 Towever. direct inversion ix al expensive process and only possible for siiall problems like this one.," However, direct inversion is an expensive process and only possible for small problems like this one."
 Finally. the table also contains the iteration count for approximate inversion of Arp. which is our preferred method for more complex ecolmctiies in 3 dimensious.," Finally, the table also contains the iteration count for approximate inversion of ${\cal
 A}_{FD}$, which is our preferred method for more complex geometries in 3 dimensions."
 It will be explained im detail in the next example., It will be explained in detail in the next example.
 The eieenvalues iu table 1 are estimates obtained by PETSc during the linear solve., The eigenvalues in table \ref{tab:Laplace2D} are estimates obtained by PETSc during the linear solve.
 This solver was developed primarily for elliptic problems m αποΊσα relativity., This solver was developed primarily for elliptic problems in numerical relativity.
 Accordingly we now solve au equation that has been of considerable interest in that field over the last few vears (see c.g. [?.1) aud references therein).," Accordingly we now solve an equation that has been of considerable interest in that field over the last few years (see e.g. \cite{Cook-Choptuik-etal:1993, Cook:2000} and references therein)."
 Readers not familiar with relativity can simply view this problem as another test example for our newsolyert., Readers not familiar with relativity can simply view this problem as another test example for our new.
. We solve . . ↕∪↥⋅↑∐↸∖↕⋯∐⊳⊓∪∐∠⋅∶∠⋅≺⇁≀⋅∙⋅∕∕∙−∙⋟," We solve for the function $\psi=\psi(x,y,z)$."
↔⇂−∶⊀⇂−≺⊽≀⋅∙⋅∕∕∙∶⋟↕↴∖↴⋜↧↨↘↽∐∪↖↖⇁∐∙↻∪↴∖↴↕↑↕↖↽↸∖↕∏∐↸⊳⊓∪∐∙9 . ⋅⋅⋟⋅ ⋜⋯≼↧↑∐↸∖↸⊳∪∐∏⋯↑⋜↕⊓∪∐⋜↧↕≼↧∪⋯⋜↧↕∐↕↴∖↴⊡⋮⋟↖↖⇁↕↑∐↑↖↖↽∪↸∖⊼↸⊳↕↴∖↴↸∖≼↧↴∖↴↻∐↸∖↥⋅↸∖↴∖↴∙ The radii ry» aud centers of the spheres are given.," $A^2=A^2(x,y,z)$ is a known, positive function, and the computational domain is $\mathbbmss{R}^3$ with two excised spheres, The radii $r_{1/2}$ and centers of the spheres are given."
 The function c iust satisfy a Dirichlet boundary couditiou at infinity. aud Robin boundary couditious at the," The function $\psi$ must satisfy a Dirichlet boundary condition at infinity, and Robin boundary conditions at the"
between the FWHM of the lines we identify as Fe compared with those we identify as FeII.,between the FWHM of the lines we identify as Fe compared with those we identify as Fe.
 We computed a simple model which determines the absorption for a thin gas layer which rotates with the disc., We computed a simple model which determines the absorption for a thin gas layer which rotates with the disc.
" This gas layer is assumed to be uniform in height, and is split into multiple sub-layers so the effect of shear can be accounted for."," This gas layer is assumed to be uniform in height, and is split into multiple sub-layers so the effect of shear can be accounted for."
 We use the binary parameters given in ? for these calculations., We use the binary parameters given in \citet{Littlefair08} for these calculations.
 We considered various locations for the absorbing gas layer., We considered various locations for the absorbing gas layer.
" Firstly, we assumed it was very close to the white dwarf, and in the plane of the disc itself."," Firstly, we assumed it was very close to the white dwarf, and in the plane of the disc itself."
" In this configuration, the light which is absorbed originates in the lower hemisphere of the white dwarf and passes through the region of the (inclined) accretion disc which is on the line-of-sight to the observer."," In this configuration, the light which is absorbed originates in the lower hemisphere of the white dwarf and passes through the region of the (inclined) accretion disc which is on the line-of-sight to the observer."
" However, in this configuration our model predicts much broader absorption lines than we observe."," However, in this configuration our model predicts much broader absorption lines than we observe."
" In order for us to obtain line widths that match the observation, we require the observing region to be located in the outer accretion disc."," In order for us to obtain line widths that match the observation, we require the observing region to be located in the outer accretion disc."
" Since the binary is inclined, this also requires the absorbing region to be elevated above the plane of the accretion disc, otherwise it would not intersect the light emitted from the white dwarf."," Since the binary is inclined, this also requires the absorbing region to be elevated above the plane of the accretion disc, otherwise it would not intersect the light emitted from the white dwarf."
 In Figure 8 we plot a cartoon which illustrates the small region of the disc we are probing., In Figure \ref{fig:cartoon} we plot a cartoon which illustrates the small region of the disc we are probing.
 We define a height h above the accretion disc for the gas layer., We define a height $h$ above the accretion disc for the gas layer.
 The h we select also determines a minimum distance rin for theobserved region of the gas layer., The $h$ we select also determines a minimum distance $r_{in}$ for the region of the gas layer.
" Note that the gas layer is not necessarily truncated at rin, but any material closer in does not lie on the line-of-sight to the white dwarf and so is not observed."," Note that the gas layer is not necessarily truncated at $r_{in}$, but any material closer in does not lie on the line-of-sight to the white dwarf and so is not observed."
" We do truncate the gas layer at the outer radius of the accretion disc, which we"," We do truncate the gas layer at the outer radius of the accretion disc, which we"
in the MOPSIC. which is part of the GILDAS software package (http://www.tram.fr/IRAMFR/GILDAS).,"in the MOPSIC, which is part of the GILDAS software package (http://www.iram.fr/IRAMFR/GILDAS)."
" The root mean square level in the maps was typically 8 to 10 mJy/beam with the FWHM beam size being ~10.5"".", The root mean square level in the maps was typically 8 to 10 mJy/beam with the FWHM beam size being $\sim 10.5''$.
" The maps were loaded into AIPS. and flux densities of point sources were extracted using the task ""JMFIT™."," The maps were loaded into AIPS, and flux densities of point sources were extracted using the task “JMFIT”."
 The 870 um observations were carried out in December 2007 using the Large APEX Bolometer Camera (LABOCA: ?)) at the APEX telescope., The 870 $\mu$ m observations were carried out in December 2007 using the Large APEX Bolometer Camera (LABOCA; \citealt{siri09}) ) at the APEX telescope.
 The field of view of the camera (~ 11%) was mapped around each source (except for sources [H.08-0.13 and 41.12-0.11. and 41.12-0.22. 41.16-0.20 and 4].23-0.20. which are close enough to be covered in two separate fields of view) using raster scans.," The field of view of the camera $\sim 11'$ ) was mapped around each source (except for sources 41.08–0.13 and 41.12–0.11, and 41.12–0.22, 41.16–0.20 and 41.23–0.20, which are close enough to be covered in two separate fields of view) using raster scans."
 The time spent per map ranged from 5 to 22 minutes depending on the strength of the source and the weather conditions., The time spent per map ranged from 5 to 22 minutes depending on the strength of the source and the weather conditions.
 The data were reduced using the Bolometer array data Analysis (BoA: Schuller et al..," The data were reduced using the Bolometer array data Analysis (BoA; Schuller et al.,"
 in prep.), in prep.)
 software., software.
 The data reduction strategy was similar to that described by ? and involved flagging of bad pixels. removal of correlated noise. despiking. low-frequency filtering and first-order baselining of the timestream followed by building a map using appropriate weighting.," The data reduction strategy was similar to that described by \citet{schul09} and involved flagging of bad pixels, removal of correlated noise, despiking, low-frequency filtering and first-order baselining of the timestream followed by building a map using appropriate weighting."
 These steps were carried out in several iterations. with a source model from the previous iteration being subtracted from the data. enabling more aggressive application of procedures such as despiking.," These steps were carried out in several iterations, with a source model from the previous iteration being subtracted from the data, enabling more aggressive application of procedures such as despiking."
" The final lc noise level in the maps ranged from 30 to 75 mJy/beam. the FWHM beam size being ~18.2""."," The final $1\sigma$ noise level in the maps ranged from 30 to 75 mJy/beam, the FWHM beam size being $\sim 18.2''$."
" As with the MAMBO data. flux densities of point sources were extracted using the AIPS task ""JMFIT."," As with the MAMBO data, flux densities of point sources were extracted using the AIPS task “JMFIT”."
 In addition to our observations above. we also used existing data at mid-infrared wavelengths to determine the SEDs.," In addition to our observations above, we also used existing data at mid-infrared wavelengths to determine the SEDs."
 While the point source catalogs of the GLIMPSE and 2MASS surveys provided data from 8.0 to 1.24 ym. the point source catalog of 24 jim sources from the MIPSGAL survey has not been released.," While the point source catalogs of the GLIMPSE and 2MASS surveys provided data from 8.0 to 1.24 $\mu$ m, the point source catalog of 24 $\mu$ m sources from the MIPSGAL survey has not been released."
 Hence. we determined 24 jim flux densities directly from the mosaics released by the MIPSGAL team.," Hence, we determined 24 $\mu$ m flux densities directly from the mosaics released by the MIPSGAL team."
 To measure the flux densities. we followed the procedure outlined in the cookbook in the Spitzer Science Center website (http://ssc.spitzer.caltech.edu/dataanalysistools/cookbook/25/).," To measure the flux densities, we followed the procedure outlined in the cookbook in the Spitzer Science Center website (http://ssc.spitzer.caltech.edu/dataanalysistools/cookbook/25/)."
 For sources that were saturated in the MIPS images. we used flux densities from the Galactic Plane survey carried out with the MSX satellite (2)..," For sources that were saturated in the MIPS images, we used flux densities from the Galactic Plane survey carried out with the MSX satellite \citep{egan03}."
 No continuum emission was seen in [4 out of 20 sources at both 3.6 em and 1.3 cm., No continuum emission was seen in 14 out of 20 sources at both 3.6 cm and 1.3 cm.
 At the native resolution. sIX sources were detected at 1.3 em. of which five were detected at 6.9 mm and two were detected at 3.6 em and at 6 cm in the CORNISH survey (?)..," At the native resolution, six sources were detected at 1.3 cm, of which five were detected at 6.9 mm and two were detected at 3.6 cm and at 6 cm in the CORNISH survey \citep{purc08}."
 The latter two sources have spectral indices consistent with optically thin free-free emission. while the three sources detected at 6.9 mm have spectral indices greater than 0.3. which is consistent with optically thick free-free emission.," The latter two sources have spectral indices consistent with optically thin free-free emission, while the three sources detected at 6.9 mm have spectral indices greater than 0.3, which is consistent with optically thick free-free emission."
 The nature of the source 42.70-0.15. that was detected only at 1.3 em is not clear.," The nature of the source 42.70–0.15, that was detected only at 1.3 cm is not clear."
 However. its 3c flux density limit at 3.6 em would be consistent with optically thick emission between these wavelengths.," However, its $3\sigma$ flux density limit at 3.6 cm would be consistent with optically thick emission between these wavelengths."
 One of the concerns in the above analysis is the presence of extended emission that is resolved in our observations., One of the concerns in the above analysis is the presence of extended emission that is resolved in our observations.
area of the active region (sunspot area and plage area) where all values are in wHem (?).,area of the active region (sunspot area and plage area) where all values are in $\mu$ Hem \citep{Chapman97}.
" The equatorial dipole moment of the sunspot group is then assumed to be proportional to the area of the group, Ar, multiplied by the separation between the opposing polarities which we take to be proportional to An?"," The equatorial dipole moment of the sunspot group is then assumed to be proportional to the area of the group, $A_R$, multiplied by the separation between the opposing polarities which we take to be proportional to $A_R^{1/2}$."
" For a sunspot group j in the northern hemisphere and near the equator, with area Ag; and central longitude $;, the axis of the equatorial dipole is orientated in the direction ¢;+90°."," For a sunspot group $j$ in the northern hemisphere and near the equator, with area $A_{R,j}$ and central longitude $\phi_j$, the axis of the equatorial dipole is orientated in the direction $\phi_j \pm 90^{\circ}$."
 The component ofthe dipole in the equatorial plane in direction $ is thus proportional to AR;cos(¢@—$;)," The component ofthe dipole in the equatorial plane in direction $\phi$ is thus proportional to $A_{R,j}^{3/2} \cos(\phi-\phi_j)$."
" For a number of sunspots groups in the northern hemisphere, the component of the resulting dipole moment in the direction ¢ is proportional to Since the sunspot groups in the southern hemisphere have the opposite polarity orientation, the corresponding component of the dipole moment is proportional to Therefore dipole moment of all the spot groups from both hemispheres is proportional to Clearly, m depends on the sunspots which are included in the sum as well as the direction 9."," For a number of sunspots groups in the northern hemisphere, the component of the resulting dipole moment in the direction $\phi$ is proportional to Since the sunspot groups in the southern hemisphere have the opposite polarity orientation, the corresponding component of the dipole moment is proportional to Therefore dipole moment of all the spot groups from both hemispheres is proportional to Clearly, $m$ depends on the sunspots which are included in the sum as well as the direction $\phi$ ."
" We define m(t,v,9) to include in the sum all spots emerging between times { and t+T."," We define $m(t,\tau,\phi)$ to include in the sum all spots emerging between times $t$ and $t+\tau$."
 We are not able to reproduce the longitudes at which sunspots have appeared from 1700 onwards or even those longitudes where nesting has occurred., We are not able to reproduce the longitudes at which sunspots have appeared from 1700 onwards or even those longitudes where nesting has occurred.
 Our aim is to obtain statistical information about the degree of nesting by measuring the degree of non-randomness present in the RGO data., Our aim is to obtain statistical information about the degree of nesting by measuring the degree of non-randomness present in the RGO data.
 We do this by creating three copies of the RGO records., We do this by creating three copies of the RGO records.
 In the first copy we replace the observed longitudes with randomly chosen longitudes from a uniform distribution., In the first copy we replace the observed longitudes with randomly chosen longitudes from a uniform distribution.
 We use the subscript notation man for this dataset., We use the subscript notation $m_{\mathrm{ran}}$ for this dataset.
 The second copy has the longitudes of its spots changed so that they all appear at 0? in the northern hemisphere and 180? in the southern hemisphere., The second copy has the longitudes of its spots changed so that they all appear at $^{\circ}$ in the northern hemisphere and $^{\circ}$ in the southern hemisphere.
 This choice maximizes the equatorial magnetic dipole moment since sunspot groups in opposite hemispheres have opposite polarity orientations in accordance with Hale's law., This choice maximizes the equatorial magnetic dipole moment since sunspot groups in opposite hemispheres have opposite polarity orientations in accordance with Hale's law.
 We use the subscript notation roa for this dataset., We use the subscript notation $m_{\mathrm{ord}}$ for this dataset.
 The third copy retains the observed longitudes and uses the subscript notation mops., The third copy retains the observed longitudes and uses the subscript notation $m_{\mathrm{obs}}$.
" For a given T we measure the amount of nonrandomness, c(t)."," For a given $\tau$ we measure the amount of nonrandomness, $\mathrm{c}(\tau)$."
 We assume that the magnitude of the observed sunspot dipole moment is the magnitude of a linear combination of the random and ordered datasets and seek the c(t) which minimizes the rms differences., We assume that the magnitude of the observed sunspot dipole moment is the magnitude of a linear combination of the random and ordered datasets and seek the $\mathrm{c}(\tau)$ which minimizes the rms differences.
" Figure 6 shows Mops, Mran and Με for the example τ= months."," Figure \ref{fig:M_vs_time} shows $M_{\mathrm{obs}}$, $M_{\mathrm{ran}}$ and $M_{\mathrm{c}}$ for the example $\tau=6$ months."
" The difference between Moy, and M;4, during activity maxima reflects the amount of nesting, the nonrandom longitude distribution."," The difference between $M_{\mathrm{obs}}$ and $M_{\mathrm{ran}}$ during activity maxima reflects the amount of nesting, the nonrandom longitude distribution."
" The similarity of all the curves near the minima is a consequence of the fact that, when there are few spot groups, they are automatically highly ordered."," The similarity of all the curves near the minima is a consequence of the fact that, when there are few spot groups, they are automatically highly ordered."
 We show the dependence of c on r in Figure 7.., We show the dependence of $\mathrm{c}$ on $\tau$ in Figure \ref{fig:C_vs_DeltaT}.
 For different studies the appropriate value of τ will vary., For different studies the appropriate value of $\tau$ will vary.
 For irradiance studies t 21 month would seem to be an appropriate choice because it is the instantaneous clustering of the sunspots which is important., For irradiance studies $\tau=$ 1 month would seem to be an appropriate choice because it is the instantaneous clustering of the sunspots which is important.
" For surface flux transport simulations, τ6 months is more relevant as this is approximatley the time it takes for the emerging flux to be sheared by differential rotation (?,p. 162).."," For surface flux transport simulations, $\tau=6$ months is more relevant as this is approximatley the time it takes for the emerging flux to be sheared by differential rotation \citep[][p.~162]{Schrijver00}. ."
 In this section we consider thearea distribution of sunspot groups based upon the RGO dataset., In this section we consider thearea distribution of sunspot groups based upon the RGO dataset.
" As previously stated, we consider each group only at the time when it has its maximum"," As previously stated, we consider each group only at the time when it has its maximum"
progenitor star. where we normalize their value according to the (vpical radius ancl mass inlerred [rom observations of the few supernovae (SNe) associated with long GRBs.,"progenitor star, where we normalize their value according to the typical radius and mass inferred from observations of the few supernovae (SNe) associated with long GRBs."
 While the jet propagates inside the star (he head dissipates most of ils energv and (he engine must continuously power (he jet in order (o support ils propagation., While the jet propagates inside the star the head dissipates most of its energy and the engine must continuously power the jet in order to support its propagation.
" For a successIul jet breakout. ihe engine working time. /,. must be larger than the breakouttime. /;."," For a successful jet breakout, the engine working time, $t_e$, must be larger than the breakouttime, $t_b$ ."
 If /.«/j the jet [ails lo escape and a regular long GRB is not observed.," If $t_e <
t_b$ the jet fails to escape and a regular long GRB is not observed."
 Once the jet breaks out [rom (he star it produces (he observed emission al large distances from the stellar surface., Once the jet breaks out from the star it produces the observed emission at large distances from the stellar surface.
" Because of relativistie effects (Sari&Piran1997). the observed duration of the prompt 5-ravs emission. ἐς. rellects the Gime that the engine operates after the jet breaksout! The distribution of /, is therefore a convolution of the distributions of the engine activity time and the breakout time combined with cosmological redshift effects."," Because of relativistic effects \citep{SariPiran97} the observed duration of the prompt $\gamma$ -rays emission, $t_\gamma$, reflects the time that the engine operates after the jet breaks: The distribution of $t_\gamma$ is therefore a convolution of the distributions of the engine activity time and the breakout time combined with cosmological redshift effects."
 Under very general conditions. Eq.," Under very general conditions, Eq."
" 2 results in a flat. distribution of /, for durations significantly shorter than the (vpical breakout time.", \ref{tgamma} results in a flat distribution of $t_\gamma$ for durations significantly shorter than the typical breakout time.
 To see it. consider first a single value of [y ancl ignore. for simplicity. cosmological redshift and detector threshokl effects.," To see it, consider first a single value of $t_b$ and ignore, for simplicity, cosmological redshift and detector threshold effects."
" The probability that a GRD has a duration /, equals. in this case. to the probability that the engine work (ime is /4Εν."," The probability that a GRB has a duration $t_\gamma$ equals, in this case, to the probability that the engine work time is $t_\gamma+t_b$."
" Namely. where p, is the probability distribution of observed chivations and p. is the probability distribution of engine working times."," Namely, where $p_\gamma$ is the probability distribution of observed durations and $p_{e}$ is the probability distribution of engine working times."
" Expanding p.(/,+£5) around ἐς=ἰν gives: Thus ον)£2p,(/5)= const For any / that is sullidently smaller (han ἐν.", Expanding $p_e(t_b+t_\g)$ around $t_e=t_b$ gives: Thus $p_e(t_b+t_\g) \approx p_e(t_b) =$ const for any $t_\g$ that is sufficiently smaller than $t_b$ .
 Moreover.if pO.) is a smooth function that does not vary rapidly in the vicinity of /.£z/j over a," Moreover,if $p_e(t_e)$ is a smooth function that does not vary rapidly in the vicinity of $t_e \approx
t_b$ over a"
"velocity resolution of 300 km s~!, matching the one used by Combes et al.,","velocity resolution of 300 km $^{-1}$, matching the one used by Combes et al.,"
 around the central velocity of the line profile., around the central velocity of the line profile.
" Integrating spatially over the emission area and over the velocity range covered, we find an integrated flux density of ~0.72 Jy km s! over a 300 km s! channel, which is about 2.3 times deeper that Combes et al."," Integrating spatially over the emission area and over the velocity range covered, we find an integrated flux density of $\sim0.72$ Jy km $^{-1}$ over a 300 km $^{-1}$ channel, which is about 2.3 times deeper that Combes et al."
 observations., observations.
" Hence, this explains why the CO emission line was undetected previously."," Hence, this explains why the CO emission line was undetected previously."
" The CO emission line from H18214-643 is located ~2.0"" 1 synthesized beam; ~8.8 south-east of the radio-core(about position, and 2.8"" (~12.3 kpc) south-east of the optical center (host galaxy)."," The CO emission line from H1821+643 is located $\sim2.0''$ (about 1 synthesized beam; $\sim8.8$ kpc) south-east of the radio-core position, and $2.8''$ $\sim12.3$ kpc) south-east of the optical center (host galaxy)."
" Based on analysis of PSF subtracted HST optical images, Floydetal.(2004) find a nebulous optical structure."," Based on analysis of PSF subtracted HST optical images, \citet{Floyd2004} find a nebulous optical structure."
" Floyd et al.,"," Floyd et al.,"
" however, point out that such optical feature may correspond to an artifact of the PSF subtraction procedure."," however, point out that such optical feature may correspond to an artifact of the PSF subtraction procedure."
" The CO position coincides with the location of this optical structure, giving support to the reality of such feature."," The CO position coincides with the location of this optical structure, giving support to the reality of such feature."
" Indeed, the bulk of the molecular gas is located at only 0.52”, well within the uncertainty in the position, from an optically faint peak to its north-east (Fig. 1))."," Indeed, the bulk of the molecular gas is located at only $0.52''$, well within the uncertainty in the position, from an optically faint peak to its north-east (Fig. \ref{fig:1}) )."
" Thus, this suggest that the CO emission is possibly associated with either an optically faint but gas-rich companion galaxy that is merging with the giant elliptical host, or with an asymmetric tail-like structure within the host galaxy reminiscent of a previous merger event."," Thus, this suggest that the CO emission is possibly associated with either an optically faint but gas-rich companion galaxy that is merging with the giant elliptical host, or with an asymmetric tail-like structure within the host galaxy reminiscent of a previous merger event."
" Note, however, that the powerful AGN is hosted by a giant elliptical galaxy, and large amounts of molecular gas (M(H2)>5x10? such as that observed here (see below) are rare in Mo)early-type galaxies (Crockeretal.2011)."," Note, however, that the powerful AGN is hosted by a giant elliptical galaxy, and large amounts of molecular gas $_2$ $>5\times10^9$ $_\odot$ ) such as that observed here (see below) are rare in early-type galaxies \citep{Crocker2011}."
". Only 3 out of 56 early-type galaxies in the Virgo cluster have molecular gas masses >10? Mo (Youngetal.2011), being statistically unlikely that the molecular gas observed in H1821+643 is related to the host galaxy."," Only 3 out of 56 early-type galaxies in the Virgo cluster have molecular gas masses $>10^9$ $_\odot$ \citep{Young2011}, being statistically unlikely that the molecular gas observed in H1821+643 is related to the host galaxy."
" Although it is still plausible that the gas could be related to a tail, relic of a past merger, the substantial amounts of molecular gas, its offset location with respect to the elliptical host and the existence of a starburst component in the optical-to-radio SED (see below) suggest that the CO emission actually comes from a different galaxy."," Although it is still plausible that the gas could be related to a tail, relic of a past merger, the substantial amounts of molecular gas, its offset location with respect to the elliptical host and the existence of a starburst component in the optical-to-radio SED (see below) suggest that the CO emission actually comes from a different galaxy."
" If the molecular gas is indeed related to the host galaxy, it would imply a highly asymmetric gas distribution, possibly associated with a previous interaction or spiral arm, in an unusually gas-rich elliptical galaxy, likely related with starburst activity."," If the molecular gas is indeed related to the host galaxy, it would imply a highly asymmetric gas distribution, possibly associated with a previous interaction or spiral arm, in an unusually gas-rich elliptical galaxy, likely related with starburst activity."
" Deep high-resolution IR observations, transparent to obscuration, are necessary to decide between both possibilities."," Deep high-resolution IR observations, transparent to obscuration, are necessary to decide between both possibilities."
" The CO emission is also well aligned with the radio-axis to the south-east suggesting a physical association between the radio-plume emission, the diffuse optical source and the large amounts of molecular gas."," The CO emission is also well aligned with the radio-axis to the south-east suggesting a physical association between the radio-plume emission, the diffuse optical source and the large amounts of molecular gas."
" According to Blundell&Rawlings (2001),, the radio jet appears to have been pointed within 7° within the last 0.3 Gyr."," According to \citet{Blundell2001}, , the radio jet appears to have been pointed within $7^\circ$ within the last 0.3 Gyr."
" Thus, the probability that the CO emission is by chance located within 7° is ~2x7/360 or3.9%,, where the factor 2 accounts for both radio lobes (north and south)."," Thus, the probability that the CO emission is by chance located within $^\circ$ is $\sim2\times7/360$ or, where the factor 2 accounts for both radio lobes (north and south)."
 Such probability is similar to the one found based on the alignment between the radio-axis and the CO position angle., Such probability is similar to the one found based on the alignment between the radio-axis and the CO position angle.
" The CO luminosity (in units of K km s! pc?) can be derived from the line intensity following Solomonetal.(1997) as, where f.AVScodv is the velocity integrated CO flux density in units of Jy km s!, νους is the observing frequency in GHz, and Dy is the luminosity distance in Mpc at the source redshift z."," The CO luminosity (in units of K km $^{-1}$ $^2$ ) can be derived from the line intensity following \citet{Solomon1997} as, where $\int_{\Delta V} S_\mathrm{CO}d v$ is the velocity integrated CO flux density in units of Jy km $^{-1}$, $\nu_\mathrm{obs}$ is the observing frequency in GHz, and $D_\mathrm{L}$ is the luminosity distance in Mpc at the source redshift $z$."
" From our integrated CO flux, we derive Lig=(1.02:0.2)x101? (K km s! pc?) ."," From our integrated CO flux, we derive $L'_\mathrm{CO}=(1.0\pm0.2)\times10^{10}$ (K km $^{-1}$ $^2$ )."
" The luminosity of the CO 1—0 emission line is commonly used to estimate the molecular gas mass, M(H»), by using the relation Χοο]ιοο."," The luminosity of the CO $1-0$ emission line is commonly used to estimate the molecular gas mass, $_2$ ), by using the relation $_2$ $X_\mathrm{CO} L'_\mathrm{CO}$."
" Xco is the CO luminosity to gas mass Μ(Η.)--conversion factor which we assume to be 0.8 Mc (K km s! pc?)-1, as it was found for local IR luminous galaxies (Downes&Solomon1998)."," $X_\mathrm{CO}$ is the CO luminosity to gas mass conversion factor which we assume to be 0.8 $_{\odot}$ (K km $^{-1}$ $^2$ $^{-1}$, as it was found for local IR luminous galaxies \citep{Downes1998}."
. We thus estimate a molecular gas mass of (8.02:1.7)x10? Me for the CO emitting source., We thus estimate a molecular gas mass of $(8.0\pm1.7)\times10^{9}$ $_\odot$ for the CO emitting source.
" This molecular gas mass is at the high end of the mass range found for local ULIRGs, and is typical of luminous submillimeter galaxies and massive disk galaxies at high-redshift (Greveetal.2005;Solomon&VandenBout2005;Daddietal. 2010).."," This molecular gas mass is at the high end of the mass range found for local ULIRGs, and is typical of luminous submillimeter galaxies and massive disk galaxies at high-redshift \citep{Greve2005,Solomon2005,Daddi2010}."
" Due to the relatively low significance of the CO detection, it is not possible to determine the actual geometry of the CO emitting source."," Due to the relatively low significance of the CO detection, it is not possible to determine the actual geometry of the CO emitting source."
"We estimate the dynamical mass based on the CO profile using disk geometry, which is commonly assumed for CO sources.","We estimate the dynamical mass based on the CO profile using disk geometry, which is commonly assumed for CO sources."
" However, we remark that our adoption of a disk geometry is not well established for our CO source and we use"," However, we remark that our adoption of a disk geometry is not well established for our CO source and we use"
through thiu-cloud cover aud ~ ssecine.,through thin-cloud cover and $\sim$ seeing.
 The observations were performed uxiug the nod-and-shufüe mode. with the RLOO erating centered atSIODA:: this setup was chosen to minimize the effect of moon ilhuuination aud to obtain good sky subtraction.," The observations were performed using the nod-and-shuffle mode, with the R400 grating centered at; this setup was chosen to minimize the effect of moon illumination and to obtain good sky subtraction."
 To avoid second-order contamination frou blue lieht in the red part of these spectra. the OCG blocking filter was used.," To avoid second-order contamination from blue light in the red part of these spectra, the OG515 blocking filter was used."
" Iu addition. multi-object spectroscopic515. (MOS) observations of galaxies in the Abell 3827 field were obtained on 2007 Sep 7 (UT). during dark time. in photometric couditious and with seeing between 07 aud 079,"," In addition, multi-object spectroscopic (MOS) observations of galaxies in the Abell 3827 field were obtained on 2007 Sep 7 (UT), during dark time, in photometric conditions and with seeing between $\farcs$ 7 and $\farcs$ 9."
 Two masks were observed with 0775 slit widths. using the D600 erating centered at5220A.," Two masks were observed with $\farcs$ 75 slit widths, using the B600 grating centered at."
. Divine the observations. all spectra were dithered to account for the eps iu the GMOS CCDs.," During the observations, all spectra were dithered to account for the gaps in the GMOS CCDs."
 Spectra were reduced aud calibrated in the standard manner for cach mode., Spectra were reduced and calibrated in the standard manner for each mode.
 We have measured radial velocitics for 67 galaxies using cross-correlation techniques or emission-liue fitting., We have measured radial velocities for 67 galaxies using cross-correlation techniques or emission-line fitting.
 Of these GT. galaxies. 56 are within £2500 kin | of the cluster redshift.," Of these 67 galaxies, 56 are within $\pm\,2500$ km $^{-1}$ of the cluster redshift."
 Using the biaweight estimator for location aud scale (Beersatal.LO90).. we calculate au average velocity for the cluster of 4 155 lans. | aud a lue-ofsight velocity dispersion of 11124125 kin with 55 member galaxies.," Using the bi-weight estimator for location and scale \citep{beers90}, we calculate an average velocity for the cluster of $\,\pm\,$ 155 km $^{-1}$ and a line-of-sight velocity dispersion of $\,\pm\,$ 125 km $^{-1}$, with 55 member galaxies."
 The virial mass inside a raciuδα of 0.3 ot Ape is (3.86!058)&103123AL... with uncertainties at the confidence intervals.," The virial mass inside a radius of 0.3 $^{-1}$ Mpc is $(3.86_{-0.28}^{+0.46}) \times 10^{14}\,h^{-1}$, with uncertainties at the confidence intervals."
 Iu the cluster core. we confirm membership for four οKER he central elliptical galaxies.," In the cluster core, we confirm membership for four of the central elliptical galaxies."
 There are also two stars 1 he foreground (see Fie., There are also two stars in the foreground (see Fig.
 1)., 1).
 Three of the core galaxie: jive similar colors (see Table 1) aud them velocities arc consistent with the average cluster velocity (within  300 aus 1), Three of the core galaxies have similar colors (see Table 1) and their velocities are consistent with the average cluster velocity (within $\sim$ 300 km $^{-1}$ ).
 Galaxy N. Eis slightly bluer than N.1. N.2 aud N3 aud las a measured radial velocity ~ 1000 kii s! üeher than the svstemic velocity.," Galaxy N.4 is slightly bluer than N.1, N.2 and N.3 and has a measured radial velocity $\sim$ 1000 km $^{-1}$ higher than the systemic velocity."
 Galaxy N.5 lay not )o in the cluster. but the backeround.," Galaxy N.5 may not be in the cluster, but the background."
 Thauks to the superb nuage quality delivered ly GAIOS and the Cemini South telescope. we lave been able to detect several strounglv-dleused features around the central cD galaxy in the core of the cluster.," Thanks to the superb image quality delivered by GMOS and the Gemini South telescope, we have been able to detect several strongly-lensed features around the central cD galaxy in the core of the cluster."
 A thin. extended tangential arc (D.l in Fie.," A thin, extended tangential arc (B.1 in Fig."
" 1) is seen atf ~ sonthecast from the central galaxy. with apparent leneth ~11"" aud thickness ~O77."," 1) is seen at $\sim$ south-east from the central galaxy, with apparent length $\sim\,11\arcsec$ and thickness $\sim\,0\farcs7$."
" Surrounding the ceutral core galaxies is a second aud more prominent arc featurea laghh-imaenified. rine-shaped configuration of four iniages around the central cD galaxy; with a radius of ~δν,"," Surrounding the central core galaxies is a second and more prominent arc feature—a highly-magnified, ring-shaped configuration of four images around the central cD galaxy, with a radius of $\sim$."
 Oulv three of the arc components are visible in the inset color image in Fie., Only three of the arc components are visible in the inset color image in Fig.
 d (À.l. À.2 and A.3).," 1 (A.1, A.2 and A.3)."
 The fourth image (A. lin Fig., The fourth image (A.4 in Fig.
 1) is between galaxies N.1 and N.2 and is obscured by the halo of the cD galaxy., 1) is between galaxies N.1 and N.2 and is obscured by the halo of the cD galaxy.
 All four features can be seen in detail in the bottoniright panel in Fig., All four features can be seen in detail in the bottom-right panel in Fig.
 d (sce the caption for details)., 1 (see the caption for details).
 These features are very renuiuiscent of the lensed images seen imn C1002111651 (Colleyetal.1996).. although the source ealaxy iu the present case is closer than 2~1 (see below).," These features are very reminiscent of the lensed images seen in Cl0024+1654 \citep{colley1996}, although the source galaxy in the present case is closer than $z\sim1$ (see below)."
 The redshifts of the lensed sources were determined οι loug-sht observations taken on 2007 Nov 15-16. using the RI50 erating centered at7150À.," The redshifts of the lensed sources were determined from long-slit observations taken on 2007 Nov 15-16, using the R150 grating centered at."
. For the feature with the rine-shaped configuration (system A). we derived a redshift for one of the images located north-west of the cluster core (A.2).," For the feature with the ring-shaped configuration (system A), we derived a redshift for one of the images located north-west of the cluster core (A.2)."
 Two emission lines (fa and A6582À)) are seen in the spectrüuni for A.2: unfortunatelv[NII] uo other lines were detected. due to liebt contamination from the ¢D halo.," Two emission lines $H\alpha$ and [N $ \lambda\,6583$ ) are seen in the spectrum for A.2; unfortunately no other lines were detected, due to light contamination from the cD halo."
 From the two eiiission lines. the estimated redshift for A.2 is 2=0.20113+ 0.00072.," From the two emission lines, the estimated redshift for A.2 is $z = 0.20443\,\pm\,0.00073$ ."
 The sviuaetric distribution of mages in svstem A aud the rine-shaped structure connecting arcs A.l. A.2 and A.B (and possibly A.l) suggest that the images are of the same source galaxy.," The symmetric distribution of images in system A and the ring-shaped structure connecting arcs A.1, A.2 and A.3 (and possibly A.4) suggest that the images are of the same source galaxy."
 Therefore. we assume that all four images come from the same source and have the sale redshift.," Therefore, we assume that all four images come from the same source and have the same redshift."
 This assunptiou should be taken with caution. without spectroscopy of the remaining svsteni A images for coufirmation. but is supported by our nodeling. discussed in the next section.," This assumption should be taken with caution, without spectroscopy of the remaining system A images for confirmation, but is supported by our modeling, discussed in the next section."
 The subtraction of light from the halo of the cD ealaxy. as well as subtraction of all galaxies aud the stars frou he cluster core. highlehts the ring-shaped structure connecting arcs A.l. A.2. À3 aud possibly À.1 (see Fig.," The subtraction of light from the halo of the cD galaxy, as well as subtraction of all galaxies and the stars from the cluster core, highlights the ring-shaped structure connecting arcs A.1, A.2, A.3 and possibly A.4 (see Fig."
 1)., 1).
 The existence of this «structure only ~157.+ kpe from he ceuter of the cD galaxy gives us a unique opportunity o study the dynamics aud mass distribution of a large BCG at unprecedeutedly small scales.," The existence of this structure only $\sim 15\,h^{-1}$ kpc from the center of the cD galaxy gives us a unique opportunity to study the dynamics and mass distribution of a large BCG at unprecedentedly small scales."
 Using several enission-liues (JOM) A3727A.. 273. [OTI] AX1959.5007 A)) seen in the spectrum of the south-cast aueential arc. D.1. we derived a redshift of ;=0.10825+ 1.00072.The relatively long. thin appearance of this arc suggests that the source galaxy lies very close to the ner old caustic. assimnuius the leus cluster has an elliptical," Using several emission-lines ([OII] $\lambda\,3727$, $H\!\beta$, [OIII] $\lambda\lambda\,4959,5007$ ) seen in the spectrum of the south-east tangential arc, B.1, we derived a redshift of $z=0.40825\,\pm\,0.00072$ .The relatively long, thin appearance of this arc suggests that the source galaxy lies very close to the inner fold caustic, assuming the lens cluster has an elliptical"
"5 lists the nearest ijonizius source iu cach cluster along with its estimated spectral type aud. Lxiuau continui ποπτν, the projected separation between the head of cach tail aud. this stax. the expected fiux of Lxauan coutimmiun radiatiou iu photous per square centimeter per second. and the predicted flux of Pan cCluission under the assumption that each comet tail head is optically thick to τα. coutimmiun radiation.","3 lists the nearest ionizing source in each cluster along with its estimated spectral type and Lyman continuum luminosity, the projected separation between the head of each tail and this star, the expected flux of Lyman continuum radiation in photons per square centimeter per second, and the predicted flux of $\alpha$ emission under the assumption that each comet tail head is optically thick to Lyman continuum radiation."
 This estimate ignores the attenuation of the Lyman contimmun radiation by dust., This estimate ignores the attenuation of the Lyman continuum radiation by dust.
 Finally. the last column in Table 3 gives the ratio of the observational lint on the Pan emission derived from Table 2 divided by the expected flux under the assumption that cach star with a dust tail contains a hemispherical. high censity core that is opaque to the Lyman continua.," Finally, the last column in Table 3 gives the ratio of the observational limit on the $\alpha$ emission derived from Table 2 divided by the expected flux under the assumption that each star with a dust tail contains a hemispherical, high density core that is opaque to the Lyman continuum."
 Iu these estimates. we use the projected distance between the nearest ionizing star and each tail. au assumption that will underestimate the true separation aud overestimate the incident Lyman contiuuu flux.," In these estimates, we use the projected distance between the nearest ionizing star and each tail, an assumption that will underestimate the true separation and overestimate the incident Lyman continuum flux."
" However, we also ignore the contributions of other OD stars in cach WIT region to the total Lyman continui radiation field."," However, we also ignore the contributions of other OB stars in each HII region to the total Lyman continuum radiation field."
 This assumption will tend ο result iu an underestimnate of the incident radiation field., This assumption will tend to result in an underestimate of the incident radiation field.
 The estimated incident radiation fields are xobablv uncertain to about a factor of two., The estimated incident radiation fields are probably uncertain to about a factor of two.
 Thus. the xedieted enmission measures and Pao fluxes are ordor-ofimagnitucde estimates;," Thus, the predicted emission measures and $\alpha$ fluxes are order-of-magnitude estimates."
 However. as shown by Table 3. he predicted fluxes are many orders of mmaenitude higher han the observational limits.," However, as shown by Table 3, the predicted fluxes are many orders of magnitude higher than the observational limits."
 The observations indicate the Spitser-detected tails are deficieut in Pao emission. by about a factor of one willion compared to a simple model in which cach dust ail contains a plioto-ablatiug core of dense material., The observations indicate the -detected tails are deficient in $\alpha$ emission by about a factor of one million compared to a simple model in which each dust tail contains a photo-ablating core of dense material.
 As discussed above. the emission measure. aud therefore the Huxes of livdrogen recombination lines are to first order iudepeudoeut of the radius of the ionization front.," As discussed above, the emission measure, and therefore the fluxes of hydrogen recombination lines are to first order independent of the radius of the ionization front."
 They depend ou the incident Lyinan continui radiatiou field. and inversely as the distance between the object and source of ionization squared.," They depend on the incident Lyman continuum radiation field, and inversely as the distance between the object and source of ionization squared."
 Thus. a set of optically thick. neutral coucdensations ought to increase in Pao flux as the inverse square of the separation.," Thus, a set of optically thick, neutral condensations ought to increase in $\alpha$ flux as the inverse square of the separation."
 If these dust tails were simular to the proplvds in the Orion Nebula. they would be orders of mmaeuitude brighter than the observed limits.," If these dust tails were similar to the proplyds in the Orion Nebula, they would be orders of magnitude brighter than the observed limits."
 The abseuce of detectable lydrogen recombination line fluxes might be explained if the disks are optically thin to Lyman continu radiation and the eas-to-dust ratio is significautlv lower than iu the ISM., The absence of detectable hydrogen recombination line fluxes might be explained if the disks are optically thin to Lyman continuum radiation and the gas-to-dust ratio is significantly lower than in the ISM.
 For an assunaed radius of the circtuustelar environment. rp=10? AU. the critical density above which the gas is opaque to the fiux of Έλαια continua radiation is npcOF.apri)!:2=yeDSK+104Fi;1/2Fs where £4o is in units of 1012.ple Lyman continuum photons stem? and rs is in units of 10? AU.," For an assumed radius of the circumstellar environment, $r_I = 10^3$ AU, the critical density above which the gas is opaque to the flux of Lyman continuum radiation is $n_I > (3F/\alpha_B r_I)^{1/2} = 2.8\times 10^4F_{12}^{1/2} r_3^{-1/2}$ where $F_{12}$ is in units of $10^{12}$ Lyman continuum photons ${\rm s^{-1} cm^{-2}}$ and $r_3$ is in units of $10^3$ AU."
 This iuplies au cussion ineasure. EALzmος10Ευcum)Spe.," This implies an emission measure, $EM \approx 6 \times 10^6 F_{12}~{\rm cm^{-6} pc}$."
 In contrast. our Pao flux limits (Table 2) tmuiply enuission measures of EMz1673 cu“pe for the Pao. flux limit from the NGC 22Li τα απο EMz637 aupe for the Pan fixes from the tail iu IC 1396.," In contrast, our $\alpha$ flux limits (Table 2) imply emission measures of $EM \approx 1673$ ${\rm cm^{-6} pc}$ for the $\alpha$ flux limit from the NGC 2244 tail and $EM \approx 637$ ${\rm cm^{-6} pc}$ for the $\alpha$ fluxes from the tail in IC 1396."
 Thus. for an incident. Lyman continu flux Z4». the Pan deficit would be between 2 and 12<107.," Thus, for an incident Lyman continuum flux $F_{12}$, the $\alpha$ deficit would be between $2$ and $12 \times 10^{-5}$."
 Table 3 gives the actual deficits for cach tail eiveu the projected separations from the nearest O-star. the Lyman coutinnuu bIuniuostv of the star (also in Table 3). aud the flux lauits of Table 2.," Table 3 gives the actual deficits for each tail given the projected separations from the nearest O-star, the Lyman continuum luminosity of the star (also in Table 3), and the flux limits of Table 2."
 If the ciremustellar cuviromment is optically thin to +ιο Lyman coutimmun aud has a coustaut density auk radius py. the upper limut for the Wvdrogen mass is ALdsEM|pe])er; * which implies masses of 1.7<OAD. aud 2.7«10.OAD. for our two limiting cases of EM © 637 cusSpe and L673 cur“pe respectively.," If the circumstellar environment is optically thin to the Lyman continuum and has a constant density and radius $r_I$, the upper limit for the Hydrogen mass is $M\approx \mu m_H (11.18 EM[pc])^{1/2}r_I^{5/2}$ , which implies masses of $1.7 \times 10^{-6} {\rm M_{\odot}}$ and $2.7 \times 10^{-6} {\rm M_{\odot}}$ for our two limiting cases of EM $\approx$ 637 ${\rm cm^{-6} pc}$ and 1673 ${\rm cm^{-6} pc}$ respectively."
 These are very conservative upper bounds since observations show that circustellar euvironnients have /fatteuc disk-like geometrics., These are very conservative upper bounds since observations show that circumstellar environments have flattened disk-like geometries.
 The eusssiou measure and derive density limits would be au order of magnitude lower for a disk ecometry., The emission measure and derived density limits would be an order of magnitude lower for a disk geometry.
 The mass of a typical protoplanetary disk is about 0.1-0.001 M... (Andrews&Williams2005)., The mass of a typical protoplanetary disk is about 0.1-0.001 ${\rm M_{\odot}}$ \citep{Andr05}.
. ον inferred gas mass upper uits for the optically thin case are at least two orders of mmaenitude lower than the eas niasses expected in a typical T Tai disk. implying eas to dust ratios 107 to LO! times lower than iu the ISM," Our inferred gas mass upper limits for the optically thin case are at least two orders of magnitude lower than the gas masses expected in a typical T Tauri disk, implying gas to dust ratios $10^2$ to $10^4$ times lower than in the ISM."
 Balogetal.(2006) modeled the gram propertics and tail morphologics for the three evaporating disks discussed in this paper., \citet{Balo06} modeled the grain properties and tail morphologies for the three evaporating disks discussed in this paper.
 The models utilized optical constants for astronomical silicates and assumed a sinele eyain size., The models utilized optical constants for astronomical silicates and assumed a single grain size.
 It was found that the tail morplologics (faintucss at 5 gan and the apparent lengths) placed strong coustraints on the erain sizes., It was found that the tail morphologies (faintness at 8 $\mu$ m and the apparent lengths) placed strong constraints on the grain sizes.
 The minim sizes were of order ~ (0.01 gan in radius - smaller grains cnutted too strongly at 8 jan to be iucluded iu the tails in significant nuubers, The minimum sizes were of order $\sim$ 0.01 $\mu$ m in radius - smaller grains emitted too strongly at 8 $\mu$ m to be included in the tails in significant numbers.
 Grains above ~ 1 jnu in radius tended to be too cold to fit the tail leueths., Grains above $\sim$ 1 $\mu$ m in radius tended to be too cold to fit the tail lengths.
 The spectra predicted by these models for the objects iu IC 1396 aud NGC 2261 are shown iu Figures 2 and 30 for comparison with our IRS spectra., The spectra predicted by these models for the objects in IC 1396 and NGC 2264 are shown in Figures \ref{fig:1396IRSSL} and \ref{fig:2264IRSSL} for comparison with our IRS spectra.
 The fit is satisfactory at wavelengths louger than ~ 12 sau. At shorter wavelengths. there is significautf enuüssion frou the “comet head” both from the stellar photosphere aud οι an excess bridging from 2.1 jun to 12 gan and," The fit is satisfactory at wavelengths longer than $\sim$ 12 $\mu$ m. At shorter wavelengths, there is significant emission from the ""comet head"" both from the stellar photosphere and from an excess bridging from 2.4 $\mu$ m to 12 $\mu$ m and"
as to the origin of the slope discrepancy.,as to the origin of the slope discrepancy.
 We have shown that the Jurcsik relatious lead to RR Lyrae radii that are iu good agreement with both Daacde-Wessclik radii aud with theoretical radii., We have shown that the Jurcsik relations lead to RR Lyrae radii that are in good agreement with both Baade-Wesselink radii and with theoretical radii.
 Ou the other haud. there exists a strong discrepancy betweeu the slope of the theoretical L irelation aud the slope of the emipirical Juresik relation.," On the other hand, there exists a strong discrepancy between the slope of the theoretical – relation and the slope of the empirical Jurcsik relation."
 We have reviewed the physical an numerical uncertainties that euter the theoretical calcilations., We have reviewed the physical and numerical uncertainties that enter the theoretical calculations.
 The discrepancy exists already at the level of purely radiative modeling., The discrepancy exists already at the level of purely radiative modelling.
 We have further shown that our 1iuodel equations for turbulent couvection do not alor the slope ofthe relation despiteL the large nunuber (8) of adjustable piruneters (αν) that wei have atf οιr disposal., We have further shown that our model equations for turbulent convection do not alter the slope of the relation despite the large number (8) of adjustable parameters $\alpha$ 's) that we have at our disposal.
 Consequently. apart frou the uncertainties 1iberent ina LD recipe for turbulent convection. it appears that the discrepancy is not caused by a deficiency of the theoretical models.," Consequently, apart from the uncertainties inherent in a 1D recipe for turbulent convection, it appears that the discrepancy is not caused by a deficiency of the theoretical models."
 We note that the same slope cdiscrepaucy is nuplicit in older aud more recent independent calculations (TugeleOO Then 1972. Bono 1997).," We note that the same slope discrepancy is implicit in older and more recent independent calculations (Tuggle Iben 1972, Bono 1997)."
 The derivation of the empirical relations makes use of a Ixurucz color - ttransformation., The derivation of the empirical relations makes use of a Kurucz color - transformation.
 By computing the behavior of color versus temperature over the pulsation evcle we lave shown that the use of static cuvelopes is in fact a very good approximation. and therefore cannot be the culprit.," By computing the behavior of color versus temperature over the pulsation cycle we have shown that the use of static envelopes is in fact a very good approximation, and therefore cannot be the culprit."
 Finally. the shape of the evolutionary tracks. through their potential avoidance of certain regious of the instability strip cannot be responsible cither.," Finally, the shape of the evolutionary tracks, through their potential avoidance of certain regions of the instability strip cannot be responsible either."
"rotation rate is favored (vsini,24.522.4 km s! . allowing a broader range of spin-orbit angles (\=192.0*/5* degrees).","rotation rate is favored $v\sin i_\star = 4.5\pm 2.4$ km $^{-1}$ ), allowing a broader range of spin-orbit angles $\lambda =
192.0_{-8.7}^{+15.6}$ degrees)."
 These results are also illustrated by the dotted lines in Figure ].., These results are also illustrated by the dotted lines in Figure \ref{fig:hat14}.
" One possible reason for the smaller result for vsiné, is differential rotation: the RM effect depends on the rotation rate over the range of latitudes spanned by the transit chord. which may differ from the spectroscopically-estimated equatorial rotation rate."," One possible reason for the smaller result for $v\sin i_\star$ is differential rotation: the RM effect depends on the rotation rate over the range of latitudes spanned by the transit chord, which may differ from the spectroscopically-estimated equatorial rotation rate."
 The near-grazing transit of HAT-P-14. in particular. may produce the most extreme possible difference.," The near-grazing transit of HAT-P-14, in particular, may produce the most extreme possible difference."
 Further observations with a higher signal-to-noise ratio might be able to identify the specific signal of differential rotation (Gaudi Winn 2007)., Further observations with a higher signal-to-noise ratio might be able to identify the specific signal of differential rotation (Gaudi Winn 2007).
 The case of HAT-P-4 was more complicated. partly because our RV observations revealed a third body in the system.," The case of HAT-P-4 was more complicated, partly because our RV observations revealed a third body in the system."
 Figure 2. illustrates that a single planet on a circular orbit no longer provides a satisfactory description of the data., Figure \ref{fig:hat4_orb} illustrates that a single planet on a circular orbit no longer provides a satisfactory description of the data.
 The best-fitting model has V7=925.7 with 18 degrees of freedom., The best-fitting model has $\chi^2 = 925.7$ with 18 degrees of freedom.
 The residuals have an rms of 15.9 m s and are highly correlated. with almost all of the most recent RVs lying above the model curve.," The residuals have an rms of 15.9 m $^{-1}$ and are highly correlated, with almost all of the most recent RVs lying above the model curve."
 Allowing the orbit to be eccentric reduces \7 to 906.0. but the residuals still have an rms of 15 m s! and show the same pattern.," Allowing the orbit to be eccentric reduces $\chi^2$ to 906.0, but the residuals still have an rms of 15 m $^{-1}$ and show the same pattern."
 A much better fit is obtained when the circular orbit is supplemented by a constant acceleration 7., A much better fit is obtained when the circular orbit is supplemented by a constant acceleration $\dot{\gamma}$.
 This model. illustrated in the lower panel of Figure 2.. gives \7=137.0 and more randomly scattered residuals with an rms residual of 6.8 m s''.," This model, illustrated in the lower panel of Figure \ref{fig:hat4_orb}, , gives $\chi^2=137.0$ and more randomly scattered residuals with an rms residual of 6.8 m $^{-1}$."
 An MCMC analysis gives ~0.0026 m s! day.," An MCMC analysis gives $\dot{\gamma}= 0.0246\pm
0.0026$ m $^{-1}$ $^{-1}$."
 The constant acceleration may represent the RV variation due to à companion star or planet whose orbit is longer than the three-year span of our observations., The constant acceleration may represent the RV variation due to a companion star or planet whose orbit is longer than the three-year span of our observations.
 Assuming it is à low-mass body on a nearly circular orbit. we may set ~~GM.sini./az. etcgiving an order-of-magnitude constraint However. given the limited time sampling of our data it is also possible that the companion has a shorter period.," Assuming it is a low-mass body on a nearly circular orbit, we may set $\dot{\gamma} \sim GM_c \sin i_c/a_c^2$, giving an order-of-magnitude constraint However, given the limited time sampling of our data it is also possible that the companion has a shorter period."
 We have not found any compelling two-planet models. but we are continuing to gather additional RV data and will report elsewhere on the results.," We have not found any compelling two-planet models, but we are continuing to gather additional RV data and will report elsewhere on the results."
 The pertinent conclusions for this study are (1) there is a third body in the system. and (2) the mass and orbital parameters of HAT-P-4b are subject to systematic errors due to the unknown influence of the third body on the RV data.," The pertinent conclusions for this study are (1) there is a third body in the system, and (2) the mass and orbital parameters of HAT-P-4b are subject to systematic errors due to the unknown influence of the third body on the RV data."
 For the latter reason. in our analysis of the transit data we did not employ a prior constraint on K.. as we did for HAT-P-14.," For the latter reason, in our analysis of the transit data we did not employ a prior constraint on $K_\star$, as we did for HAT-P-14."
 Another difference in our analysis is that we fitted the available photometric data along with the transit-night RV data. because the new data offer better constraints on the transit ephemeris. depth. duration. and partial duration.," Another difference in our analysis is that we fitted the available photometric data along with the transit-night RV data, because the new data offer better constraints on the transit ephemeris, depth, duration, and partial duration."
 Our photometric model was taken from the Transit Light Curve project (see. e.g.. Holman et al.," Our photometric model was taken from the Transit Light Curve project (see, e.g., Holman et al."
 2006. Winn et al.," 2006, Winn et al."
 2009c)., 2009c).
 In brief. we used the Mandel Agol (2002) formulas for a quadratic limb-darkening law. as implemented by Pall(2008).," In brief, we used the Mandel Agol (2002) formulas for a quadratic limb-darkening law, as implemented by Páll(2008)."
 The linear coefficient was allowed to vary freely. and the," The linear coefficient was allowed to vary freely, and the"
Our main source for the construction of our target. List was?.. with the Edinburgh Cape objects comingfrom 2 and private communications.,"Our main source for the construction of our target list was\citet{Kilkenny88}, with the Edinburgh Cape objects comingfrom \citet{Kilkenny97} and private communications."
 Over the course of our study we discovered a number of sources which have been misidentified as sdB stars in these catalogues. which we listin Table S...," Over the course of our study we discovered a number of sources which have been misidentified as sdB stars in these catalogues, which we listin Table \ref{tab:misclassifications}."
 At the time of writing they are all still [listed as sdDs in the SIMDAD astronomical database. although the misidentification has previously. been reported for six of these eleven. objects.," At the time of writing they are all still listed as sdBs in the SIMBAD astronomical database, although the misidentification has previously been reported for six of these eleven objects."
 For the other five we give our new classifications. obtained using the same spectral fitting techniquesdescribed in Section 4.1..," For the other five we give our new classifications, obtained using the same spectral fitting techniquesdescribed in Section \ref{sec:tefflogg}."
 In this paper we present a large number of racial velocity measurements of bright sdB stars., In this paper we present a large number of radial velocity measurements of bright sdB stars.
 The aim of this project was to detect binary systems through variations in these racial velocity. measurements. anc hence determine the orbital parameters of those svstems.," The aim of this project was to detect binary systems through variations in these radial velocity measurements, and hence determine the orbital parameters of those systems."
 Εις is a continuation ofthework begun in ? and 7.., This is a continuation ofthework begun in \citet{Maxted02} and \citet{MoralesRueda03}.
 We presented: a total of 28 new binary systems with their parameters. in all cases determining the orbital period to much better than 5 per cent.," We presented a total of 28 new binary systems with their parameters, in all cases determining the orbital period to much better than $5$ per cent."
 We determined elfective temperatures ancl surface eravities For 13 of the sdDs in these svstems., We determined effective temperatures and surface gravities for 13 of the sdBs in these systems.
 The parameters are consistent with the sdBs being extreme horizontal branch stars (ELLB) with two exceptions. PCOO34|186 and DOG1244|113. which we classify as post-ELB stars.," The parameters are consistent with the sdBs being extreme horizontal branch stars (EHB) with two exceptions, PG0934+186 and PG1244+113, which we classify as post-EHB stars."
 As well as the 28) solved systems we presented measurements for 108 other sdBs., As well as the 28 solved systems we presented measurements for 108 other sdBs.
 20 of these show strong signs of binaritv. but our data is insullicient as vet to constrain the svstem parameters.," 20 of these show strong signs of binarity, but our data is insufficient as yet to constrain the system parameters."
 The remaining LOO stars show no significant racial velocity variations and are likely to be either single sdBs or binary svstems in which the orbital period is long enough to push the radial velocity. variations below our detection threshold., The remaining 100 stars show no significant radial velocity variations and are likely to be either single sdBs or binary systems in which the orbital period is long enough to push the radial velocity variations below our detection threshold.
 Our best.estimate for the binary fraction is 46 56 per cent., Our bestestimate for the binary fraction is $46$ – $56$ per cent.
 Finally. we note that none ofthe binaries we report in this paper show the near-infrared. colours which indicate the presence of a G-Ix-tvpe companion star.," Finally, we note that none ofthe binaries we report in this paper show the near-infrared colours which indicate the presence of a G-K-type companion star."
 Phe companion stars in our solved systems are therefore likely either white ewarfs or M dwarls., The companion stars in our solved systems are therefore likely either white dwarfs or M dwarfs.
 There is some evidence for a cool companion in the spectra of 32 of the remaining sdBs. including 5 which show radial velocity variations.," There is some evidence for a cool companion in the spectra of 32 of the remaining sdBs, including 5 which show radial velocity variations."
 These are strong candidates for future stud., These are strong candidates for future study.
 CAIC and PRAT are supported under grant ST00250041 from the Science and Technology Facilities Council (SLEC)., CMC and TRM are supported under grant ST/F002599/1 from the Science and Technology Facilities Council (STFC).
 The results presented in this paper are based on observations made with the Isaac. Newton Telescope operated on the island. of La Palma by the Isaac Newton Group in the Spanish Observatorio del Roque de los Muchachos of the Institutio cle Astrofisica de Canarias and observations made with the 1.9 m telescope operated by the South African Astronomical Observatory., The results presented in this paper are based on observations made with the Isaac Newton Telescope operated on the island of La Palma by the Isaac Newton Group in the Spanish Observatorio del Roque de los Muchachos of the Institutio de Astrofisica de Canarias and observations made with the 1.9 m telescope operated by the South African Astronomical Observatory.
 his research has made use of «λος Astrophysics Data System Bibliographic Services and the SIMDAD database. operated at CDS. Strasbourg. France.," This research has made use of NASA's Astrophysics Data System Bibliographic Services and the SIMBAD database, operated at CDS, Strasbourg, France."
 This publication makes use of data products rom the Two Micron All Sky Survey. which is a joint oject of the University of Massachusetts ancl the Infrared 'pocessing anc Analysis Center/Calilormia Institute of Technology. funded by the National Aeronautics and Space Administration ancl the National Science Foundation.," This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation."
 We hank the referee for helpful comments., We thank the referee for helpful comments.
scaling of angular momentum (ransler with interaction parameters.,scaling of angular momentum transfer with interaction parameters.
 However. it should be pointed out Chat for more distant encounters. the error bar in (he angular momentum change is larger for 10 000 particles Chan 50 000 because of the reduced resolution al (he outer edge.," However, it should be pointed out that for more distant encounters, the error bar in the angular momentum change is larger for 10 000 particles than 50 000 because of the reduced resolution at the outer edge."
 The dise mass is mg; = 0.001 M..., The disc mass is $m_{disc}$ = 0.001 $M _{\sun}$.
 A gap of 10 AU between the star and the inner edee of the disc has been assumed in order (o save computer time ancl (o avoid additional complex caleulation of direct star/disc interactions., A gap of 10 AU between the star and the inner edge of the disc has been assumed in order to save computer time and to avoid additional complex calculation of direct star/disc interactions.
 To a first approximation the resulis can be generalized to larger disc sizes. since the disc mass is verv small and the effects of the disc particles on the stus very limited.," To a first approximation the results can be generalized to larger disc sizes, since the disc mass is very small and the effects of the disc particles on the stars very limited."
 The same holds for the density distribution: the results can to converted to an arbitrary. density profile O(r) bv simple multiplication with O(7)/O4(r)2zr.Nr. where Ar is the radial extent of each annulus of material aud Ο4(70) the 1/r-density profile used in above simulations.," The same holds for the density distribution: the results can to converted to an arbitrary density profile $\Theta(r)$ by simple multiplication with $\Theta(r)/\Theta_{-1}(r) 2\pi r \Delta r$, where $\Delta r$ is the radial extent of each annulus of material and $\Theta_{-1}(r)$ the $1/r$ -density profile used in above simulations."
 Fie., Fig.
 laa) shows a tvpical example of the change of the angular momentum per particle (or per mass unit: (hese are equivalent in (his case. since all simulation particles are of equal mass) as a [unetion of the distance from the central star in one of our simulations.," \ref{fig:ang_rad_simple}a a) shows a typical example of the change of the angular momentum per particle (or per mass unit: these are equivalent in this case, since all simulation particles are of equal mass) as a function of the distance from the central star in one of our simulations."
 Outside the original dise radius (£004U) only a gain in angular momentum can occur. since (here were obviously no particles in (liis region belore the encounter.," Outside the original disc radius (100AU) only a gain in angular momentum can occur, since there were obviously no particles in this region before the encounter."
 Inside (is initial radius there is a broad region where each particle loses angular momentum and still further inside is a small region where angular momentum is gained (but to a much lesser extent)., Inside this initial radius there is a broad region where each particle loses angular momentum and still further inside is a small region where angular momentum is gained (but to a much lesser extent).
 Close to the central star (not shown here) the angular momentum remains unaltered., Close to the central star (not shown here) the angular momentum remains unaltered.
 Apparently the inner angular moment eain has nol been seen previously: its phivsical origin remains vel io be determined., Apparently the inner angular momentum gain has not been seen previously: its physical origin remains yet to be determined.
 The change of angular momentum of the dise is usually measured by siumming over all particles still bound to the central star after the encounter., The change of angular momentum of the disc is usually measured by summing over all particles still bound to the central star after the encounter.
 Although the gain of angular momentum per particle outside the initial disc radius is larger than the loss insicle. the total amounts to a loss of angular momentum in (he disc. since there are many more particles inside the disc than outside (see Fig.," Although the gain of angular momentum per particle outside the initial disc radius is larger than the loss inside, the total amounts to a loss of angular momentum in the disc, since there are many more particles inside the disc than outside (see Fig."
 1bb))., \ref{fig:ang_rad_simple}b b)).
 sunmmine over all bound particles - as is usually done - might therefore not be appropriate. eiven (hat the problem at hand is how to lose angular momentum from (he central region of the disce.," Summing over all bound particles - as is usually done - might therefore not be appropriate, given that the problem at hand is how to lose angular momentum from the central region of the disc."
 The particles outside the original dise do not really [ligure in (his process. apart from compensating to some degree for the loss in the more central areas of the disc by (heir," The particles outside the original disc do not really figure in this process, apart from compensating to some degree for the loss in the more central areas of the disc by their"
a gas molecule to escape vields estimates of mass loss rates.,a gas molecule to escape yields estimates of mass loss rates.
 We take the mass loss rate to be (7): where £45; is the extreme UV. (XUV) Ilux from the star (wavelengths from 0.1 to 100 nm). evaluated at. the planets orbital distance.," We take the mass loss rate to be \citep{2007A&A...472..329E}: where $F_{xuv}$ is the extreme UV (XUV) flux from the star (wavelengths from 0.1 to 100 nm), evaluated at the planet's orbital distance."
" A, is the planets radius. and G is the gravitational constant."," $R_p$ is the planet's radius, and G is the gravitational constant."
 This equation. assumes that the planets optical cross-section in the NUV. is sd. as suggested by 2? ancl ?..," This equation assumes that the planet's optical cross-section in the XUV is $\pi R_p^2$ , as suggested by \citet{2007A&A...472..329E} and \citet{2004Icar..170..167Y}."
 The factors ο and Ai; (each of which may range from 0 to 1) are discussed The parameter ec represents the fraction of the incoming energv that is carried away by the escaping gas., The factors $\epsilon$ and $K_{tide}$ (each of which may range from 0 to 1) are discussed The parameter $\epsilon$ represents the fraction of the incoming energy that is carried away by the escaping gas.
 Some of the incoming energy mav drive chemistry or be lost through radiative cooling of the planet. rather than driving escape.," Some of the incoming energy may drive chemistry or be lost through radiative cooling of the planet, rather than driving escape."
 In these cases. c is less than 1 (????).. ," In these cases, $\epsilon$ is less than 1 \citep{2007ApJ...658L..59H, 2007Icar..187..358H, 2004Icar..170..167Y, 2009A&A...506..399L}."
"For example. estimates of the mass loss rate for LL 209458 b. suggest 4M,diτοLOM efs(eig... 7)). whieh corresponds to €0.1."," For example, estimates of the mass loss rate for HD 209458 b suggest $\frac{dM_p}{dt} \sim 4\times10^{10}$ g/s, \citealt{2008SSRv..139..437Y}) ), which corresponds to $\epsilon \sim 0.4$."
 On the other hand. an analysis of atmospheric escape from Venus suggests 0.15 may. be a more appropriate value for ο (2)..," On the other hand, an analysis of atmospheric escape from Venus suggests $0.15$ may be a more appropriate value for $\epsilon$ \citep{1996JGR...10126039C}."
 2. indicate that ο is probably less than 0.6 for hot Jupiters but will depend on the host star's activity and the planets atmospheric composition. among other things.," \citet{2009A&A...506..399L} indicate that $\epsilon$ is probably less than 0.6 for hot Jupiters but will depend on the host star's activity and the planet's atmospheric composition, among other things."
 Moreover. vaporization and removal of Mercury exosphere involve processes that may be dilferent. from. evaporations of gas giant planets (?).. and evolution of the evaporating planct’s composition may result in a time-dependent e-value (?2)..," Moreover, vaporization and removal of Mercury's exosphere involve processes that may be different from evaporation of gas giant planets \citep{2007SSRv..132..433K}, and evolution of the evaporating planet's composition may result in a time-dependent $\epsilon$ -value \citep{2009ApJ...693...23M}."
 However. Equation | provides a reasonable estimate of evaporation of planetary mass by NUV tux. which plays an important. if not dominant. role for Mercury. as well as easeous planets.," However, Equation \ref{eqn:dMpdt} provides a reasonable estimate of evaporation of planetary mass by XUV flux, which plays an important, if not dominant, role for Mercury as well as gaseous planets."
 Here. we consider a range of time-invariant € values: 0. Q.1. 0.25. 0.5. and. 1.," Here, we consider a range of time-invariant $\epsilon$ values: 0, 0.1, 0.25, 0.5, and 1."
 We include ο=0 to compare orbital evolution rates with and without mass loss., We include $\epsilon = 0$ to compare orbital evolution rates with and without mass loss.
 Note that we do not include any possible dependence ofο on the solid. planet’s composition. even though mass loss may be more rapid if the planet were more enriched in volatiles.," Note that we do not include any possible dependence of $\epsilon$ on the solid planet's composition, even though mass loss may be more rapid if the planet were more enriched in volatiles."
 The detailed: modeling required to quantify that ellect is beyond the scope of this studs., The detailed modeling required to quantify that effect is beyond the scope of this study.
 For planets close to their host stars. gas molecules may escape from a planet even. with a velocity significantly less than the usual escape velocity. because they only have to reach the Roche lobe.," For planets close to their host stars, gas molecules may escape from a planet even with a velocity significantly less than the usual escape velocity because they only have to reach the Roche lobe."
 This reduction in the required velocity results in an increased escape rate for a given rate of energv input., This reduction in the required velocity results in an increased escape rate for a given rate of energy input.
 In Equation l.. Ay... represents this reduction in required escape energy (2).," In Equation \ref{eqn:dMpdt}, $K_{tide}$ represents this reduction in required escape energy \citep{2007A&A...472..329E}."
" Ayia. is given by: where £ is the ratio of the Roche radius (=(AL,/3A1, a) to R,.", $K_{tide}$ is given by: where $\xi$ is the ratio of the Roche radius $=(M_p/3M_*)^{1/3} a$ ) to $R_p$ .
 For a planet far from its host star. where the Roche radius is much greater than A. νε approaches one. indicating no enhancement of the escape rate.," For a planet far from its host star, where the Roche radius is much greater than $R_p$, $K_{tide}$ approaches one, indicating no enhancement of the escape rate."
 For LED 209458 b. with e=0.047 AU. ων=0.65.," For HD 209458 b, with $a = 0.047$ AU, $K_{tide} = 0.65$."
 As tides bring a planet closer and closer to its host star. Aia decreases to reflect the reduction in required escape In emploving this formulation for the mass loss rate. we treat its elects as if the mass loss is isotropic. that ds. as if the flux of escaping gas is the same in all directions.," As tides bring a planet closer and closer to its host star, $K_{tide}$ decreases to reflect the reduction in required escape In employing this formulation for the mass loss rate, we treat its effects as if the mass loss is isotropic, that is, as if the flux of escaping gas is the same in all directions."
 In reality. gas may escape preferentially through the L1 Lagrange point between the planet and star. where the net acceleration is close to zero (?2).. which may allect the planets orbital evolution.," In reality, gas may escape preferentially through the L1 Lagrange point between the planet and star, where the net acceleration is close to zero \citep{2003ApJ...588..509G}, which may affect the planet's orbital evolution."
 However. incorporating this directionalitv may require a 3-D. hiydrodyvnamic model and so is bevond the scope of our study.," However, incorporating this directionality may require a 3-D hydrodynamic model and so is beyond the scope of our study."
 We discuss this point more in Section We also include the evolution of the stellar SUV lux in our model., We discuss this point more in Section We also include the evolution of the stellar XUV flux in our model.
 The stellar XUV flux powers mass loss. but as stars age. typically their rotation rates drop with time. which reduces their magnetic activity anc NUN radiation(6.9... 7).," The stellar XUV flux powers mass loss, but as stars age, typically their rotation rates drop with time, which reduces their magnetic activity and XUV radiation, \citealt{1972ApJ...171..565S}) )."
" 7. studied the evolution of the NUV flux from C-type stars like ColtoT-7 and suggested fy. evolves according to: where A=20.7 eres | cm7. 34=123. fer, ds the stars age in billions of vears and e is in AU."," \citet{2005ApJ...622..680R} studied the evolution of the XUV flux from G-type stars like CoRoT-7 and suggested $F_{xuv}$ evolves according to: where $\lambda = 29.7$ ergs $^{-1}$ $^{-2}$, $\beta = 1.23$, $t_{Gyr}$ is the star's age in billions of years and $a$ is in AU."
 This equation is an empirical fit to data for observations of stars with ages ereater than 0.1 Cyr ancl therefore may. only apply ord >0.1 Gaver., This equation is an empirical fit to data for observations of stars with ages greater than $0.1$ Gyr and therefore may only apply for $t > 0.1$ Gyr.
 However. planets in our solar svstem (and. »esumably. in others) developed: atmospheres in less than a few tens of millions of vears after formation (?)..," However, planets in our solar system (and, presumably, in others) developed atmospheres in less than a few tens of millions of years after formation \citep{1986JGR....91..291A}."
 Also. [ore §=0.1 Gyr. ColtoT-7s XUV tux was probably somewhat [larger than at later times.," Also, before $t = 0.1$ Gyr, CoRoT-7's XUV flux was probably somewhat larger than at later times."
 Thus. atmospheric oss mav also have been significant at earlier times than considered. here.," Thus, atmospheric loss may also have been significant at earlier times than considered here."
 In. this sense. our results provide a lower imit to the total mass lost by ColtoT-7 b. However. that earlier. period of time was much shorter than the time span we consider. and so the total mass lost may have en not much greater than suggested by our calculations.," In this sense, our results provide a lower limit to the total mass lost by CoRoT-7 b. However, that earlier period of time was much shorter than the time span we consider, and so the total mass lost may have been not much greater than suggested by our calculations."
 Without more data on stellar NUV at carly times. stopping our models at /=O.1 Gar is probably an adequate approximation for Gur CoRkol-7 b receives sullicient insolation that vaporization of its rocky surface may produce a tenuous— atmosphere(e.g...," Without more data on stellar XUV at early times, stopping our models at $t = 0.1$ Gyr is probably an adequate approximation for our CoRoT-7 b receives sufficient insolation that vaporization of its rocky surface may produce a tenuous atmosphere, \citealt{2009ApJ...703L.113S}) )."
 7)) ? suggested that the rate of vaporization of rocky material may be orders of magnitude larger than the atmospheric mass loss rate. sullicientlvy rapid to continually resupply the atmosphere. so in our model. we allow evaporation to continue even after a model planet loses its entire original H/Ie To determine the planets radius after. the original Il/lle envelope is lost. we use Equations 7 S from ?.. along with corrections given in ?.. ," \citet{2009arXiv0907.3067V} suggested that the rate of vaporization of rocky material may be orders of magnitude larger than the atmospheric mass loss rate, sufficiently rapid to continually resupply the atmosphere, so in our model, we allow evaporation to continue even after a model planet loses its entire original H/He To determine the planet's radius after the original H/He envelope is lost, we use Equations 7 8 from \citet{2007ApJ...659.1661F}, along with corrections given in \citet{2007ApJ...668.1267F}. ."
These equations provide mass-racdius relations for solid. planets mace of ice and rock or rock and iron. which requires us to make some assumption about the planet's composition in order to determine its radius. given à mass.," These equations provide mass-radius relations for solid planets made of ice and rock or rock and iron, which requires us to make some assumption about the planet's composition in order to determine its radius, given a mass."
 Figure 1. shows the range of compositions. allowed bv ColtoT-7 b's observed. mass and. radius based. on the equations from ?.., Figure \ref{fig:plot_Rp_v_Mp_mf} shows the range of compositions allowed by CoRoT-7 b's observed mass and radius based on the equations from \citet{2007ApJ...659.1661F}.
 hat. figure shows that. if made of ice and rock. CololT-7 b's ice mass fraction. (EME) may lie between 0.01 and0.212.," That figure shows that, if made of ice and rock, CoRoT-7 b's ice mass fraction (IMF) may lie between 0.01 and0.212."
 I£ made of rock and iron. ColtoT-7 bs rock mass fraction (RAIL) maylie in the range 0.719 to 0.983.," If made of rock and iron, CoRoT-7 b's rock mass fraction (RMF) maylie in the range 0.719 to 0.983."
 These latter numbers are similar to compositions for ColtoT-7 bproposed by 7.. , These latter numbers are similar to compositions for CoRoT-7 bproposed by \citet{2009arXiv0907.3067V}. .
In our models. we consider the allowed. minimum. middle ancl maximum values forboth," In our models, we consider the allowed minimum, middle and maximum values forboth"
the K-K relation to eq.,the K-K relation to eq.
 I. and is plotted in Fig.," 4, and is plotted in Fig."
 1D. again together with the corresponding function for eraphite. as obtained by Taft aud Philipp.," 1B, again together with the corresponding function for graphite, as obtained by Taft and Philipp."
 Finally. the extinction eficicney for both materials was obtained from where A ds the current waveleusth aud the graiu shape parameter is set at L—1/3. asstuning the erains are roughly spherical.," Finally, the extinction efficiency for both materials was obtained from where $\lambda$ is the current wavelength and the grain shape parameter is set at $L=1/3$, assuming the grains are roughly spherical."
 Oulv the 7 resonance part of Qa. which is of interest for comparison with IS spectra. is plotted in Fie.," Only the $\pi$ resonance part of $Q/a$, which is of interest for comparison with IS spectra, is plotted in Fig."
 1C. for eraphene (the narrowest peak. 0.75 eV or 0.6 pau 1) aud eraphite (the widest peak. 1.5 eV or L2 yan 1).," 1C, for graphene (the narrowest peak, 0.75 eV or 0.6 $\mu$ $^{-1}$ ) and graphite (the widest peak, 1.5 eV or 1.2 $\mu$ $^{-1}$ )."
 Achuittedly. the Caussian function was introduced in eq.," Admittedly, the Gaussian function was introduced in eq."
 3 as an to 6tain the desired feature at the right location with the right width aud the right intensity., 3 as an to obtain the desired feature at the right location with the right width and the right intensity.
" However. since this IS feature is so sensitive to the dielectric fuction in its spectral range. aud if our graphene model is indeed. valid. our procedure can also be considered as a guide to the theory of eraphene. based on accurate astronolucal ποσοοί»,"," However, since this IS feature is so sensitive to the dielectric function in its spectral range, and if our graphene model is indeed valid, our procedure can also be considered as a guide to the theory of graphene, based on accurate astronomical measurements."
 This is welcome. suce even theories of the joiut density of states must rely on empirical coupling cucreics," This is welcome, since even theories of the joint density of states must rely on empirical coupling energies"
because the innermost planet has reached the edge of the inner cavity. where it remains trapped. or because two adjacent bodies migrate differentially.,"because the innermost planet has reached the edge of the inner cavity, where it remains trapped, or because two adjacent bodies migrate differentially."
 In some simulations. resonant interaction and the associated eccentricity growth resulted in close encounters and collisions between bodies. resulting in the formation of à more massive planet.," In some simulations, resonant interaction and the associated eccentricity growth resulted in close encounters and collisions between bodies, resulting in the formation of a more massive planet."
 At later times. once the more chaotic phase of evolution is over. the system behaves more quiescently. and generally reaches a stationary state with each remaining body in resonance with its closest neighbours on non migrating orbits due to planet trapping at the cavity edge.," At later times, once the more chaotic phase of evolution is over, the system behaves more quiescently, and generally reaches a stationary state with each remaining body in resonance with its closest neighbours on non migrating orbits due to planet trapping at the cavity edge."
 In two of the three simulations. we found that disc dispersal led to strong scattering and collisions between bodies. resulting in eventually only one planet being left in the system.," In two of the three simulations, we found that disc dispersal led to strong scattering and collisions between bodies, resulting in eventually only one planet being left in the system."
 In the remaining case. we found that three planets could exist stably in a three body mean motion resonance over very long times.," In the remaining case, we found that three planets could exist stably in a three body mean motion resonance over very long times."
 The results of our simulations indicate. that resonant systems of planets with masses in the ~ few—Earth mass range may be common in close binary systems., The results of our simulations indicate that resonant systems of planets with masses in the $\sim$ few–Earth mass range may be common in close binary systems.
 Indeed. the inner edge of a circumbinary dise plays the role of a barrier. which raises the possibility that two bodies become resonantly trapped.," Indeed, the inner edge of a circumbinary disc plays the role of a barrier, which raises the possibility that two bodies become resonantly trapped."
 This resonance trapping. however. usually prevents the two planets from approaching each other so that they are unable to merge to form a larger body.," This resonance trapping, however, usually prevents the two planets from approaching each other so that they are unable to merge to form a larger body."
 When nore than two bodies are introduced we found that the resulting scattering and close encounters could lead to planetary growth. however. but we have been unable to run sufficient numbers of simulations to. provide reliable statistical information about the range of outcomes.," When more than two bodies are introduced we found that the resulting scattering and close encounters could lead to planetary growth, however, but we have been unable to run sufficient numbers of simulations to provide reliable statistical information about the range of outcomes."
 A number of outstanding issues remain when considering planet formation in. eircumbinary— disces., A number of outstanding issues remain when considering planet formation in circumbinary discs.
 The first is understanding where the planetary cores can form due to planetesimal accretion., The first is understanding where the planetary cores can form due to planetesimal accretion.
 The formation of an eccentric disc around a binary system. and the presence of the binary itself. can increase the velocity dispersion to a value where accretion does not occur except far out in the disc.," The formation of an eccentric disc around a binary system, and the presence of the binary itself, can increase the velocity dispersion to a value where accretion does not occur except far out in the disc."
 The second is understanding the growth of cores once they have formed. either through continued planetesimal accretion. or through giant impacts.," The second is understanding the growth of cores once they have formed, either through continued planetesimal accretion, or through giant impacts."
 The migration of the core toward the inner cavity may actually place it in a region where planetesimal accretion is reduced due to increased velocity dispersion. and at present there have been no simulations of giant impact growth which include the gas disc.," The migration of the core toward the inner cavity may actually place it in a region where planetesimal accretion is reduced due to increased velocity dispersion, and at present there have been no simulations of giant impact growth which include the gas disc."
 If a giant core can be formed. however. then it is likely that a gas giant planet will result.," If a giant core can be formed, however, then it is likely that a gas giant planet will result."
 As it grows. a giant planet should leave the edge of the cavity and undergo type II migration (e.g. Lin Papaloizou 1993: Nelson al.," As it grows, a giant planet should leave the edge of the cavity and undergo type II migration (e.g. Lin Papaloizou 1993; Nelson al."
 2000). until it becomes trapped into the 4:] resonance with the binary.," 2000), until it becomes trapped into the 4:1 resonance with the binary."
 This appears to be the most likely fate of a giant planet migrating in a cireumbinary disc (Nelson 2003). although trapping is not necessarily permanent and simulations indicate a finite probablity of the planet being ejected from the system due to close encounters with the binary.," This appears to be the most likely fate of a giant planet migrating in a circumbinary disc (Nelson 2003), although trapping is not necessarily permanent and simulations indicate a finite probablity of the planet being ejected from the system due to close encounters with the binary."
 The evolution of the planet during its growth from a core into a gas glant. however. has not been explored. and this will be the subject of a future paper.," The evolution of the planet during its growth from a core into a gas giant, however, has not been explored, and this will be the subject of a future paper."
It is clear that (AR) satisfies the coudition Inserting (2)> in thetrace equation (8)). the Ricci scalar is obtained as Then we arrive at the solutious of Eq.(7)) where eds a constaut.,"It is clear that $h(R)$ satisfies the condition Inserting $\varphi(R)$ in thetrace equation \ref{eq8}) ), the Ricci scalar is obtained as Then we arrive at the solutions of \ref{eq7}) ) where $c$ is a constant."
 We eau see our solutions are il agreement with the exact solutions (9))., We can see our solutions are in agreement with the exact solutions \ref{eq9}) ).
" Also oue can check neglecting R. C, aud ΟΠΠΟ) iu Eq.(1.. 21) js reasonable."," Also one can check neglecting $R$, $G_{\mu\nu}$ and $h(R)/\varphi(R)$ in \ref{eq1}, \ref{eq2}) ) is reasonable."
" For this model (1)=ΤΠ,", For this model $\varphi(R)=1/R$.
 Solving trace equation (8)) and field equatious (7)) we obtain aud where M is a constant., Solving trace equation \ref{eq8}) ) and field equations \ref{eq7}) ) we obtain and where $M$ is a constant.
 Therefore the space time metric for enipty space in this model is This generalized eravitational poteutial has two terius.," Therefore the space time metric for empty space in this model is We can see, the generalized Newtonian potential is This generalized gravitational potential has two terms."
" The first terii is the standard Newtoniui potential aud the secoud term make à coustant acceleration, |ve2[. which is indepeudeut of the mass of star."," The first term is the standard Newtonian potential and the second term make a constant acceleration, $+\sqrt{\epsilon/24}$, which is independent of the mass of star."
 Tu (SaffariandRaliwar2005) this metric is used to address the Pioneers anomalous., In \citep{safari} this metric is used to address the Pioneer's anomalous.
 Based on equivalence between £8) exavitv aud Braus-Dicke theory with w=0. it was argued that this theory Is Inconsistent with solar svstei tests (Chiba2003).," Based on equivalence between $f(R)$ gravity and Brans-Dicke theory with $\omega=0$, it was argued that this theory is inconsistent with solar system tests \citep{a5}."
. Tudeed by this approach the Post-Newtouian parameter is found as 5ppy=1/2 while the measurements indicate ~ppy=1|(2.242.3)ς107 (Bertottietal.2003)., Indeed by this approach the Post-Newtonian parameter is found as $\gamma_{PPN}=1/2$ while the measurements indicate $\gamma_{PPN}=1+(2.1\pm 2.3)\times10^{-5}$ \citep{cas}.
. Also we must uote that using equivalence between f(R) exavity aud scalar tensor eravity oue ca- find models which are consistent with the solar svstem tests., Also we must note that using equivalence between $f(R)$ gravity and scalar tensor gravity one can find models which are consistent with the solar system tests.
 This consistency can be made by eiving the scalar a biel mass or exploiting the so-called clameleocoe effect(AlotaandBarrow9001IKhouryaud.Woeltiua2007).," This consistency can be made by giving the scalar a high mass or exploiting the so-called chameleon \citep{New1, r9,
r10, r11}."
. However. when oue is using equivalence betwee- F(R) evavity aud scalar tensor gravitv. the continuity of scalar Ποια or its equivalent. the Ricci scalar. at the latter boundary is crucial couditiou which is not the case iun Eimstein gravitv.," However, when one is using equivalence between $f(R)$ gravity and scalar tensor gravity, the continuity of scalar field or its equivalent, the Ricci scalar, at the matter boundary is crucial condition which is not the case in Einstein gravity."
 But iu this work we dout adopt the continuity of Ricci scalar for solving the field equations., But in this work we don't adopt the continuity of Ricci scalar for solving the field equations.
 Dustead. we suppose that when jp teuds to zero we arrive at the Einsteim eravity.," Instead, we suppose that when $\mu$ tends to zero we arrive at the Einstein gravity."
 Thus we find a solution for 1/R model which is radically different from other solutions m (Erickeeketal.," Thus we find a solution for $1/R$ model which is radically different from other solutions in \citep{us1, us2}."
 For this model we have where ACR) fulfills the condition 0.," For this model we have where $h(R)$ fulfills the condition $\mathop {\lim }\limits_{R \to
0} \left[ {h(R)/\varphi (R)} \right] = 0$ ."
 Solving Eqs.(7..8)) we obtain where a?=ULI and M is a coustaut.," Solving \ref{eq7}, \ref{eq8}) ) we obtain where $\alpha^{3}=4/147$ and $M$ is a constant."
 Therefore the metric for space time 1s, Therefore the metric for space time is
"and sini=1.50,7.40 My, andMgsini0.84, 1.57 Mj, respectively;Udry-20011,, Μι; ?), (7.22,16.2M5;?), (0.71,2.20and4.45My; and (0.56,1.88My; ?).","and \\citep[$M_2 \sin i = 1.50$,
7.40~M$_{\rm J}$ , and $M_2 \sin i =0.84$, 1.57 $_{\rm J}$ , respectively;][], \citep[0.16,
0.38 M$_{\rm J}$ , \citep[7.22, 16.2~M$_{\rm
 J}$, \citep[0.71, 2.20 and 4.45~M$_{\rm J}$ and \citep[0.56, 1.88~M$_{\rm J}$ ."
" The inversion process is only able to treat these systems under the hypothesis that the orbits of the different planets in a given system arenot coplanar, since Eq. (4))"," The inversion process is only able to treat these systems under the hypothesis that the orbits of the different planets in a given system are coplanar, since Eq. \ref{Eq:integral2}) )"
 to hold requires random orbital inclinations., to hold requires random orbital inclinations.
 The case of coplanar and non-coplanar orbits are discussed separately in the remaining of this section., The case of coplanar and non-coplanar orbits are discussed separately in the remaining of this section.
" Figure 2 compares the solutions U(M2) obtained from the Lucy-Richardson algorithm (after 2 and 20 iterations, denoted Ψο and Ψορ, respectively) and from the formal solution of Abel’s integral equation with smoothing lengths hoy; and 2 hogy on ®(Y) (the corresponding solutions are denoted V,,.. and V5,,..)."," Figure \ref{Fig:2} compares the solutions $\Psi(M_2)$ obtained from the Lucy-Richardson algorithm (after 2 and 20 iterations, denoted $\Psi_2$ and $\Psi_{20}$, respectively) and from the formal solution of Abel's integral equation with smoothing lengths $h_{\rm opt}$ and 2 $h_{\rm opt}$ on $\Phi(Y)$ (the corresponding solutions are denoted $\Psi_{h_{\rm opt}}$ and $\Psi_{2h_{\rm opt}}$ )."
" The solutions 2&om the two methods basically agree with each other, although solutions with different degrees of smoothness may be obtained with each method."," The solutions from the two methods basically agree with each other, although solutions with different degrees of smoothness may be obtained with each method."
" On the one hand, Woo and V,,. exhibit high-frequency fluctuations that may be traced back to the statistical fluctuations in the imput data."," On the one hand, $\Psi_{20}$ and $\Psi_{h_{\rm opt}}$ exhibit high-frequency fluctuations that may be traced back to the statistical fluctuations in the input data."
" This can be seen by noting that the peaks Jpresent in V,,.. correspond in fact tothe high-frequency Jfluctuations already present in $,,, (Fig. 1)).", This can be seen by noting that the peaks present in $\Psi_{h_{\rm opt}}$ correspond in fact tothe high-frequency fluctuations already present in $\Phi_{h_{\rm opt}}$ (Fig. \ref{Fig:1}) ).
 These [fluctuations should thus not be given much credit., These fluctuations should thus not be given much credit.
" The same explanation holds true for Φορ, since it was argued it Sect."," The same explanation holds true for $\Psi_{20}$, since it was argued in Sect."
" 3 that the solutions W, resulting from a large ""number of iterations tend to match 4 at increasingly small scales (i.e., higher frequencies) where statistical jfluctuations become dominant."," \ref{Sect:LR} that the solutions $\Psi_r$ resulting from a large number of iterations tend to match $\Phi$ at increasingly small scales (i.e., higher frequencies) where statistical fluctuations become dominant."
" On the other hand, V2 aad Wo), are much smoother, and are probably better -matches to the actual distribution."," On the other hand, $\Psi_{2}$ and $\Psi_{2h_{\rm opt}}$ are much smoother, and are probably better matches to the actual distribution."
 The local maximum Jaround Mz~1 is very likely however an artifact of |the strong detection bias against low-mass companions., The local maximum around $M_2 \sim 1$ $_{\rm J}$ is very likely however an artifact of the strong detection bias against low-mass companions.
 The most striking feature of the Ψ(λ49) distribution ?üisplayed in Fig., The most striking feature of the $\Psi(M_2)$ distribution displayed in Fig.
" 2 is its decreasing character, reaching zero for the first time around M»=10 Mj, and in any case well before 13.6 My."," \ref{Fig:2} is its decreasing character, reaching zero for the first time around $M_2 = 10$ $_{\rm J}$, and in any case well before 13.6 $M_{\rm J}$."
" The latter value, corresponding to the minimum stellar mass for igniting Deuterium, does not in any way mark the transition between giant planets and brown dwarfs, as sometimes proposed."," The latter value, corresponding to the minimum stellar mass for igniting Deuterium, does not in any way mark the transition between giant planets and brown dwarfs, as sometimes proposed."
" That transition, which is thus likely to occur at smaller masses, must rely instead on the different mechanisms governing the formation of planets and brown dwarfs."," That transition, which is thus likely to occur at smaller masses, must rely instead on the different mechanisms governing the formation of planets and brown dwarfs."
" Another argument favouring a giant-planet/brown-dwarf transition mass smaller than 13.6 Mj is provided e.g., by the observation offree-floating (and thus most likely stellar) objects with masses probably smaller than 10 My in the o Orionis star cluster (?).."," Another argument favouring a giant-planet/brown-dwarf transition mass smaller than 13.6 $_{\rm J}$ is provided e.g., by the observation of (and thus most likely ) objects with masses probably smaller than 10 $_{\rm J}$ in the $\sigma$ Orionis star cluster \citep{Zapatero-2000}."
" The W(M2»)distribution nevertheless clearly exhibits a tail of objects clustering around My, due to (Mpsini=11.5 Mj), (14.3 Mj), (15.0 My) and (16.2 Mj)."," The $\Psi(M_2)$distribution nevertheless clearly exhibits a tail of objects clustering around $M_2 \sim
15$ $_{\rm J}$ , due to $M_2 \sin i =
11.5$ $_{\rm J}$ ), (14.3 $_{\rm J}$ ), (15.0 $_{\rm J}$ ) and (16.2 $_{\rm J}$ )."
" It would be interesting to investigate whether these systems differ fromthose with smaller masses in some identifiable way (periods, eccentricities,"," It would be interesting to investigate whether these systems differ fromthose with smaller masses in some identifiable way (periods, eccentricities,"
predicted simulation temperatures at 2=5 for the harder ionising spectra.,predicted simulation temperatures at $z=5$ for the harder ionising spectra.
 Note that the agreement would be even worse if the heating contribution from X-rays were included in the simulations., Note that the agreement would be even worse if the heating contribution from X-rays were included in the simulations.
 These results are thus consistent. with a predonminance |sources with relatively soft. (a2 3) ionising spectra during hydrogen reionisation. and also with an epoch of reionisation (most likely driven by quasars) which was not Cully underway. until lower redshift (c.g. 2)).," These results are thus consistent with a predominance of sources with relatively soft $\alpha \geq 3$ ) ionising spectra during hydrogen reionisation, and also with an epoch of reionisation (most likely driven by quasars) which was not fully underway until lower redshift (e.g. \citealt{McQuinn2009}) )."
 We therefore conclude that ifà population of sources with rather hard spectra. such as mini-quasars (?)) or population-IHll stars (2)) were responsible for reionising hydrogen. their 'ontribution must be either sub-dominant at all recdshifts r confined. predominantly at carly times (2> 9). such ju there has been sullicient time for the ICM temperature o cool and doubly ionised helium to recombine by z76.," We therefore conclude that if a population of sources with rather hard spectra, such as mini-quasars \citealt{Madau04}) ) or population-III stars \citealt{Bromm01}) ) were responsible for reionising hydrogen, their contribution must be either sub-dominant at all redshifts or confined predominantly at early times $z\geq 9$ ), such that there has been sufficient time for the IGM temperature to cool and doubly ionised helium to recombine by $z\simeq 6$."
 This is not surprising as population-LLL stars are believed o be present at z«9. but. compared. to population-LL stars. in. negligide numbers (scc e.g. 77)).," This is not surprising as population-III stars are believed to be present at $z<9$, but, compared to population-II stars, in negligible numbers (see e.g. \citealt{Tornatore.Ferrara.Schneider_2007,Maio_etal_2010}) )."
 7? have also recently pointed: out that relative metal abundances in he LGAL sugges population-LL stars. produced the bulk of ivdrogen ionising photons during reionisation.," \cite{Becker12}
 have also recently pointed out that relative metal abundances in the IGM suggest population-II stars produced the bulk of hydrogen ionising photons during reionisation."
 Similarly. although mini-quasars have been investigated by a number ofauthors as possible sources of ionising photons. the general agreement is that their contribution is not dominant (seo e.g. 7?)).," Similarly, although mini-quasars have been investigated by a number of authors as possible sources of ionising photons, the general agreement is that their contribution is not dominant (see e.g. \citealt{Madau04,Miralda-Escude.Haehnelt.Rees_2000}) )."
 In addition. à model in which reionisation were dominated by mini-quasars. would. most. likely overpredict also the observed soft. X-ray background ?..," In addition, a model in which reionisation were dominated by mini-quasars would most likely overpredict also the observed soft X-ray background \cite{Salvaterra.Haardt.Ferrara_2005}."
 We. next compare the jonisingEN emiüssiv2.ilv used in our simulations to observational estimates based. on recent measurements of the galaxy luminosity function at d«z S., We next compare the ionising emissivity used in our simulations to observational estimates based on recent measurements of the galaxy luminosity function at $4<z<8$ .
 Fortiis purpose. we compute the ionising emissivity from ealaxies using the recent fit to the redshift evolution of the galaxy luminosity function. presented by 7..," For this purpose, we compute the ionising emissivity from galaxies using the recent fit to the redshift evolution of the galaxy luminosity function presented by \cite{Bouwens11}."
" We assume a spectra energy distribution c,x7° for 012«A3000. and c,x£ (Le a —3)for A«912. «ith an additional factor of six break at the Lyman limit (e.g. 773)."," We assume a spectral energy distribution $\epsilon_{\nu} \propto \nu^{0}$ for $912 {\rm \AA}<\lambda<3000 \rm
\AA$ and $\epsilon_{\nu} \propto \nu^{-3}$ (i.e. $\alpha=3$ ) for $\lambda< 912 \rm \AA$, with an additional factor of six break at the Lyman limit (e.g. \citealt{Leitherer99,Madau99}) )."
 In addition. we adoX two cdillerent redshift evolutions or the faint-enel slope: he ? best fit app=1.840.05(z6). and a steeper faint end. slope of apr=1.9Q.1(z6).," In addition, we adopt two different redshift evolutions for the faint-end slope: the \cite{Bouwens11} best fit $\alpha_{\rm LF} = -1.84 - 0.05(z-6)$, and a steeper faint end slope of $\alpha_{\rm LF} = -1.9 - 0.1(z-6)$."
 These choices are intented to represent he considerable observational uncertainty in the faint-end slope., These choices are intented to represent the considerable observational uncertainty in the faint-end slope.
 The resulting emissivities are cisplav« Las the hatched regions in Figure 9.. with the results from he two cillerent [aint end slope evolutions shown in each mel.," The resulting emissivities are displayed as the hatched regions in Figure \ref{fig:emissivity}, with the results from the two different faint end slope evolutions shown in each panel."
 The evan and orange hatching assume ionising photon escape fractions TEN0.2 anc fou= 05. while|o the lower and upper limits to the hatching correspond to the emissivity obtained by integrating the ? Iuminositv function [fit to a lower magnitude Limit ob Aly=15 and Αν 10. respectively.," The cyan and orange hatching assume ionising photon escape fractions of $f_{\rm esc}=0.2$ and $f_{\rm esc}=0.5$ , while the lower and upper limits to the hatching correspond to the emissivity obtained by integrating the \cite{Bouwens11} luminosity function fit to a lower magnitude limit of $M_{\rm UV}=-18$ and $M_{\rm UV}=-10$ , respectively."
 Ehese limits roughly οorrespond to the magnitude limit of the observational data and the expected magnitude of a galaxy in a halo with virial. temperature 2.1011 (7)). pos»ectivelv.," These limits roughly correspond to the magnitude limit of the observational data and the expected magnitude of a galaxy in a halo with virial temperature $2\times 10^{4}\rm\,K$ \citealt{Trenti10}) ), respectively."
 These are compared. to the emissivities used in moclels €1.2-a3 (solid curvo) and €1.6-03 (dashed. curve).," These are compared to the emissivities used in models ${\mathcal
 E}$$\alpha$ 3 (solid curve) and ${\mathcal E}$ $\alpha$ 3 (dashed curve)."
 Observational οmnstraints on the enissivitv at. z6 (red circles with error bars) derived," Observational constraints on the emissivity at $z\leq
6$ (red circles with error bars) derived"
teiirporal features of the radiation respectively.,temporal features of the radiation respectively.
 Iu the final section we παΊο our fincines and eive a brief discussion on implications of the bubble., In the final section we summarize our findings and give a brief discussion on implications of the bubble.
 We asstune that a burst itself arises from a series of explosive reconnection events., We assume that a burst itself arises from a series of explosive reconnection events.
 After the CRB. we are Ieft with a highly maenetized. rapidly rotating compact object.," After the GRB, we are left with a highly magnetized, rapidly rotating compact object."
 Let's first assume that it is a iiillisecoucd maguctar with period P. surface maguctic field streneth B.. moment of inertia Z. radius At. aud angele between the rotation axis and maenetic dipole moment 0.," Let's first assume that it is a millisecond magnetar with period $P$, surface magnetic field strength $B_s$, moment of inertia $I$, radius $R_{\rm M}$, and angle between the rotation axis and magnetic dipole moment $\theta$."
" Since such a pulsar loses its rotational cucrey through the maeuetic dipole torque. the huuinosity of a resulting relativistic wind is given by where By44,=D;sinf10:6G. Rape=Rayολο, aud ο "," Since such a pulsar loses its rotational energy through the magnetic dipole torque, the luminosity of a resulting relativistic wind is given by where $B_{\perp,14}=B_s\sin\theta/10^{14}{\rm G}$, $R_{\rm
M,6}=R_{\rm M}/10^6{\rm cm}$, and $P_{\rm ms}=P/1\,{\rm ms}$."
"Because of spiu-down. this huninosity will evolve with time as where fis the observer time in units of dav. συ=Μα...rep2_R©prTuas avs is the “initial” spin-down nuescale of the magnetar at the onset of the afterelow. Py=Poms.sLins is the rotation period at this time. and {ο={Ποοοι”, "," Because of spin-down, this luminosity will evolve with time as where $t$ is the observer time in units of day, $T_{{\rm
M},0}=0.58B_{\perp,14}^{-2}I_{45}R_{{\rm M},6}^{-6}P_{0,{\rm
ms}}^2$ days is the “initial"" spin-down timescale of the magnetar at the onset of the afterglow, $P_0=P_{0,{\rm ms}}\times 1\,{\rm
ms}$ is the rotation period at this time, and $I_{45}=I/10^{45}{\rm g}\,{\rm cm}^2$."
Usov (1992) also asmuued that he early spin-down could be due to eravitational wave radiation besides magnetic dipole radiation., Usov (1992) also assumed that the early spin-down could be due to gravitational wave radiation besides magnetic dipole radiation.
 However. the uunmnositv for gravitational wave radiation depends ou he stellar ellipticity which is poorly known. aud so we welected the effect of this mechanisin ou spin-down.," However, the luminosity for gravitational wave radiation depends on the stellar ellipticity which is poorly known, and so we neglected the effect of this mechanism on spin-down."
 We now diseuss another case in which the central object isa rapidly rotating black hole surrounded by an accretion disk., We now discuss another case in which the central object is a rapidly rotating black hole surrounded by an accretion disk.
 If the amplified maeuetic feld iu the disk docs uot evolve significantly with time during fallback of the ejecta. the rotational enerev of this black hole will be eradually extracted by the Blandford-Zuajels mechauisui. whose huninosity is approximated by where f(4)—1.[a|vbe?)2|2oa*f8 for the rotation parameter e<1l. My is the black-hole mass. and Bypass is the disk ficld strength in units of 1017 C (Lee. Wijers Brown 2000).," If the amplified magnetic field in the disk does not evolve significantly with time during fallback of the ejecta, the rotational energy of this black hole will be gradually extracted by the Blandford-Znajek mechanism, whose luminosity is approximated by where $f(a)=1-[(1+\sqrt{1-a^2})/2]^{1/2}\simeq a^2/8$ for the rotation parameter $a\ll 1$, $M_{\rm BH}$ is the black-hole mass, and $B_{{\rm BH},15}$ is the disk field strength in units of $10^{15}$ G (Lee, Wijers Brown 2000)."
 We note that a typical value (~107eres ly of the wind hunilosity dn equations (1) zud (3) has been invoked by Rees Alésszárros (2000) το explain the observed ion lues from some CRBs within the framework of the collapsar/hiyperuova imodel.," We note that a typical value $\sim 10^{47}\,{\rm erg}\,{\rm s}^{-1}$ ) of the wind luminosity in equations (1) and (3) has been invoked by Rees Mésszárros (2000) to explain the observed iron lines from some GRBs within the framework of the collapsar/hypernova model."
 Cousicderine the rotational enerev of the black hole Epi=αλλ-(a?/S)Mpc? for a<1 (Lee et al.," Considering the rotational energy of the black hole $E_{\rm
BH}=f(a)M_{\rm BH}c^2\simeq (a^2/8)M_{\rm BH}c^2$ for $a\ll 1$ (Lee et al."
" 2000). aud κκιο Fou=L,,. we find that the Dlaudford-Zuajek mechanisin vields spin-down. similarly to the magnetar case. if the accreted angular momenta is ucelected because its rate secnus to be below the torque driven bv the Blaucdtord-Zuajek miechanisni at tine of davs after the burst for typical parameters in the collapsar/lvpernova inodel."," 2000), and assuming $\dot{E}_{\rm BH}=L_w$, we find that the Blandford-Znajek mechanism yields spin-down, similarly to the magnetar case, if the accreted angular momentum is neglected because its rate seems to be below the torque driven by the Blandford-Znajek mechanism at time of days after the burst for typical parameters in the collapsar/hypernova model."
 Thus. we obtain a crude evolution law of the above DIuninuosity. where Topi=LOBpaas(Atou3M.)loy0.3)?2 davs is the “initial” spin-down time of the black hole and og is the “initial” rotation parameter.," Thus, we obtain a crude evolution law of the above luminosity, where $T_{{\rm BH},0}=1.0B_{{\rm BH},15}^{-2}(M_{\rm
BH}/3M_\odot)^{-1}(a_0/0.2)^{-2}$ days is the “initial"" spin-down time of the black hole and $a_0$ is the “initial"" rotation parameter."
" One cau easily see that equations (2) aud (1) are simular. which iuplics that our discussion of a relativistic wind bubble in the renaming text for magnuetars should be valid for black holes,"," One can easily see that equations (2) and (4) are similar, which implies that our discussion of a relativistic wind bubble in the remaining text for magnetars should be valid for black holes."
" Similarly to the Crab Nebula. a rotating magnetar at the center of an afterglow eeuerates a highly relativistic wine domünated by the eucrey flux of ele pairs. with bulk Lorentz factor of 24.—1010""."," Similarly to the Crab Nebula, a rotating magnetar at the center of an afterglow generates a highly relativistic wind dominated by the energy flux of $^+$ $^-$ pairs, with bulk Lorentz factor of $\gamma_w\sim 10^4-10^7$ ."
" Atovan (1999) argue that the Crab pulsar initially had 5,—LO! to interpret the measured radio spectrum of the Crab Nebula.", Atoyan (1999) argued that the Crab pulsar initially had $\gamma_w\sim 10^4$ to interpret the measured radio spectrum of the Crab Nebula.
 We adop ~=10! asa fiducial value in our calculations., We adopt $\gamma_w=10^4$ as a fiducial value in our calculations.
" Because 5, isch larger than the Loreutz factor of the 1iedimun swep up by the fireball. this wind passes through a shock frou and decelerates to match the expansion velocity of the swept-up medium."," Because $\gamma_w$ is much larger than the Lorentz factor of the medium swept up by the fireball, this wind passes through a shock front and decelerates to match the expansion velocity of the swept-up medium."
 Therefore. a relativistic wind bubble. as a result of mteraction of the wind with the medium. should include two shocks: a reverse shock that propagates iuto the cold wind and a forward shock that propagates into the ambient medium.," Therefore, a relativistic wind bubble, as a result of interaction of the wind with the medium, should include two shocks: a reverse shock that propagates into the cold wind and a forward shock that propagates into the ambient medium."
 Thus. there are four regions separated in the bubble by these shocks: (1) the uushocked medi. (2) the forward-shocked medium. (3) the reverse-shocked wind gas. and (1) the uushocked cold wind. where reeious 2 and 3 are separated by a contact discoutimty.," Thus, there are four regions separated in the bubble by these shocks: (1) the unshocked medium, (2) the forward-shocked medium, (3) the reverse-shocked wind gas, and (4) the unshocked cold wind, where regions 2 and 3 are separated by a contact discontinuity."
 For simplicity. we here assume that two initially-forming orward shocks during interactions of the fireball both with he mecditun aud with the wind have eventually iierged to oue forward shock. aud also neglect effects of the barvon oading Whose mass is much less than the sept-up mass. when the observer tine far exceeds the initial deceleration inescale.," For simplicity, we here assume that two initially-forming forward shocks during interactions of the fireball both with the medium and with the wind have eventually merged to one forward shock, and also neglect effects of the baryon loading whose mass is much less than the swept-up mass, when the observer time far exceeds the initial deceleration timescale."
" We denote ο aud P as the barvou umber deusity aud oessure of region ""/ m its own rest frame respectively. and >; is the Lorentz factor of region ""7 measured in the ocal mediunmn's rest frame."," We denote $n_i$ and $P'_i$ as the baryon number density and pressure of region $i$ "" in its own rest frame respectively, and $\gamma_i$ is the Lorentz factor of region $i$ "" measured in the local medium's rest frame."
" We derive the relative Loreutz actor of region3 measured ium the rest frame of region | as 53422(1/26/510/0)=twf(233)29d doy 5,o93> T. nmapleus a relativistic reverse shock."," We derive the relative Lorentz factor of region3 measured in the rest frame of region 4 as $\gamma_{34}\simeq
(1/2)(\gamma_w/\gamma_3+\gamma_3/\gamma_w)\simeq
\gamma_w/(2\gamma_3)\gg 1$ for $\gamma_w\gg \gamma_3\gg 1$ , implying a relativistic reverse shock."
" Because the electron number density iu the rest frame comoving with the uushocked wind is a,=Ly.Cor?52a,e) Gvhere r is the radius of region 3 and a, is the electrou. mass). according to the jump conditions for a relativistic shock (Blandford Melee 1976). the pressure of region 3 is calculated by Neglecting the prescuce of the reverse shock and the radiative eucrev loss of region 2. aud ΠΕ anambicut iterstella iuediuu with coustaut density of ny. the properties of the shocked 1iediuu iu region 2 should satisfy"," Because the electron number density in the rest frame comoving with the unshocked wind is $n_4=L_w/(4\pi r^2\gamma_w^2m_ec^3)$ (where $r$ is the radius of region 3 and $m_e$ is the electron mass), according to the jump conditions for a relativistic shock (Blandford McKee 1976), the pressure of region 3 is calculated by Neglecting the presence of the reverse shock and the radiative energy loss of region 2, and assuming anambient interstellar medium with constant density of $n_1$ , the properties of the shocked medium in region 2 should satisfy"
5truciu BStrueiu Cnr Cnr Cnr Cnr σα1) clussilü ον CLUSSS ciusshbxlO 1102 ciutir clurd cutis clures ,"5truein 8truein cmr8 cmr8 cmr8 cmr8 cmr10 cmssi10 cmss10 cmss8 cmssbx10 2 cmti7 cmr6 cmti8 cmr8 \def\ref{\par\noindent\hangindent 15pt}
 "
224085)) is an extremely active RS CVn star. exhibiting both strong chromospheric and coronal activity. in addition to photometric variability caused by spots.,") is an extremely active RS CVn star, exhibiting both strong chromospheric and coronal activity, in addition to photometric variability caused by spots."
 It has had one of the highest observed photometric amplitudes (2) and is also claimed to be the brightest X-ray source within 50 pe (?).., It has had one of the highest observed photometric amplitudes \citep{doyle1989} and is also claimed to be the brightest X-ray source within 50 pc \citep{makarov2003}.
 More recent studies of II. Peg include estimates of its differential rotation using ground-based (?) and satellite observations (2).., More recent studies of II Peg include estimates of its differential rotation using ground-based \citep{roettenbacher2011} and satellite observations \citep{siwak2010}.
 Both studies indicate a weak solar-type differential rotation., Both studies indicate a weak solar-type differential rotation.
 The latter study based on data from the MOST satellite confirmed previous findings that flares must be related to active regions. since they are more frequently observed when the most spotted hemisphere ts visible (e.g.229) .," The latter study based on data from the MOST satellite confirmed previous findings that flares must be related to active regions, since they are more frequently observed when the most spotted hemisphere is visible \citep[e.g.][]{mohin1993,teriaca1999,frasca2008}."
 has been spectroscopically monitored for nearly 20 years with the SOFIN spectrograph at the Nordic Optical Telescope (La Palma. Spain).," has been spectroscopically monitored for nearly 20 years with the SOFIN spectrograph at the Nordic Optical Telescope (La Palma, Spain)."
 Its observations have enabled both Doppler imaging and long-term studies of the spot activity to be performed (222)..," Its observations have enabled both Doppler imaging and long-term studies of the spot activity to be performed \citep{berdyugina1998,berdyugina1999,
 lindborg2011}."
 A summary of these results can be found in our earlier paper ?.. which was based on observations from 1994 to 2002.," A summary of these results can be found in our earlier paper \citet{lindborg2011}, which was based on observations from 1994 to 2002."
 In these Doppler images. there were usually two active longitudes. one of which was persistent and stronger almost throughout the nine years of spectroscopic monitoring.," In these Doppler images, there were usually two active longitudes, one of which was persistent and stronger almost throughout the nine years of spectroscopic monitoring."
 Major changes. occurring on a timescale of shorter than one year. could also be seen.," Major changes, occurring on a timescale of shorter than one year, could also be seen."
 For instance. the spot activity could switch to another. weaker. longitude. for a short period of time.," For instance, the spot activity could switch to another, weaker, longitude, for a short period of time."
 These shifts are similar to the ‘flip-flops’ described by Berdyugina and collaborators (???).. except that they were not found to follow any periodicity.," These shifts are similar to the 'flip-flops' described by Berdyugina and collaborators \citep{berdyugina1998,BT98,berdyugina1999}, except that they were not found to follow any periodicity."
 In addition to the short-term shifts of the active longitudes. we also discovered a drift in the active regions with respect to the orbital rotation frame. indicating that the spot-generating structure rotated slightly faster than the tidally locked binary system.," In addition to the short-term shifts of the active longitudes, we also discovered a drift in the active regions with respect to the orbital rotation frame, indicating that the spot-generating structure rotated slightly faster than the tidally locked binary system."
 Rapidly rotating late-type stars with deep convective envelopes are expected to exhibit very little differential rotation. because it is theoretically predicted to be suppressed in the rapid rotation regime (e.g.2): for Peg.. this has also been observationally confirmed (??)..," Rapidly rotating late-type stars with deep convective envelopes are expected to exhibit very little differential rotation, because it is theoretically predicted to be suppressed in the rapid rotation regime \citep[e.g.][]{KR99}; for , this has also been observationally confirmed \citep{siwak2010,roettenbacher2011}. ."
 However. several stars have been found to contradict this prediction (e.g.222).," However, several stars have been found to contradict this prediction \citep[e.g.][]{hackman2001,jeffers2008,frasca2011}."
 Furthermore it is unclear whether surface differential rotation can be recovered by following the motion of large spots (?).., Furthermore it is unclear whether surface differential rotation can be recovered by following the motion of large spots \citep{korhonen2011}.
 Nevertheless. in the case of suppressed differential rotation. the dynamos working in rapidly rotating stars are expected to be of the a7-type. the magnetic field generation only being due to the inductive action of convective turbulence.," Nevertheless, in the case of suppressed differential rotation, the dynamos working in rapidly rotating stars are expected to be of the $\alpha^2$ -type, the magnetic field generation only being due to the inductive action of convective turbulence."
 According to both linear (e.g.?) and nonlinear solutions (e.g.?) of the a7- dynamo equations. the nonaxisymmetric modes become more easily excited in the rapid rotation regime.," According to both linear \citep[e.g.][]{krause1980meanfield} and nonlinear solutions \citep[e.g.][]{MBBT95} of the $\alpha^2$ -dynamo equations, the nonaxisymmetric modes become more easily excited in the rapid rotation regime."
 The #7=1 mode. representing an azimuthally varying field changing sign once over the full longitude span. is commonly the preferred field configuration.," The $m\!=\!\!1$ mode, representing an azimuthally varying field changing sign once over the full longitude span, is commonly the preferred field configuration."
 This is unsurprising. as the largest scale mode is always the least affected by diffusive effects.," This is unsurprising, as the largest scale mode is always the least affected by diffusive effects."
 The nonaxisymmetric modes. turn out to be waves migrating m the azimuthal direction. not necessarily having the rotation period of the star (e.g.?)..," The nonaxisymmetric modes turn out to be waves migrating in the azimuthal direction, not necessarily having the rotation period of the star \citep[e.g.][]{krause1980meanfield}."
 Both slower and faster dynamo waves can occur. depending for example on the profile and properties of the turbulent transport coeficients.," Both slower and faster dynamo waves can occur, depending for example on the profile and properties of the turbulent transport coefficients."
 The faster waves with dipole symmetry (S1) were found to be preferred in linear models with simple profiles (?).. and slower waves with quadrupolar symmetry (AI) in more complicated nonlinear models when solving for the dynamics (?)..," The faster waves with dipole symmetry (S1) were found to be preferred in linear models with simple profiles \citep{krause1980meanfield}, and slower waves with quadrupolar symmetry (A1) in more complicated nonlinear models when solving for the dynamics \citep{tuominen2002starspot}."
 From the viewpoint of dynamo theory. two migratory active longitudes are thus an expected result.," From the viewpoint of dynamo theory, two migratory active longitudes are thus an expected result."
 We. therefore. interpret the behaviour seen in during 1994-2002 as a manifestation of such a dynamo wave.," We, therefore, interpret the behaviour seen in during 1994–2002 as a manifestation of such a dynamo wave."
 The objective of the present paper was to inspect whether the azimuthal dynamo wave persisted on the star in the most recent observational data of the object., The objective of the present paper was to inspect whether the azimuthal dynamo wave persisted on the star in the most recent observational data of the object.
 Our observations were made using theSOFIN high-resolution écchelle spectrograph at the 2.56m Nordic Optical Telescope, Our observations were made using theSOFIN high-resolution écchelle spectrograph at the 2.56m Nordic Optical Telescope
luminosity correlation for steepespectrunr quasars to be mmagnitudes (0.6οκ).,luminosity correlation for steep-spectrum quasars to be magnitudes dex).
 Vhis will be slightly smaller than e. since there should exist radio galaxies whose intrinsic optical luminosities are on average fainter than those of the quasars.," This will be slightly smaller than $\sigma$, since there should exist radio galaxies whose intrinsic optical luminosities are on average fainter than those of the quasars."
 However. Serjeant et al," However, Serjeant et al."
s sample was dominated by distant (z2 1) quasars whose high bluminosities would imply a large opening angle for the torus.,'s sample was dominated by distant $z>1$ ) quasars whose high luminosities would imply a large opening angle for the torus.
 At large opening angles. most objects are observed as quasars. and. therefore the dispersion in optical magnitude in a quasar saniple is close to the intrinsic dispersion of the parent AGN population (Figure 2)).," At large opening angles, most objects are observed as quasars, and therefore the dispersion in optical magnitude in a quasar sample is close to the intrinsic dispersion of the parent AGN population (Figure \ref{fig:disp}) )."
 We thus adopt σ=0.6., We thus adopt $\sigma = 0.6$.
 The second constrain is the observed quasar (or radio galaxy) fraction: we mus constrain 6) such that the fraction in our simulation matches that in a racio-selected sample., The second constraint is the observed quasar (or radio galaxy) fraction; we must constrain $\theta_0$ such that the fraction in our simulation matches that in a radio-selected sample.
 Following the reasoning of Laing et ((1904) anc others. when comparing our results to observed. quantities. we consider only those radio galaxies with prominen emission lines.," Following the reasoning of Laing et (1994) and others, when comparing our results to observed quantities, we consider only those radio galaxies with prominent emission lines."
" At z«0.43 and OF«à13"". where there exist. high-quality optical spectra allowing classification on the basis of broad Lla (Laing et 11994). there are 15 such radio galaxies and 9 quasars."," At $z < 0.43$ and $0\hour < \alpha <
13\hour$, where there exist high-quality optical spectra allowing classification on the basis of broad $\alpha$ (Laing et 1994), there are 15 such radio galaxies and 9 quasars."
" The observed radio galaxy fraction is therefore 0.625. corresponding to 6,=50° for 6=0.6 (Figure 3: the value is extremely insensitive to the choice of 7)."," The observed radio galaxy fraction is therefore 0.625, corresponding to $\theta_0 = 50\degree$ for $\sigma = 0.6$ (Figure \ref{fig:frg}; the value is extremely insensitive to the choice of $\sigma$ )."
 Figure 4 illustrates the amount by which quasars are expected to be overluminous compared: to radio galaxies as a function of @) and o. ancl for the values determined previously. we find £Loso?/iLnco?=L8.," Figure \ref{fig:rellum} illustrates the amount by which quasars are expected to be overluminous compared to radio galaxies as a function of $\theta_0$ and $\sigma$, and for the values determined previously, we find $\langle L_{\rm QSO} \rangle / \langle L_{\rm RG} \rangle = 1.8$."
 This is elfectively the same as the dilference in [luminosities for the two groups of objects. which could therefore he explained if the cenussion line luminosity is directly proportional to the ionizing luminosity.," This is effectively the same as the difference in luminosities for the two groups of objects, which could therefore be explained if the emission line luminosity is directly proportional to the ionizing luminosity."
 ts constaney of equivalent width over approximately three orders. of magnitude in luminosity (Miller ct 11992) indicates that it is., Its constancy of equivalent width over approximately three orders of magnitude in luminosity (Miller et 1992) indicates that it is.
 Why then is a similar ellect not seen in the luminosity. since both emission lines correlate well with optical luminosity over several orders of magnitude in quasars?," Why then is a similar effect not seen in the luminosity, since both emission lines correlate well with optical luminosity over several orders of magnitude in quasars?"
 We believe that the answer is once again due to a jas introduced by scatter., We believe that the answer is once again due to a bias introduced by scatter.
 Although the effective ionization xvwameter of the NLR. does not correlate with luminosity (the range of quasar luminosities is [ar larger than the observed range o£ C. 3xlogSeg. SSaunders et 11989). the most luminous objects at a given redshift will also have the highest ionization parameters.," Although the effective ionization parameter of the NLR does not correlate with luminosity (the range of quasar luminosities is far larger than the observed range of $U$, $-3 \leq \log U
\leq -1$; Saunders et 1989), the most luminous objects at a given redshift will also have the highest ionization parameters."
 Figure 5. shows the dependence of aand Iluminosity as a function of ionization parameter for an optically-thick plane-parallel slab., Figure \ref{fig:linelum} shows the dependence of and luminosity as a function of ionization parameter for an optically-thick plane-parallel slab.
 Phis was calculated with, This was calculated with
Fig. Ifa,Fig. \ref{fig:ngc193all}( (
) MD the total-intensity clistribution over 1193 at 4.9 mi1 and 4.05-aresee FAVLAL resolution.,a) shows the total-intensity distribution over 193 at 4.9 GHz and 4.05-arcsec FWHM resolution.
 The svninie thei:jets. appear broaden rapidly ancl also »end away fron initial straight path as they reach the micpoints of symmetric lobes., The symmetrical jets appear to broaden rapidly and also bend away from their initial straight path as they reach the midpoints of symmetric lobes.
 Phe lobes both have. well-defined leading edges that approximate ares of circles in oojection on the sky (the W lobe having a Iarger radius of curvature). but they lack hot spots that might mark the ermination points of the jets.," The lobes both have well-defined leading edges that approximate arcs of circles in projection on the sky (the W lobe having a larger radius of curvature), but they lack hot spots that might mark the termination points of the jets."
 A broad. faint emission bridge ills the central region of the source between the lobes and appears to be wider than the lobes in the N-S direction in he centre of the source (similar to the “wines” observed in some FRAIL sources: )).," A broad, faint emission bridge fills the central region of the source between the lobes and appears to be wider than the lobes in the N-S direction in the centre of the source (similar to the “wings” observed in some II sources; )."
 lies l((b) and (c) taken together show that. broad. “caps” of emission can be delineated in both lobes by enhanced intensity graciients ancl by lower-than-average values of af., Figs \ref{fig:ngc193all}( (b) and (c) taken together show that broad “caps” of emission can be delineated in both lobes by enhanced intensity gradients and by lower-than-average values of $\alpha^{4.9}_{1.4}$.
 ‘The intensity gracients are largest at the outer edges of these caps. but there are also gradient. features within the lobes (marked with arrows on Fig.," The intensity gradients are largest at the outer edges of these caps, but there are also gradient features within the lobes (marked with arrows on Fig."
 Ibb) which coincide with the edges of the IEatter-specteum region., \ref{fig:ngc193all}b b) which coincide with the edges of the flatter-spectrum region.
 The regions where the jets are most prominent have low spectral indices around 0.6., The regions where the jets are most prominent have low spectral indices around 0.6.
 The spectral index at the trailing edges of the caps steepens smoothly to àz0.9 where the emission merges with the ooader. symmetric lobes.," The spectral index at the trailing edges of the caps steepens smoothly to $\alpha \approx
0.9$ where the emission merges with the broader, symmetric lobes."
 The most. dilfuse lobe emission las spectral indices increasing from σε] at the edges of the most elongated parts of the lobes to z 1.4 near the centre of he source. à spectral-index. pattern characteristic of lobed radio sources of both FR classes (see Section 4.2)).," The most diffuse lobe emission has spectral indices increasing from $\approx$ 1 at the edges of the most elongated parts of the lobes to $\approx$ 1.4 near the centre of the source, a spectral-index pattern characteristic of lobed radio sources of both FR classes (see Section \ref{lobes}) )."
 There are regions of emission with spectral index approaching 2 at he edges of the faintest emission on the N and S edges of he Fig. {4} , There are regions of emission with spectral index approaching 2 at the edges of the faintest emission on the N and S edges of the Fig. \ref{fig:ngc193all}( (
shows that the distribution of the apparent magnetic Lield direction over the lobes is basically. circumferential. while the magnetic field in the jets is perpendicular to the jets over most. of their lengths.,"d) shows that the distribution of the apparent magnetic field direction over the lobes is basically circumferential, while the magnetic field in the jets is perpendicular to the jets over most of their lengths."
 The structure of the apparent. magnetic field in both themjets and the lobes appears regular. anc characteristic of commonly found in jets of FRE sources and in the lobes of both FR. classes.," The structure of the apparent magnetic field in both the jets and the lobes appears regular, and characteristic of that commonly found in jets of I sources and in the lobes of both FR classes."
 Fig. 2((, Fig. \ref{fig:ngc193hires}( (
a) shows the total-intensity distribution over the jets at GOllz with a resolution of L.6aaresec EFWHIA,a) shows the total-intensity distribution over the jets at GHz with a resolution of arcsec FWHM.
L Both jets are fairly straight. ancl similar in overall appearance. exhibiting rapid lateral expansion just bevond the distance from the unresolved: nuclear radio source at which the I2 jet is markedly brighter than the W jet.," Both jets are fairly straight and similar in overall appearance, exhibiting rapid lateral expansion just beyond the distance from the unresolved nuclear radio source at which the E jet is markedly brighter than the W jet."
 Their edges are well delineated by steep transverse intensity| eracients near the midpoint of the source (Fig., Their edges are well delineated by steep transverse intensity gradients near the midpoint of the source (Fig.
 2bb)., \ref{fig:ngc193hires}b b).
 The surface brightnesses of both jets decrease smoothly with distance from the nucleus except at a distance of =? aaresec. where there are more sudden drops. mostly clearly visible as “ares” crossing the overalljets on the eracicnt image (indicated on Fig.," The surface brightnesses of both jets decrease smoothly with distance from the nucleus except at a distance of $\approx$ arcsec, where there are more sudden drops, mostly clearly visible as “arcs” crossing the jets on the gradient image (indicated on Fig."
 2bb)., \ref{fig:ngc193hires}b b).
 The spectral index of the jet-dominated emission appears to steepen with distance from the nucleus. but this is almost certainly the result of superposition of dimming jets on steeper-specteum Lobe emission.," The overall spectral index of the jet-dominated emission appears to steepen with distance from the nucleus, but this is almost certainly the result of superposition of dimming jets on steeper-spectrum lobe emission."
 The prominent arc in the E jet is associated with a slight llattening in the spectrum (Fig., The prominent arc in the E jet is associated with a slight flattening in the spectrum (Fig.
 2cc)., \ref{fig:ngc193hires}c c).
 Fig. 2€(, Fig. \ref{fig:ngc193hires}( (
d) shows the intensity and. apparent. magnetic Leld. distributions over the inner zdb5-arcsec regions of both jets at 1.35-aresec resolution.,d) shows the intensity and apparent magnetic field distributions over the inner $\approx$ 45-arcsec regions of both jets at 1.35-arcsec resolution.
 The magnetic Ποιά organization over both jets is quite regular. with the field closest. to the jet. axis. predominantly orientated perpendicular to the axis.," The magnetic field organization over both jets is quite regular, with the field closest to the jet axis predominantly orientated perpendicular to the axis."
 Phe apparent field. at the edges of the inner jets is parallel to the rapidlv-expanding outer isophotes., The apparent field at the edges of the inner jets is parallel to the rapidly-expanding outer isophotes.
 Farther from the nucleus. the edge field directions in both jets converge towards the axis to form almost circular patterns.," Farther from the nucleus, the edge field directions in both jets converge towards the axis to form almost circular patterns."
 Fig. 2((, Fig. \ref{fig:ngc193hires}( (
0) shows the 4.9-Cllz total intensity and apparent distributions over the inner 15 aresecn NM5 kpe) of the jets at Q.45-aresec resolution.,e) shows the 4.9-GHz total intensity and apparent magnetic-field distributions over the inner 15 arcsec $\approx$ 4.5 kpc) of the jets at 0.45-arcsec resolution.
 Both Jese faint inner regions in the first 2 aresce [rom the nucleus before they brighten and subsequently flare., Both jets exhibit faint inner regions in the first $\approx$ 2 arcsec from the nucleus before they brighten and subsequently flare.
 There is xight. non-axisvmumietric knot structure in the first zd aresec of the I2 jet. downstream of the llaring point1999)... where both jets brighten abruptly.," There is bright, non-axisymmetric knot structure in the first $\approx$ 4 arcsec of the E jet, downstream of the flaring point, where both jets brighten abruptly."
 Fig., Fig.
 3. shows the total-intensitv. brightness and spectral-index. clistributions over the whole ofME 0206|35 at two resolutions.," \ref{fig:0206all} shows the total-intensity, brightness gradient, and spectral-index distributions over the whole of 0206+35 at two resolutions."
 Fig. Bl, Fig. \ref{fig:0206all}( (
a) at 1.4 Gilg. 4.5-aresee ENLIM resolution. shows that the larec scale structure consists of two lobes. overlapping and circular in cross-section. with well-defined: outer edges to the NW. and SE of the source. -- on fainter cdilfuse emission to the N and S. ΕΙ ((,"a) at 1.4 GHz, 4.5-arcsec FWHM resolution, shows that the large scale structure consists of two lobes, overlapping and circular in cross-section, with well-defined outer edges to the NW and SE of the source, superimposed on fainter diffuse emission to the N and S. Fig. \ref{fig:0206all}( ("
b). at CGiGLIz and 1ucesec resolution. shows the source in more detail but with reduced. sensitivity to the argest-scale emission.,"b), at GHz and 1.2-arcsec resolution, shows the source in more detail but with reduced sensitivity to the largest-scale emission."
 At this resolution. both lobes show sharp outer boundaries.," At this resolution, both lobes show sharp outer boundaries."
 The roughly circular. edge of the AW lobe protruces bevond the dilfuse. emission. whereas he corresponding feature in the SEE is well inside the outer x»undary. of the source and is most obvious in the IE of the obe. close to the termination of the jet.," The roughly circular edge of the NW lobe protrudes beyond the diffuse emission, whereas the corresponding feature in the SE is well inside the outer boundary of the source and is most obvious in the E of the lobe, close to the termination of the jet."
 Lf the orientation of 6zz40 determined for the inner jets (Laing Briclle. in preparation) also applies to the lobes. then they are srestumably ellipsoidal with an axial ratio 1.6.," If the orientation of $\theta \approx 40^\circ$ determined for the inner jets (Laing Bridle, in preparation) also applies to the lobes, then they are presumably ellipsoidal with an axial ratio $\approx$ 1.6."
 Fig. 3((, Fig. \ref{fig:0206all}( (
b) also shows some internal structure in both andjets.,b) also shows some internal structure in both jets.
 The NW t has the brighter base. and both bends brightens as it enters its lobe. alter which its path meancers.," The NW jet has the brighter base, and both bends and brightens as it enters its lobe, after which its path meanders."
 The SE counter-jel appears to expand. more rapidly initially. then also meanders as it enters its lobe.," The SE counter-jet appears to expand more rapidly initially, then also meanders as it enters its lobe."
 Figs 3 and (d) show the L4-Gllz intensity gradient images mofM0206|35 at 4.5-aresce ancl 1.2-aresec. resolution.," Figs \ref{fig:0206all}( (c) and (d) show the 1.4-GHz intensity gradient images of 0206+35 at 4.5-arcsec and 1.2-arcsec resolution, respectively."
 The edges of thejets are clearly marked. by enhanced intensity gradients. at both resolutions. while significant internal structure is also apparent in the lobes.," The edges of the jets are clearly marked by enhanced intensity gradients at both resolutions, while significant internal structure is also apparent in the lobes."
 oth lobes exhibit strong brightness gradients at their outer edges in these displavs. corresponding to the sharp boundaries noted. earlier.," Both lobes exhibit strong brightness gradients at their outer edges in these displays, corresponding to the sharp boundaries noted earlier."
 Phere is a particularly striking correlation. between the main features of these intensitv-eracdient images and of the two 14 to 4.9-Cillz spectral-index images shown as Figs. 2 3((, There is a particularly striking correlation between the main features of these intensity-gradient images and of the two 1.4 to 4.9-GHz spectral-index images shown as Figs. \ref{fig:0206all}( (
0) and (£D).,e) and (f).
 Phe emission inside the brightest intensity gracients around the jet has a much lower spectral index. typically «0.65. than the 20.7 to 1.0 spectral index that is prevalent over the rest of the two spherical," The emission inside the brightest intensity gradients around the jet has a much lower spectral index, typically $<$ 0.65, than the $\approx$ 0.7 to 1.0 spectral index that is prevalent over the rest of the two spherical"
"We now use the standard: result. of spherical collapse. that a forming halo has a linear 0,4=94. where the critical density of collapse ὃς is independent of mass (and equals 1.69 in the Einstein de-Sitter limit. valid over a wide range of shifts!)).","We now use the standard result of spherical collapse, that a forming halo has a linear $\delta\rtot = \dc$, where the critical density of collapse $\dc$ is independent of mass (and equals 1.69 in the Einstein de-Sitter limit, valid over a wide range of )."
" We also assume that we are considering sullicientlv ↓⋜⊔⋅⋏∙≟∢⊾⊳∖≼∙⋜↧⇂⋖⋅≱∖≱∖∪↿⇂⋯↿↙↘∣↕⊓↕≼∼⋜⋯∣⋡∢⋅↥↓⋅⋖⊾⋜⋯⊾∠⇂⋜↧⊳∖⋜↧↓≻⋖⊾↓⋅⋯↓⋅∣⋡⋜∐↓∪⊔∪⇂ 4 . Oe. (or 05,4). and that 7j and ry are also small quantities."," We also assume that we are considering sufficiently large scales so that $\delta\rtot^l$ can be treated as a perturbation of $\delta\rtot$ (or $\delta\rtot^s$ ), and that $r\rl$ and $r\rs$ are also small quantities."
 Then the mean barvon fraction in halos is frbs-crb(1rj)... and the lowest order perturbation is derived to be =.," Then the mean baryon fraction in halos is =, and the lowest order perturbation is derived to be =."
. We now use the actual value of r(&) (see 3). specifically the fact that it approaches a constant on scales below the BAOs. with a value (depending on redshift but not A) that we denote rpx« (for. Large. Scale. Structure) following Naoz&Barkana(2007).," We now use the actual value of $r(k)$ (see 3), specifically the fact that it approaches a constant on scales below the BAOs, with a value (depending on redshift but not $k$ ) that we denote $\rLSS$ (for Large Scale Structure) following \citet{NB07}."
. ας. in the just-derived equations we can (reat nr=ress as a Constant (at a given redshift). since most of the density 9. needed to form a halo comes from scales well below the BAO scale.," Thus, in the just-derived equations we can treat $r\rs = \rLSS$ as a constant (at a given redshift), since most of the density $\dc$ needed to form a halo comes from scales well below the BAO scale."
 Thus. the mean barvon fraction in halos is frb=crb(l.. while on large scales (1.0. small &) the Ποιατοι is orf=ut.," Thus, the mean baryon fraction in halos is =, while on large scales (i.e., small $k$ ) the fluctuation is =."
 The remaining issue is the ellect of non-linear collapse. and the relation between the barvon fraction in the linearly-extrapolated halo perturbation and the barvon fraction in the actual virialized halo.," The remaining issue is the effect of non-linear collapse, and the relation between the baryon fraction in the linearly-extrapolated halo perturbation and the baryon fraction in the actual virialized halo."
 We show simulation results in 3 that only test the mean barvon fraction in halos (i.e. equation 2.2)) but do so over a range of redshifts. and suggest that halo collapse enhances the effect. and results in an effective value of russ that is amplified by a factor olf several.," We show simulation results in 3 that only test the mean baryon fraction in halos (i.e., equation \ref{barf}) ) but do so over a range of redshifts, and suggest that halo collapse enhances the effect and results in an effective value of $\rLSS$ that is amplified by a factor of several."
 One wav to understand this enhancement is to considerthe variation of ress with time., One way to understand this enhancement is to considerthe variation of $\rLSS$ with time.
" Lt declines (in absolute value) approximately as rxLéa (where a=l/(l]|2) is the scale factor). since (05,4.01) const. while OoX0 (until the cosmological constant becomes significant at low redshift)."," It declines (in absolute value) approximately as $r \propto 1/a$ (where $a=1/(1+z)$ is the scale factor), since $(\delta\rtot-\delta\rb) \approx $ const while $\delta\rtot \propto a$ (until the cosmological constant becomes significant at low redshift)."
 The decline of ress with time is of critical importance. since we are computing it according to linear theory. and it may not be appropriate to extrapolate ress all the wav to the halo formation time when we evaluate it in equation (2.2)).," The decline of $\rLSS$ with time is of critical importance, since we are computing it according to linear theory, and it may not be appropriate to extrapolate $\rLSS$ all the way to the halo formation time when we evaluate it in equation \ref{barf}) )."
 The barvon Uuetuations. which were erased on small scales. before. cosmic recombination. later continuously catch up with the dark matter (and thus with the total matter as well) in linear perturbation theory.," The baryon fluctuations, which were erased on small scales before cosmic recombination, later continuously catch up with the dark matter (and thus with the total matter as well) in linear perturbation theory."
 llowever. once a perturbation begins to form a halo and enters the non-linear stage of collapse. we expect that the rapid collapse will bring with it only the barvons already present within the perturbation. and the continued. decline of the linear-theory ris will become irrelevant for the halo eas content.," However, once a perturbation begins to form a halo and enters the non-linear stage of collapse, we expect that the rapid collapse will bring with it only the baryons already present within the perturbation, and the continued decline of the linear-theory $\rLSS$ will become irrelevant for the halo gas content."
 The upshot is that the simulations suggest that i£ ve use the linear-theory rpss (ancl. we assume. more generally for r(A)) then we must multiply it by an effective amplification factor ον: The resulting fluctuations in the luminosity density (equation 2)) are — Μα ο. μμ where is an ellective. bias factor. that measures the overall dependence of galaxy luminosity on. the underlving difference ;N between the barvon and total density Iluctuations.," The upshot is that the simulations suggest that if we use the linear-theory $\rLSS$ (and, we assume, more generally for $r(k)$ ) then we must multiply it by an effective amplification factor $A_r$: The resulting fluctuations in the luminosity density (equation \ref{dL}) ) are = + ) + [r(k) - ], where is an effective bias factor that measures the overall dependence of galaxy luminosity on the underlying difference $\Delta$ between the baryon and total density fluctuations."
 We have assumed thus far that we observe a fixed galaxy population. regardless of the varvine luminosity of its members.," We have assumed thus far that we observe a fixed galaxy population, regardless of the varying luminosity of its members."
 In reality. observed samples are limited by Dux. or equivalently by luminosity if for simplicity we consider ealaxies at à single redshift.," In reality, observed samples are limited by flux, or equivalently by luminosity if for simplicity we consider galaxies at a single redshift."
 Suppose the fraction of galaxies above luminosity £ is where ó is the luminosity function.," Suppose the fraction of galaxies above luminosity $L$ is F(L) = ), where $\phi$ is the luminosity function."
 Phen the observed number density of galaxies is =.. and the luminosity density of these galaxies is Pisas).. where —— o((L')dL.," Then the observed number density of galaxies is =, and the luminosity density of these galaxies is _L = L, where L = )."
. We assume for simplicity that the same luminosity distribution holds in clillerent regions. except that the luminosity. of all galaxies. is enhanced or diminished uniformly. in response to changes in the total density and the halo barvon fraction. as discussed in 2.1..," We assume for simplicity that the same luminosity distribution holds in different regions, except that the luminosity of all galaxies is enhanced or diminished uniformly in response to changes in the total density and the halo baryon fraction, as discussed in \ref{basic}."
 Hf à sample only includes galaxies above a detection threshold Livin. then we can analyze the variations of F'(L) by keeping © fixed and varving the cllective threshold Livin. while in pL we also include the perturbation in the luminosity of cach galaxy.," If a sample only includes galaxies above a detection threshold $\Lmin$, then we can analyze the variations of $F(L)$ by keeping $\phi$ fixed and varying the effective threshold $\Lmin$ , while in $\rho\rL$ we also include the perturbation in the luminosity of each galaxy."
 From equation (2.3)) we obtain a relative [uctuation, From equation \ref{FL}) ) we obtain a relative fluctuation _F =
LMC X-3 exhibit trausitious between the low/lard aud hieh states. and exhibit canonical spectra (7.andenceswithin)..,"LMC X-3 exhibit transitions between the low/hard and high states, and exhibit canonical spectra \citep[and references within]{mr03}."
 It is interesting to compare these systems with the ULXs. as many ULXs are thought to be ," It is interesting to compare these systems with the ULXs, as many ULXs are thought to be HMXBs."
IINENDs.. ? conducted a survey of all point sources 2107 ere in NADALNewton observations of 32 nearby galaxies., \citet{wmr06} conducted a survey of all point sources $>$ $^{38}$ erg $^{-1}$ in XMM-Newton observations of 32 nearby galaxies.
 They modeled the spectra of —100 X-ray sources. cach with £100 counts.," They modeled the spectra of $\sim$ 100 X-ray sources, each with $\ga$ 400 counts."
 Ten sources exhibited spectra consisteu with NGC300 N-I in Obs., Ten sources exhibited spectra consistent with NGC300 X-1 in Obs.
" 2 aud. {, iucludiug three ULXs."," 2 and 4, including three ULXs."
 A further 23 sources exhibited spectra consistent with Obs., A further 23 sources exhibited spectra consistent with Obs.
 1 of NGC300 X-1 and IC10 N-1. including 10 ULNs.," 1 of NGC300 X-1 and IC10 X-1, including 10 ULXs."
 1010 XN-1 is the first coufriined ΕΙWR system. aud NCGC300. N-1 is thought to be one also.," IC10 X-1 is the first confirmed BH+WR system, and NGC300 X-1 is thought to be one also."
 They are known to have strikinely simular X-ray periods and bpuuiuosities. so we studied them iu detail.," They are known to have strikingly similar X-ray periods and luminosities, so we studied them in detail."
 We have examined for the first tine the PDS of NCC300 X-1. in four NAINENewton observations. aud also of 10 N-1. in one NAINENewton observation.," We have examined for the first time the PDS of NGC300 X-1, in four XMM-Newton observations, and also of IC10 X-1, in one XMM-Newton observation."
 We find that they are well described by a shuple power law with spectral iudex z1. over the 0.0020.1 Iz range: such variability is characteristic of disc-acercting NBs iu their high state. or of some TAINB pulsus that are powered bv DondilIlovle accretion.," We find that they are well described by a simple power law with spectral index $\simeq$ 1, over the 0.002--0.1 Hz range; such variability is characteristic of disc-accreting XBs in their high state, or of some HMXB pulsars that are powered by Bondi-Hoyle accretion."
 We also find a striking resemblance between the Obs., We also find a striking resemblance between the Obs.
 1 NGOC200 N-1 spectrum and the IC10 N-1 spectu., 1 NGC300 X-1 spectrum and the IC10 X-1 spectrum.
 Since 1010 X-l is a confirmed WR|DII svstem. our results provide strong support for NGC300 X-1 being oue too.," Since IC10 X-1 is a confirmed WR+BH system, our results provide strong support for NGC300 X-1 being one too."
 We have considered four accretion scenarios: disc accretion or Doudi-EÉovle acerction. outo a BIT or NS.," We have considered four accretion scenarios: disc accretion or Bondi-Hoyle accretion, onto a BH or NS."
 We favour disc accretion outo a BIT for both svstems. but cannot rule out Boud-Tovle accretion onto a BU.," We favour disc accretion onto a BH for both systems, but cannot rule out Bond-Hoyle accretion onto a BH."
 The observed X-ray spectra of NCC300 N-1 allowed us to rule out Doudi-IToyle accretion outo a NS. but do not exclude cise acerction outo à NS|WR binary. or outo à NS LAINB that is merely coincident with the WR.," The observed X-ray spectra of NGC300 X-1 allowed us to rule out Bondi-Hoyle accretion onto a NS, but do not exclude disc accretion onto a NS+WR binary, or onto a NS LMXB that is merely coincident with the WR."
 Itf both NGC300 X-1 aud IC10 N-1 axe indeed WR|BI binaries. then they are spectrally distinct from the kuown DII XBs.," If both NGC300 X-1 and IC10 X-1 are indeed WR+BH binaries, then they are spectrally distinct from the known BH XBs."
 We xopose that this difference may be due to these svstenis being persistently in the high state. allowing them to keep their dise corona. like LAIC X-l: contrariwise. BIT LAINBs cuter the high state only in violent outbursts. where the corona may be ejected.," We propose that this difference may be due to these systems being persistently in the high state, allowing them to keep their disc corona, like LMC X-1; contrariwise, BH LMXBs enter the high state only in violent outbursts, where the corona may be ejected."
 NGC300 N-I and ICLO X-1 (aud possibly LMC X-1) would then comprise a new class of BIT binary. defiuec by their lieh mass transfer rates and consequent stable accretion disc configurations.," NGC300 X-1 and IC10 X-1 (and possibly LMC X-1) would then comprise a new class of BH binary, defined by their high mass transfer rates and consequent stable accretion disc configurations."
 Bernarci οἱ al. (, Bernardi et al. (
2008) argue that the most massive of these objects are likely to be prolate. viewed along the long axis (sec alse Thomas ct al.,"2008) argue that the most massive of these objects are likely to be prolate, viewed along the long axis (see alse Thomas et al."
 2007)., 2007).
 Li the shape is due to anisotropic velocities. then the velocity dispersion projected along the line of sight is S. 4. 4..," If the shape is due to anisotropic velocities, then the velocity dispersion projected along the line of sight is \ref{fig:MdynMvisbeta}, \ref{fig:MdynMvis} \ref{fig:MdynMvis}."
detected in Fig.,detected in Fig.
 2 are this time encircled in denoting clockwise/counterclockwise sense of rotation., \ref{verticalvel} are this time encircled in denoting clockwise/counterclockwise sense of rotation.
" For both time series we found a larger number of counterclockwise motions, though the difference is not very significant."," For both time series we found a larger number of counterclockwise motions, though the difference is not very significant."
 In general these events are evidenced by larger concentration of with the exception of a very few casesa in the second time series panel)) in which no sign of is detected though the areas are characterized by a junction of somewhat dark intergranular lanes., In general these events are evidenced by a larger concentration of with the exception of a very few cases in the second time series ) in which no sign of is detected though the areas are characterized by a junction of somewhat dark intergranular lanes.
 The distribution of vortices shows therefore a prominent correspondence with the location of intergranular lanes., The distribution of vortices shows therefore a prominent correspondence with the location of intergranular lanes.
 The circles in Fig., The circles in Fig.
" 3 are regions in which there is a strong concentration of tracers, forming in some cases ring-shaped features."," \ref{corchos} are regions in which there is a strong concentration of tracers, forming in some cases ring-shaped features."
" Although some of these regions display converging flows and correspond to downflows, they are not linked to very strong downflows when plotted over the image of vertical velocities (see in Fig. 2))."," Although some of these regions display converging flows and correspond to downflows, they are not linked to very strong downflows when plotted over the image of vertical velocities (see in Fig. \ref{verticalvel}) )."
 Their location also change from the first to the second time series (e.g. compare the distribution of and circles in the lower panel of Fig. 2))., Their location also change from the first to the second time series (e.g. compare the distribution of and circles in the lower panel of Fig. \ref{verticalvel}) ).
" They do not represent though the most significant part of the detected events but ~10% in both series, respectively."," They do not represent though the most significant part of the detected events but $\sim$ in both series, respectively."
" We have calculated the occurrence of the detected vortex-type events in the FOV obtaining 1.46 x 10? and 1.09 x 10? vortices Μπι for the SST/G-band series and s2,, respectively (the values increase up to 1.66 x 10? and 1.25 x 10? vortices Mm’? if including the events not associated to very strong downflows and encircled in red in Fig. 3..)"," We have calculated the occurrence of the detected vortex-type events in the FOV obtaining 1.46 $\times$ $^{-2}$ and 1.09 $\times$ $^{-2}$ vortices $^{-2}$ for the SST/G-band series and , respectively (the values increase up to 1.66 $\times$ $^{-2}$ and 1.25 $\times$ $^{-2}$ vortices $^{-2}$ if including the events not associated to very strong downflows and encircled in red in Fig. \ref{corchos}. .)"
 As discussed in the previous section we have computed horizontal velocities and their corresponding divergences from G-band observations., As discussed in the previous section we have computed horizontal velocities and their corresponding divergences from G-band observations.
 With the location of the detected vortices (circles in Fig. 3)), With the location of the detected vortices (circles in Fig. \ref{corchos}) )
" we define a binary mask used to compute some statistical properties in vortex-type areas (ie., a mask of vortices)."," we define a binary mask used to compute some statistical properties in vortex-type areas (i.e., a mask of vortices)."
 The vorticity is a vectorial magnitude whose direction is perpendicular to the XY plane., The vorticity is a vectorial magnitude whose direction is perpendicular to the $XY$ plane.
" Using the equation: Vx¥=Oi,-OV.a we compute vorticities of the horizontal flow, with v; notZ2? included in the computation."," Using the equation: $\nabla \times \vec v=\frac{\partial \vec v_{y}}{\partial x}-\frac{\partial \vec v_{x}}{\partial y}$, we compute vorticities of the horizontal flow, with $v_z$ not included in the computation."
 Previous works have already studied vorticity in the solar photosphere from observational (Wangatetal.2010) and theoretical (Stein&Nordlund1998;Shelyagetal.2011) approaches.," Previous works have already studied vorticity in the solar photosphere from observational \citep{wang1995,bonet2010} and theoretical \citep{stein1998,shelyag2011} approaches."
 Figure 4 (upper panels) shows histograms for the whole FOV of ~50x square Mm., Figure \ref{histogramas} (upper panels) shows histograms for the whole FOV of $\sim50 \times 50$ square Mm.
" Panel a) displays independent results for series 51 (red)) and s2 (blue)) and total values for both series (black)), resulting in a Maxwellian distribution of horizontal velocities, with a most probable value of 0.52 km s."," Panel a) displays independent results for series s1 ) and s2 ) and total values for both series ), resulting in a Maxwellian distribution of horizontal velocities, with a most probable value of 0.52 km $^{-1}$ ."
 Histograms in panel c) display a similar distribution of speeds using exclusively values within the mask of vortices., Histograms in panel c) display a similar distribution of speeds using exclusively values within the mask of vortices.
 The most probable speed reduces to 0.48 km s!., The most probable speed reduces to 0.48 km $^{-1}$.
" Median values combining s] and s2, referred to as median(,;,.,5, are shown in each panel in Figure 4.."," Median values combining s1 and s2, referred to as $_{(s1+s2)}$, are shown in each panel in Figure \ref{histogramas}."
" Independent median values of horizontal speed for 91/92 are 0.50/0.55 km s! in Figure 4aa and 0.51/0.44 km s! in Figure 4cc. Large horizontal velocities in the FOV are predominantly coming from recurrent exploding events (at the mesogranular scale) with large positive divergence values and hence the mean speed within vortex areas is expected to moderate as corresponding to negative divergences Same as with horizontal velocities and divergences, we have also plotted vorticity values over the whole FOV and in the mask of vortices (ie. in the areas inside circles in Figs."," Independent median values of horizontal speed for s1/s2 are 0.50/0.55 km $^{-1}$ in Figure \ref{histogramas}a a and 0.51/0.44 km $^{-1}$ in Figure \ref{histogramas}c c. Large horizontal velocities in the FOV are predominantly coming from recurrent exploding events (at the mesogranular scale) with large positive divergence values and hence the mean speed within vortex areas is expected to moderate as corresponding to negative divergences Same as with horizontal velocities and divergences, we have also plotted vorticity values over the whole FOV and in the mask of vortices (i.e. in the areas inside circles in Figs."
 2 and 3))., \ref{verticalvel} and \ref{corchos}) ).
 Figures 4bb) and 4dd) show the distribution of divergence lines)) and vorticity lines))., Figures \ref{histogramas}b b) and \ref{histogramas}d d) show the distribution of divergence ) and vorticity ).
 Independent results are again plotted forseries sl (red)) and 52 (blue)) andtotal values for both series (black))., Independent results are again plotted forseries s1 ) and s2 ) andtotal values for both series ).
" When considering the total FOV, the divergence distribution is centered at a slightly negative value"," When considering the total FOV, the divergence distribution is centered at a slightly negative value"
our model. we predict that«CRB afterelows originating at Lom6 are optically dark.,"our model, we predict that GRB afterglows originating at $z\ga 6$ are optically dark."
 The iuterveniug. partially neutral ICAL would eficicutly absorb the rest-frame UV afterelow that would otherwise have been redshitted iuto the optical baud (see also Fruchter 1999: Piro et al.," The intervening, partially neutral IGM would efficiently absorb the rest-frame UV afterglow that would otherwise have been redshifted into the optical band (see also Fruchter 1999; Piro et al."
 2002)., 2002).
 These bursts might give rise to the receuth discovered class of N-vav rich. CRBs (ce... Piro et al.," These bursts might give rise to the recently discovered class of X-ray rich GRBs (e.g., Piro et al."
 2002: see also Schueider et al., 2002; see also Schneider et al.
 2001) due to the redshitting of the frame 5-ravs muto the N-rayv baud., 2001) due to the redshifting of the source-frame $\gamma$ -rays into the X-ray band.
 The duration of a GRD reflects the characteristic timescale over which the ceutral cueine is active aud is therefore a diagnostic of the GRB progenitor., The duration of a GRB reflects the characteristic timescale over which the central engine is active and is therefore a diagnostic of the GRB progenitor.
 The distribution of CRB durations has been determined by the BATSE iustrument on board theObservatory (as swauuarizecdt in Paciesas et al., The distribution of GRB durations has been determined by the BATSE instrument on board the (as summarized in Paciesas et al.
 1999). and is observed to be bimodal with a population of short bursts centered on ων~ (0.35. and loug bursts around Ti.30s Usouveliotou ct al.," 1999), and is observed to be bimodal with a population of short bursts centered on $T_{\rm obs}\sim 0.3 {\rm \,s}$ , and long bursts around $T_{\rm obs}\sim 30 {\rm \,s}$ (Kouveliotou et al."
 1993)., 1993).
 For the definition of the burst duration. ων. we use the interval f time over which a GRD contaius from 5 to of its total observed 5-rav counts. also denoted as Toy iu +he literature.," For the definition of the burst duration, $T_{\rm obs}$, we use the interval of time over which a GRB contains from 5 to of its total observed $\gamma$ -ray counts, also denoted as $T_{90}$ in the literature."
 Since bursts originate over a broad ranec of redshifts. the question arises as to what the iutriusic distribution of durations is like.," Since bursts originate over a broad range of redshifts, the question arises as to what the intrinsic distribution of durations is like."
 For simplicity. we assunie in this section that BATSE was capable of sampling the full redshift distribution of GRBs shown in Figure 2.," For simplicity, we assume in this section that BATSE was capable of sampling the full redshift distribution of GRBs shown in Figure 2."
 This provides us with the maxinuun level of distortion that cosmological time dilation could have had on the observed distribution of burst durations., This provides us with the maximum level of distortion that cosmological time dilation could have had on the observed distribution of burst durations.
 The true iutriusic distribution of durations for the BATSE-trigecred bursts should lie in between the observed distribution aud the intrinsic one calculated in this section., The true intrinsic distribution of durations for the BATSE-triggered bursts should lie in between the observed distribution and the intrinsic one calculated in this section.
 The απνο of observed bursts iu a given bin ἐν with i1 observed duration Zio; aud a width Αμ cau be written as where Nus=NPANS is the total number of bursts iu the sample.," The number of observed bursts in a given bin $i$, with an observed duration $T_{{\rm obs,}i}$ and a width $\Delta T_{{\rm obs,}i}$, can be written as where $N_{\rm tot}=\sum_{i} \Delta N_{{\rm obs,}i}$ is the total number of bursts in the sample."
 We assune that the intrinsic distribution. dP/dt. is independent of redshift are satisfies Jy(dPdT)dT=1 aud (dP/dT)>0.," We assume that the intrinsic distribution, ${\rm
d}P/{\rm d}T$, is independent of redshift and satisfies $\int_{0}^{\infty}({{\rm d}P}/{{\rm d}T}){\rm d}T=1$ and $({{\rm d}P}/{{\rm
d}T}) \ge 0$."
" The intrinsic burst duration. T. is related to the observed one by the cosinological time dilation. T=T;,4,,;/(D) dk (LfAdi) is the CRB redshift distribution. as calculate in 82."," The intrinsic burst duration, $T$, is related to the observed one by the cosmological time dilation, $T=T_{{\rm obs,}i}/(1+z)$, and $({\rm d}f/{\rm d}z)$ is the GRB redshift distribution, as calculated in 2."
" We replace the iutegratiou in equation (10) bv a “παΊο, covering the range of intrinsic durations. 7. with the same umber of bius; μμ=23. as the observe histogram."," We replace the integration in equation (10) by a summation, covering the range of intrinsic durations, $T$, with the same number of bins, $N_{\rm bin}=23$, as the observed histogram."
 The inversion problem is then uniquely defined., The inversion problem is then uniquely defined.
 We carry out the deconvolution with the staucard iterative Lucy method., We carry out the deconvolution with the standard iterative Lucy method.
 This is a reliable technique. derived frou Daves theorem. to solve a set of linear equatious with additional coustraiuts on the unknowns (Lucy 1971).," This is a reliable technique, derived from Bayes' theorem, to solve a set of linear equations with additional constraints on the unknowns (Lucy 1974)."
 The stability of the algorithm is nuproved by liniting the change in the uukuowus iu each iteration. aud bv sx1000thiug over adjacent bius.," The stability of the algorithm is improved by limiting the change in the unknowns in each iteration, and by smoothing over adjacent bins."
 To this extent. we use equation (11b) in Baugh Efstatliou (1993) with parameter values of Jj=0.5 and e—0.9.," To this extent, we use equation (11b) in Baugh Efstathiou (1993) with parameter values of $\beta=0.8$ and $\epsilon=0.9$."
 We have verified that the solutions are not very sensitive to the choice of these parameters., We have verified that the solutions are not very sensitive to the choice of these parameters.
 The result of this inversion is shown in Fieure Ld. where we compare the derived intrinsic distribution to the observed one. οαςDonn)ANNorAN).," The result of this inversion is shown in Figure 4, where we compare the derived intrinsic distribution to the observed one, ${\rm d}P/{\rm d}T_{\rm
obs}(T_{{\rm obs,}i})=\Delta N_{{\rm obs,}i}/ (N_{\rm tot}\Delta T_{{\rm
obs,}i})$."
 Tt is evident that the intrinsic durations are svstematically shifted to shorter values due to the cosimological time lation., It is evident that the intrinsic durations are systematically shifted to shorter values due to the cosmological time dilation.
 The bimodality is preserved. with peaks that are narrower than the observed ones (note that the horizoutal scale is logarithniuic).," The bimodality is preserved, with peaks that are narrower than the observed ones (note that the horizontal scale is logarithmic)."
 The two star formation histories Iscussed in 82 leacl to simular intrinsic distributions., The two star formation histories discussed in 2 lead to similar intrinsic distributions.
 The shift to shorter durations. however. is more pronounced memi the case of star formation via II» cooling.," The shift to shorter durations, however, is more pronounced in the case of star formation via $_{2}$ cooling."
 The mean intrinsic durations characterising the first. short-duration. peak are ~0.058 for cooling due to atomic hydrogen. aud 0.03x for Πω cooling.," The mean intrinsic durations characterising the first, short-duration, peak are $\sim 0.05{\rm
\,s}$ for cooling due to atomic hydrogen, and $\sim 0.03{\rm \,s}$ for $_{2}$ cooling."
" The correspouding uunbers for the lone-duration peak are ~75 and ~5s. respectively,"," The corresponding numbers for the long-duration peak are $\sim 7{\rm \,s}$ and $\sim 5{\rm \,s}$, respectively."
 These differences imm the mean durations are a direct consequence of the fact that CRBs originate on average at somewhat higher redshift if Il) cooling is effective., These differences in the mean durations are a direct consequence of the fact that GRBs originate on average at somewhat higher redshift if $_{2}$ cooling is effective.
 Note that. statistically. the longest duration bursts. with Io=10005. are expected to originate at high :. iad this could be a successful selection strategy for observations tarectine highredslift GRBs.," Note that, statistically, the longest duration bursts, with $T_{\rm obs}\ga
1000{\rm \, s}$, are expected to originate at high $z$, and this could be a successful selection strategy for observations targeting high–redshift GRBs."
 The shift to longer durations due to the coswwological time dilation could iu part be compensated by the following subtle selection. effect which we ignore iu this paper., The shift to longer durations due to the cosmological time dilation could in part be compensated by the following subtle selection effect which we ignore in this paper.
 Sources at lugh + will on average have lower fluxes. aud observations with a eiven scusitivitv threshold will therefore only detect the brightestportion of the total cinission. thus systematically underestimating the true duration of the burst.," Sources at high $z$ will on average have lower fluxes, and observations with a given sensitivity threshold will therefore only detect the brightestportion of the total emission, thus systematically underestimating the true duration of the burst."
 Based on observations of CRB afterglows with known redshifts. Frail et al. (," Based on observations of GRB afterglows with known redshifts, Frail et al. ("
2001) have recently: presented,2001) have recently presented
in the outer magnetosphere (7Sry).,in the outer magnetosphere $r\lapprox r_{\rm lc}$ ).
 The parallel component of the inductive electric field (hen reappears. aud accelerates charges.," The parallel component of the inductive electric field then reappears, and accelerates charges."
 This needs to be examined in detail as an alternative {ο an outer gap for the production of pulsed > rays (Hinotorietal.2003:Takata2010).," This needs to be examined in detail as an alternative to an outer gap for the production of pulsed $\gamma$ rays \citep{Hetal03,Tetal10}."
. In a conventional pulsar wind model (Rees&Gunn1975:Arons2004) the magnetic field well bevond the light cvlinder has a predominantly toroidal component.," In a conventional pulsar wind model \citep{ReesGunn1974, Arons2004PWN} the magnetic field well beyond the light cylinder has a predominantly toroidal component."
 The rotation of the star and the outflow of the wind cause the direction of the toroidal component to reverse periodically with radial distance (Ixirketal.2007)., The rotation of the star and the outflow of the wind cause the direction of the toroidal component to reverse periodically with radial distance \citep{KirkLyubarskyPetri07}.
. Specific models (DogovalovContopoulos2005) imply a temporally changing electric field in the wind.," Specific models \citep{Bogovalov1999, Contopoulos2005} imply a temporally changing electric field in the wind."
 In such models the (Inicl theory determines the electric field., In such models the fluid theory determines the electric field.
 Although this electric field is varvingperiodically with Gime in an oblique rotator (Bogovalov1999).. implying a nonzero displacement current. this is not the displacement. current associated with the rotating dipole.," Although this electric field is varyingperiodically with time in an oblique rotator \citep{Bogovalov1999}, implying a nonzero displacement current, this is not the displacement current associated with the rotating dipole."
 The latter field would be present in vacuo. and its implicitly assumed absence in wind models requires (hat it be effectivelv screened by eiurents.," The latter field would be present in vacuo, and its implicitly assumed absence in wind models requires that it be effectively screened by currents."
 As with the inductive lield in the inner magnetosphere. if current screening is ineffective. the radiative electric field must be present in the wind zone.," As with the inductive field in the inner magnetosphere, if current screening is ineffective, the radiative electric field must be present in the wind zone."
 The neglect of the inductive electric field. Εμ. in models for a pulsar magnetosphere cannot be justified.," The neglect of the inductive electric field, ${\bi E}_{\rm ind}$, in models for a pulsar magnetosphere cannot be justified."
 This field is generated by the changing magnetic field and the associated displacement current of an obliquely rotating magnetic dipole., This field is generated by the changing magnetic field and the associated displacement current of an obliquely rotating magnetic dipole.
 It is impossible in principle to screen an inductive field by charges., It is impossible in principle to screen an inductive field by charges.
 We show that screening of the displacement current is possible in principle. and we calculate the required screening current density in an idealized model.," We show that screening of the displacement current is possible in principle, and we calculate the required screening current density in an idealized model."
 Screening of the components parallel to the magnetic field ancl perpendicular to the field involves dillerent physics., Screening of the components parallel to the magnetic field and perpendicular to the field involves different physics.
 The parallel component is unstable to the development ol oscillations. and screening can occur only in an average sense. where the average is over the oscillations.," The parallel component is unstable to the development of oscillations, and screening can occur only in an average sense, where the average is over the oscillations."
 Perpendicular screening involves a polarization current. driven by the displacement current. and can never be complete.," Perpendicular screening involves a polarization current, driven by the displacement current, and can never be complete."
 Nearly complete screening is possible [ου m/e«1. which is not the case in a pulsar magnetosphere. where the perpendicular inductive field is essentially unscreened. and unchanged from its vacuum: value.," Nearly complete screening is possible for $v_A^2\// c^2\ll1$, which is not the case in a pulsar magnetosphere, where the perpendicular inductive field is essentially unscreened, and unchanged from its vacuum value."
 The presence of Εμ has important consequences for understanding pulsar magnetospheres., The presence of ${\bi E}_{\rm ind}$ has important consequences for understanding pulsar magnetospheres.
 In particular. the magnetospheric plasma cannot be corotating with the star.," In particular, the magnetospheric plasma cannot be corotating with the star."
 The assumption of corotalion is central (o. conventional corotating-magnetosphere model [ον pulsars. ancl much of the physical interpretation of pulsar phenomena is based on (his assumption.," The assumption of corotation is central to conventional corotating-magnetosphere model for pulsars, and much of the physical interpretation of pulsar phenomena is based on this assumption."
 The inductive electric field is absent only in the artificial case of an aligned rotator (sina= 0)., The inductive electric field is absent only in the artificial case of an aligned rotator $\sin\alpha=0$ ).
 In the case of small obliquity. sina<I. an aligned rotator can be regarded as a zeroth orcer model. with the inductive field and its implications treated to first order in sina.," In the case of small obliquity, $\sin\alpha\ll1$, an aligned rotator can be regarded as a zeroth order model, with the inductive field and its implications treated to first order in $\sin\alpha$ ."
 Using such, Using such
uatural line width of the 2s-2p trausitious is therefore domiuated by the rapid decay of the 2p state and is D—34/2:=99.5 MHz.,"natural line width of the $2s$ $2p$ transitions is therefore dominated by the rapid decay of the $2p$ state and is $\Gamma =
A_{21}/ 2\pi = 99.8$ MHz."
" For a Lorentzian profile. where Ais either A, or A, and vy is the frequency. of the fine structure trausition."," For a Lorentzian profile, where $A$ is either $A_a$ or $A_b$ and $\nu_f$ is the frequency of the fine structure transition."
 From equations (11) and (12) we fiud have assumed that the Lyman decay. rate σι always dominates the collision rate Capt., From equations (11) and (12) we find where we have assumed that the Lyman decay rate $A_{21}$ always dominates the collision rate $C_{sp}n_i$.
" The absorption coefficient at line center then is 9 i be scaled according to the free-[ree absorption coellicient (Altenholl et al.1960)..qo0212njn, The ratio of the line opticaldepth (at line center). 77"". to the free-[ree continuum optical depth. 5. then is A =—3.0x104}sec for the 254,5 —2pyotransition and A. =0.11sec. for the 25,5—2ps;» trausitiou."," The absorption coefficient at line center then is This will be scaled according to the free-free absorption coefficient \citep{alt60}, , The ratio of the line opticaldepth (at line center), $\tau_\nu^{max}$, to the free-free continuum optical depth, $\tau_\nu^{ff}$ , then is where $K = -3.0 \times 10^{-4}$ $^{-1}$ for the $2s_{1/2}\rightarrow 2p_{1/2}$ transition and $K = 0.41$ $^{-1}$ for the $2s_{1/2}\rightarrow 2p_{3/2}$ transition."
 We assumed T=107 NK. and that all ionizing photons are captured by the HII region (Osterbrock1989).," We assumed $T = 10^4$ K, and that all ionizing photons are captured by the HII region \citep{ost89}."
. The equation of transfer is Note that the liue produces absorption or stimulated emission only. (through αν). but does uot produce significant spontaueous emission.," The equation of transfer is Note that the line produces absorption or stimulated emission only (through $\kappa_\nu$ ), but does not produce significant spontaneous emission."
 Thus. the emissivity is completely dominated by the [ree-[ree emissivity. Jj).," Thus, the emissivity is completely dominated by the free-free emissivity, $J_\nu^{ff}$."
 For an HII region uniform along a ray path the solution is where the result has been expressed in terms of brightuess temperature 75., For an HII region uniform along a ray path the solution is where the result has been expressed in terms of brightness temperature $T_b$.
 The liue streneth fin I) divided bythe continuum briglituess temperature is, The line strength (in K) divided bythe continuum brightness temperature is
Foundation.,Foundation.
 This research has made use of the NASA/LPAC Extragalactic Database (NED). which is operated. by the Jet Propulsion Laboratory. €'altech. under contract with the National Aeronautics and Space Administration.," This research has made use of the NASA/IPAC Extragalactic Database (NED), which is operated by the Jet Propulsion Laboratory, Caltech, under contract with the National Aeronautics and Space Administration."
density changes bv at most ~JdLOU dcTween l~--laud :zm10 relative to standard ACDAL cosinology.,density changes by at most $\sim\pm40$ between $z\approx1$ and $z\approx10$ relative to standard $\Lambda$ CDM cosmology.
" Following Cornish (2000) the change in A might shift (,,ay bY at most ~LOK 351 comparable to the οἱirent level of experimental uncertainty in this paralcter (Page et al."," Following Cornish (2000) the change in $\Lambda$ might shift $\lpeak$ by at most $\sim40$ $/35\sim1$, comparable to the current level of experimental uncertainty in this parameter (Page et al."
 2003)., 2003).
 Plivsically. the reason why the CAIB coustrains this theory less strongly than SNIa is because it proos higher redshifts where the density of matter 1s so muuch higher than that of dark energy that a uodest chanee in the latter has little effect.," Physically, the reason why the CMB constrains this theory less strongly than SNIa is because it probes higher redshifts where the density of matter is so much higher than that of dark energy that a modest change in the latter has little effect."
 Nore detailed study is warrauted. however. particularly iu lieht of the iucerease in experimental precision expected from the Plaucς nüssion.," More detailed study is warranted, however, particularly in light of the increase in experimental precision expected from the Planck mission."
 We hope to return to this Issue m future work., We hope to return to this issue in future work.
 These cosinological results are complementary ο earlier ones based ou the classical solur-svsteni tests and ones involving the equivaleice principle., These cosmological results are complementary to earlier ones based on the classical solar-system tests and ones involving the equivalence principle.
 It has been known for sole time that the basic 5D extension of ID Eiusteoi1i theory is consistent with these tests (alligas et al., It has been known for some time that the basic 5D extension of 4D Einstein theory is consistent with these tests (Kalligas et al.
 1995. Lin Overduin 2000. Overduin 2000: Wesson 2006 for review).," 1995, Liu Overduin 2000, Overduin 2000; Wesson 2006 for review)."
 It is important to recall here that standard cosmological models which are in ID inay be smoothly eiibedded in models which are in 5D. [For a review see the books bv Wessou (1999.2006): for an account of the embeddiugs from au astroplivsical viewpoint see Lachieze-Rey (2000).]," It is important to recall here that standard cosmological models which are in 4D may be smoothly embedded in models which are in 5D. [For a review see the books by Wesson (1999,2006); for an account of the embeddings from an astrophysical viewpoint see Lachieze-Rey (2000).]"
 The data may therefore be suggesting uot ouly that the universe has a fifth dimension. but that its structure may be nmch simpler than previously thought.," The data may therefore be suggesting not only that the universe has a fifth dimension, but that its structure may be much simpler than previously thought."
" The apparent magnitude n of a source at redshift + is defined by: where M is absolute iiaguitude. A(:) is the K-correction due to frequency shift aud the huuiuositv distance d,(2) is given by: Here the coustant σι aud functionὃν are defined πο that σι=OveLbYLOneO,,) and S,LY|teinWV.XssinhX]. respectively for Kk=(11.0.1}."," The apparent magnitude $m$ of a source at redshift $z$ is defined by: where $M$ is absolute magnitude, $K(z)$ is the K-correction due to frequency shift and the luminosity distance $\dL(z)$ is given by: Here the constant $\sigma_k$ and function${\cal S}_k$ are defined so that $\sigma_k\equiv\{\sqrt{\Omato-\Olamo-1},1,\sqrt{1-\Omato-\Olamo}\}$ and ${\cal S}_k[X]\equiv\{\sin X,X,\sinh X\}$ respectively for $k=\{+1,0,-1\}$."
 Eq. (1511 , Eq. \ref{appmag}) )
"contains terms such as AY. A(:) aud JT, that are independent of the background cosnologv aud hence not of interest to us here."," contains terms such as $M$, $K(z)$ and $\Ho$ that are independent of the background cosmology and hence not of interest to us here."
" We remove those terms by focusing. not on the apparcut magnitude itself. but on the difference or “residual” magnitude relative to a fiducial model. which we take here to be the standard flat ACDAL model with WMLADP values of O,,, aud O,,."," We remove those terms by focusing, not on the apparent magnitude itself, but on the difference or “residual” magnitude relative to a fiducial model, which we take here to be the standard flat $\Lambda$ CDM model with WMAP values of $\Omato$ and $\Olamo$ ."
 That is. we focus ou the observational quautity," That is, we focus on the observational quantity"
"Iu Appeucdix ?°? we present X. T. and 7 behaviors iu accretion disk with the dominant background viscosity a, and opacity iu the form (9)).","In Appendix \ref{app1} we present $\Sigma$, $T$, and $\tau$ behaviors in accretion disk with the dominant background viscosity $\alpha_\nu$ and opacity in the form \ref{eq:opacity}) )."
 If we now use these expressions [equations (À3)). (AG)). (A12)). and (A17))| to calculate Toomre Q in the selfluuinous viscous part of the disk at iuy we find that >1 in the optically thick case and o«l πι the optically thin case.," If we now use these expressions [equations \ref{eq:Sig_eq_thick}) ), \ref{eq:T_eq_thick}) ), \ref{eq:Sig_eq_thin}) ), and \ref{eq:T_eq_thin}) )] to calculate Toomre $Q$ in the self-luminous viscous part of the disk at $\omega>\omega_\nu$ we find that >1 in the optically thick case and <1 in the optically thin case."
 Apparently. for any reasonable dust opacity behavior (Q starts to deviate yon its mareinal stabilitv value (Qu towards higher values at the transition from the eravitoturbuleut to viscous regime (at iw= v).," Apparently, for any reasonable dust opacity behavior $Q$ starts to deviate from its marginal stability value $Q_0$ towards higher values at the transition from the gravitoturbulent to viscous regime (at $\omega=\omega_\nu$ )."
 Thus. accretion disk caunot © evavitationally uustable if its augular momentum ransport is uot donunated by eravitational torques mt is rather due to some other form of cffective viscosity.," Thus, accretion disk cannot be gravitationally unstable if its angular momentum transport is not dominated by gravitational torques but is rather due to some other form of effective viscosity."
 This consideration demonstrates explicitly how he transition between the exavitoturbuleut aud viscous warts of the disk occur., This consideration demonstrates explicitly how the transition between the gravitoturbulent and viscous parts of the disk occur.
 Note that a¢y given by equation (69)) is independent of v., Note that $\alpha_{GI}$ given by equation \ref{eq:alph}) ) is independent of $\omega$.
 This implies that ifTow accretion disk renuiünus viscous evervwlhere aud does not transition iuto the eravitoturbuleut state at all.," This implies that if, accretion disk remains viscous everywhere and does not transition into the gravitoturbulent state at all."
 Thus. at low enough i a eravitationally stable viscous disk can transport matter all the way to the ceutral object from infinitely laree distances.," Thus, at low enough $\dot m$ a gravitationally stable viscous disk can transport matter all the way to the central object from infinitely large distances."
 Condition. (91)) Is more restrictive than the coustraimt (82)) which delineates reenue in which the eravitoturbuleut constant AL disk can fransfer imaterial from infinitely Large distances without fragmentation., Condition \ref{eq:dotm_visc_limit}) ) is more restrictive than the constraint \ref{eq:stable}) ) which delineates regime in which the gravitoturbulent constant $\dot M$ disk can transfer material from infinitely large distances without fragmentation.
 Our results derived in 83 apply fo a varicty of situations in which eravitoturbuleut disks cau exist: they can be optically thick or thin. fraeiueutiug or not.," Our results derived in \ref{sect:mdot} apply to a variety of situations in which gravitoturbulent disks can exist: they can be optically thick or thin, fragmenting or not."
" Tere we classify different states in which accretiou disks can be found according to the values of M. background viscosity a, and irradiation temperature Ty."," Here we classify different states in which accretion disks can be found according to the values of $\dot M$, background viscosity $\alpha_\nu$ and irradiation temperature $T_0$."
 We also describe applications of these results to real astroplivsical svstcus and discuss their connection with the work of others., We also describe applications of these results to real astrophysical systems and discuss their connection with the work of others.
 As we saw du previous sections different parts of accretion disks can be characterized by different states eravitoturbulent. viscous. iradiated. etc.," As we saw in previous sections different parts of accretion disks can be characterized by different states — gravitoturbulent, viscous, irradiated, etc."
 Transitious between these reeines depend ou AL. Ty. aud αν.," Transitions between these regimes depend on $\dot M$, $T_0$, and $\alpha_\nu$."
 Our results allow us to single out three important cases for the different transition topologics., Our results allow us to single out three important cases for the different transition topologies.
 When external irradiation 1s strong. TyTjLl. the phase space of possible regimes cau be represcuted by Figure L..," When external irradiation is strong, T_0/T_f>1, the phase space of possible regimes can be represented by Figure \ref{fig:CaseA}."
 Caven that in the case of opacity dominated by cold dust Ty is just slightly higher than 10 Is (see eq. [][), Given that in the case of opacity dominated by cold dust $T_f$ is just slightly higher than $10$ K (see eq. \ref{eq:num_values}] ])
 it is clear that the condition. (95)) should apply to virtually all aceretion disks in the Universe as even the lowest measured teniperatures encouutered i sonie deuse molecular cores are not very different from 10 I (Di Francesco 2007)., it is clear that the condition \ref{eq:CaseA}) ) should apply to virtually all accretion disks in the Universe as even the lowest measured temperatures encountered in some dense molecular cores are not very different from $10$ K (Di Francesco 2007).
 Using Fieure 1l/ one can examine different reginies (indicated. by shading) which are relevant for a given AL., Using Figure \ref{fig:CaseA} one can examine different regimes (indicated by shading) which are relevant for a given $\dot M$.
" separating different regimes correspond to various critical transitions; 9=a, (dotted limes) iuplics a transition from viscous (above aud to the left of this curve. slaut solid shading) to eravitoturbuleut state. a=1 (solid lines) describes the onset of fragmentation (below and to the right of this curve. vertical solid shacine). T—Ty is a (long«lashed) line below which external iraciation starts to dominate disk structure (region witli horizontal dotted shadiug). 7=1 curve (short-dashed) shows a transition from au optically thick (above this curve) to an optically thin disk."," separating different regimes correspond to various critical transitions: $\alpha=\alpha_\nu$ (dotted lines) implies a transition from viscous (above and to the left of this curve, slant solid shading) to gravitoturbulent state, $\alpha=1$ (solid lines) describes the onset of fragmentation (below and to the right of this curve, vertical solid shading), $T=T_0$ is a (long-dashed) line below which external irradiation starts to dominate disk structure (region with horizontal dotted shading), $\tau=1$ curve (short-dashed) shows a transition from an optically thick (above this curve) to an optically thin disk."
 Coordinates Gn. phase space Mee. upper aud right axes) of the points A. D. C. D where these curves cross are giveu bv expressions (B93) in Appeudix ??..," Coordinates (in phase space $\dot m,\omega$ , upper and right axes) of the points A, B, C, D where these curves cross are given by expressions \ref{eq:CaseA_points}) ) in Appendix \ref{app2}."
 Uushlided region correspouds to a self-Iuminous. eravitoturbulent disk.," Unshaded region corresponds to a self-luminous, gravitoturbulent disk."
 A constant AL disk corresponds to a straight vertical line in Figure 1. cutting through citferent regines., A constant $\dot M$ disk corresponds to a straight vertical line in Figure \ref{fig:CaseA} cutting through different regimes.
" The umunerical values ou the left aud lower axes of this Figure correspoud. to a particular choice of a,=0.002. ToTg=3 (Ty=35 I) aud opacity in the form (10))."," The numerical values on the left and lower axes of this Figure correspond to a particular choice of $\alpha_\nu=0.003$, $T_0/T_f=3$ $T_0=35$ K) and opacity in the form \ref{eq:dust_opacity}) )."
" As described i 83.5 at the very low A=10 ""ALL | disk is always viscous and eravitationally stable. even at arbitrarily large distances from the central object."," As described in \ref{subsect:bkgvisc} at the very low $\dot M\lesssim 10^{-7}$ $_\odot$ $^{-1}$ disk is always viscous and gravitationally stable, even at arbitrarily large distances from the central object."
 Above this value of M the disk must be eravitoturbulent within sone range of distauces., Above this value of $\dot M$ the disk must be gravitoturbulent within some range of distances.
 For 100* AL. veο<afxE10 OA tf eravitoturbulent region extends from Q~lo?10ὃς ball the way to very large distances (where Q> 0).," For $10^{-7}$ $_\odot$ $^{-1}\lesssim\dot M\lesssim 4\times 10^{-5}$ $_\odot$ $^{-1}$ gravitoturbulent region extends from $\Omega\sim 
10^{-10}-10^{-9}$ $^{-1}$ all the way to very large distances (where $\Omega\to 0$ )."
" However. for AT above ALzds10 ON, Lb eravitoturbulent disk fragments at Q=1.1«10.1"" ! (iuthe optically thick case with opacity characterized by 3=2 fragiuentation occurs exactly at Og. see 83.1) so that a constaut AY disk cau be maintained onlv interior to this poiut."," However, for $\dot M$ above $\dot M\gtrsim 
4\times 10^{-5}$ $_\odot$ $^{-1}$ gravitoturbulent disk fragments at $\Omega\approx 1.4\times 10^{-10}$ $^{-1}$ (in the optically thick case with opacity characterized by $\beta=2$ fragmentation occurs exactly at $\Omega_f$ see \ref{subsect:frag}) ) so that a constant $\dot M$ disk can be maintained only interior to this point."
 For auv M=10 “AL. vr! the eravitoturbuleut augular momentum transport becomes weak at laree enough © so that the backeround viscosity starts to determine disk properties at small radii (upper part ofthe plot. at O2LO? Ἐ in this particular case).," For any $\dot M\gtrsim 10^{-7}$ $_\odot$ $^{-1}$ the gravitoturbulent angular momentum transport becomes weak at large enough $\Omega$ so that the background viscosity starts to determine disk properties at small radii (upper part of the plot, at $\Omega>10^{-9}$ $^{-1}$ in this particular case)."
" As an example.let us use Figure 1. to figure out where the trausitious between differeut reginues occur in adisk with AF=10 M. | n,= 0.003."," As an example,let us use Figure \ref{fig:CaseA} to figure out where the transitions between different regimes occur in adisk with $\dot M=10^{-5}$ $_\odot$ $^{-1}$ , $\alpha_\nu=0.003$ ,"
the high precision intensity measurements described by (he MOST space mission (Walker et al.,the high precision intensity measurements described by the MOST space mission (Walker et al.
 2003)., 2003).
 In contrast with the ground based observations. (he MOST space telescope has failecl to detect p-modes from Procyon (Matthews et al.," In contrast with the ground based observations, the MOST space telescope has failed to detect $p$ -modes from Procyon (Matthews et al."
 2004) at its instrumental limit., 2004) at its instrumental limit.
 This paper focusses on modeling the dvnamies and physical conditions in the outer convective lavers of Procyon. aud in making a comparison wilh the Sun.," This paper focusses on modeling the dynamics and physical conditions in the outer convective layers of Procyon, and in making a comparison with the Sun."
 We present the results of a 3D radiative hydrodynamic numerical large eddy simulation of the outer lavers of Procyon. and compare il (o a similar simulation perlormed previously lor the Sun using the same detailed physics input and numerical code (Robinson et al.," We present the results of a 3D radiative hydrodynamic numerical large eddy simulation of the outer layers of Procyon, and compare it to a similar simulation performed previously for the Sun using the same detailed physics input and numerical code (Robinson et al."
 2003)., 2003).
 The domains of both simulations extend [rom the observable atiosphleric lavers in radiative equilibrium and (he outer lavers of the convection zone. down to a depth al which the temperature eradient exceeds (he acliabatic value by only a very small amount (deep convection).," The domains of both simulations extend from the observable atmospheric layers in radiative equilibrium and the outer layers of the convection zone, down to a depth at which the temperature gradient exceeds the adiabatic value by only a very small amount (deep convection)."
 The simulation domain in both simulations includes the highly superacdiabatie transition laver (hereafter SAL) between the atmosphere ancl deep convection., The simulation domain in both simulations includes the highly superadiabatic transition layer (hereafter SAL) between the atmosphere and deep convection.
 It is in the SAL region that (he p-modes are believed (to be excited by stochastic processes., It is in the SAL region that the $p$ -modes are believed to be excited by stochastic processes.
 In the next section. we describe briellv the starting model and the procedure followed in carrving out the 3D numerical simulation.," In the next section, we describe briefly the starting model and the procedure followed in carrying out the 3D numerical simulation."
 Section 3 discusses (he main results from the Procvon simulation. emphasizing the striking dillerences in the physical conditions in the atmospheres of Procvon and the Sun.," Section 3 discusses the main results from the Procyon simulation, emphasizing the striking differences in the physical conditions in the atmospheres of Procyon and the Sun."
 The most surprising result is that the SAL. which in (he Sun changes little with (ime and is placed below the photosphere. is found to vary with time in position and amplitude in Procvon.," The most surprising result is that the SAL, which in the Sun changes little with time and is placed below the photosphere, is found to vary with time in position and amplitude in Procyon."
 As a result. the outer lavers of the SAL in Procvon are found part of the time in the optically thin atmosphere in Procvon.," As a result, the outer layers of the SAL in Procyon are found part of the time in the optically thin atmosphere in Procyon."
 Iu this region. radiative losses are large. and characteristic timescales are short. in contrast. with ihe Sun where the SAL is subphotospheric.," In this region, radiative losses are large, and characteristic timescales are short, in contrast with the Sun where the SAL is subphotospheric."
 The implications of these fundamental differences For the excitation. damping aud lifetimes ol p-mode oscillations. are discussed in general terms in Section 4. where we analvze the likely consequences for p-moce observations in Procvon. and the conditions for the detectability of p-mode oscillations either from (he intensity measurements [rom space (MOST mission). or from ground-based radial velocity measurements.," The implications of these fundamental differences for the excitation, damping and lifetimes of $p$ -mode oscillations, are discussed in general terms in Section 4, where we analyze the likely consequences for $p$ -mode observations in Procyon, and the conditions for the detectability of $p$ -mode oscillations either from the intensity measurements from space (MOST mission), or from ground-based radial velocity measurements."
 In Section 5. we summarize our resulis and the direction of future research in the light ol our stud.," In Section 5, we summarize our results and the direction of future research in the light of our study."
"Iu the last few decades. cosmology has become a data-driven field. where lieh-precision iieasurenienuts of he cosmic microwave backeround (CAIB.ο,ον,2).. woak casing (c.g.?) aud galaxy survevs (e.g.?) lave xxnitted the establishment of a standard: cosmological nodel in which the Universe is composed of barvous. dark matter aud dark cnerey,","In the last few decades, cosmology has become a data-driven field, where high-precision measurements of the cosmic microwave background \citep[CMB, e.g.,][]{wmap7}, weak lensing \citep[e.g.,][]{Schrabback:2010} and galaxy surveys \citep[e.g.,][]{Percival:2007} have permitted the establishment of a standard cosmological model in which the Universe is composed of baryons, dark matter and dark energy."
 Some Ulajor questions remain. the nature of dark matter aud dark enerev in particular is still not understood.," Some major questions remain, the nature of dark matter and dark energy in particular is still not understood."
 Similarly. the initial conditions ofthe Universe are vet to be established aud alternative models of eravity are still to be tested iu conrparison with Eiusteiu's general relativity.," Similarly, the initial conditions of the Universe are yet to be established and alternative models of gravity are still to be tested in comparison with Einstein's general relativity."
" Now surveys are underway with these science objectives. e.g. Planck for the CMD (?).. DES ergvSurvey. ο), BOSS (BarvonOscillationSurvey. ?).. LSST (LareeSvnopticSurveyTelescope.?) and Euclid (??) for weals Iensine and the study of large-scale structure with ealaxy survevs."," New surveys are underway with these science objectives, e.g. Planck for the CMB \citep{Planck}, , DES \citep[Dark Energy Survey,][]{DES:2005}, , BOSS \citep[Baryon Oscillation Spectroscopic Survey,][]{Schlegel:2007}, LSST \citep[Large Synoptic Survey Telescope,][]{LSST} and Euclid \citep{2011arXiv1110.3193L, Euclidsb} for weak lensing and the study of large-scale structure with galaxy surveys."
 In order to be beneficial. cosmological studies of these survevs need to use high-precision statistical methods. such as a full 3D analysis ou the sky where all-sky 3D survevs are available.," In order to be beneficial, cosmological studies of these surveys need to use high-precision statistical methods, such as a full 3D analysis on the sky where all-sky 3D surveys are available."
 Several tools have been developed to analyse data on the sphere. which is required for a 2D spherical harmonic CXMB analysis (22777)..," Several tools have been developed to analyse data on the sphere, which is required for a 2D spherical harmonic CMB analysis \citep{igloo,igloo2,healpix:2002,gorski:2004by,glesp}."
" Weals lensing and galaxy survey data can also be analysed tomoeraphlically νο, in 2D slices). but uulike for the CAIB. a full 3D spherical Fouricr-Bessel analysis can also be sought (?22777).."," Weak lensing and galaxy survey data can also be analysed tomographically (i.e. in 2D slices), but unlike for the CMB, a full 3D spherical Fourier-Bessel analysis can also be sought \citep{Fisher:1995,Heavens:1995,heavens3d,weak3d,2011arXiv1112.3100R}."
 Previous 3D data analyses were ou relatively small data sets (2??7).. but future surveys Like Euclid aud LSST will provide surveys with billious of galaxies. making previous methods for caleulatiug the 3D spectra unfeasiblv tine-coustuuine.," Previous 3D data analyses were on relatively small data sets \citep{Fisher:1995,Heavens:1995,Erdogdu:2006dv,Erdogdu:2005wi}, but future surveys like Euclid and LSST will provide surveys with billions of galaxies, making previous methods for calculating the 3D spectra unfeasibly time-consuming."
 Iun Section 7?7.. we present the theory behind the 3D Fourier-Bessel decomposition for infinite aud finite continuous fields as well as the usual iuiethod for a discrete survey (eg. galaxy survey).," In Section \ref{sub:FB}, we present the theory behind the 3D Fourier-Bessel decomposition for infinite and finite continuous fields as well as the usual method for a discrete survey (e.g. galaxy survey)."
 Iu section ?7.. wo present two additional equivalent formmlatious of the spherical Fourier-Dessel decomposition. one of which is ceutral to the 3DEN code.," In section \ref{sub:3equiv}, we present two additional equivalent formulations of the spherical Fourier-Bessel decomposition, one of which is central to the 3DEX code."
 In Section 3.. WO conipare the acctirmaev ud calculation time for the usual method used for calculating Fourier-Dessel cocfiicicuts and methods with the 3DEX code presented in this paper.," In Section \ref{sec:Methods Comparison}, we compare the accuracy and calculation time for the usual method used for calculating Fourier-Bessel coefficients and methods with the 3DEX code presented in this paper."
 Iu Section L. we describe the 3DEN library and eive examples of how ) use it.," In Section \ref{sec:3DEX}, we describe the 3DEX library and give examples of how to use it."
 In Section 5 we present our conclusions., In Section \ref{sec:Conclusion} we present our conclusions.
 We also oeiclude an appendix. where we discuss the subtleties of +ιο Fouricr-Bessel normalisation.," We also include an appendix, where we discuss the subtleties of the Fourier-Bessel normalisation."
 Iu observational cosimology. spherical coordinates Qvhere the observer is at the origi) are a natural choice or the analysis of cosmological fields.," In observational cosmology, spherical coordinates (where the observer is at the origin) are a natural choice for the analysis of cosmological fields."
 Iu this svstem of coordinates. eieenfunctious of the Laplacian operator are products of spherical Bessel fictionsaud spherical iunionies. Le. functious {τω with cigenvalucs f+," In this system of coordinates, eigenfunctions of the Laplacian operator are products of spherical Bessel functionsand spherical harmonics, i.e. functions $j_\ell (kr)Y_{\ell m}$ with eigenvalues $-k^2$."
 For an homogeneous threc-dinensioual field fu)= ina flat eeomoetry. the spherical Fourier-Dessel decomposition (277) is withthe inverse relation The coctficieuts way be used to caleulatethe 3D power," For an homogeneous three-dimensional field $f(\textbf{r}) = f(r,\theta,\phi)$ in a flat geometry, the spherical Fourier-Bessel decomposition \citep{Fisher:1995,heavens3d,weak3d} is withthe inverse relation The coefficients may be used to calculatethe 3D power"
"~80 z20.83 z>2.8 0,,20.3. O4=0. M/L 0.30.5. z=0.5 z20. ςΞ0.8. z0.4 ",$\sim 80$ $z=0.83$ $z>2.8$ $\Omega_m =0.3$ $\Omega_{\Lambda}=0$ $M/L$ $0.3< z < 0.5$ $z=0.5$ $z=0$ $z=0.8$ $z \approx 0.4$ 
with ZCXPM) <15.5.,with $E$ (APM) $\le 18.8$.
 In. EDBOS-2 (White et al., In FBQS-2 (White et al.
 2000. αμα {ος 17.8) 90 per cent of the QSO candidates were spectroscopically classified. and these included: 636. QSOs. none of which had z=3.5. although 12 had 3<z«3.5.," 2000, $O-E \le 2$ and $E \le 17.8$ ) 90 per cent of the QSO candidates were spectroscopically classified, and these included 636 QSOs, none of which had $z>3.5$, although 12 had $3 < z <3.5$."
 We are therefore confident that the number of QSOs we find with 3.7<2<odd is limited only by the radio and £ band Hux density limits. with no QSOs missed due to the colour selection.," We are therefore confident that the number of QSOs we find with $3.7 \le z \le 4.4$ is limited only by the radio and $E$ band flux density limits, with no QSOs missed due to the colour selection."
 The survey reported here can be considered Complementary to PBQS-2., The survey reported here can be considered complementary to FBQS-2.
 Both surveys select FIRST radio sources with starlike optical counterparts. but the colour criteria are Ofe<2? for FBQS-2 (bluc-exeess objects) and OLe52 for ur work.," Both surveys select FIRST radio sources with starlike optical counterparts, but the colour criteria are $O-E \le 2$ for FBQS-2 (blue-excess objects) and $O-E \ge 2$ for our work."
 In PBQS-2 objects were selected with O|£x210 liminate galaxies. since catalogues based on photographic dates provide poor discrimination between stellar ancl non-μαellar objects.," In FBQS-2 objects were selected with $O-E \le 2$ to eliminate galaxies, since catalogues based on photographic plates provide poor discrimination between stellar and non-stellar objects."
 The survey. selection criteria. and the QSO ractions as à function of redshift. are compared in Table 3.," The survey selection criteria, and the QSO fractions as a function of redshift, are compared in Table 3."
 Figs., Figs.
 Ga and 6b show O—.E versus A and O|FE versus z for 16 QSOs in the two surveys., 6a and 6b show $O-E$ versus $E$ and $O-E$ versus $z$ for the QSOs in the two surveys.
 Whereas the redshifts covered w EDBOS-2 range up to 3.4. our sample includes redshifts up ο 4.3 (Fig.," Whereas the redshifts covered by FBQS-2 range up to 3.4, our sample includes redshifts up to 4.3 (Fig."
 6b. Table 3). and this is à consequence of both 1e colour cut and the fainter magnitudes.," 6b, Table 3), and this is a consequence of both the colour cut and the fainter magnitudes."
 The ellicieney. of QSO selection in this work (51/78 = 65 per cent) is slightly larger than in FDOS-2 (636/1130 = 56 per cent)., The efficiency of QSO selection in this work (51/78 = 65 per cent) is slightly larger than in FBQS-2 (636/1130 = 56 per cent).
 Although our selection criteria filter out most of the QSOs below z=3.5. and these are covered by FBQS- our survey is better at rejecting galaxies. since we take the morphological information from SDSS whereas EDOS-2 used the less reliable APM discrimination.," Although our selection criteria filter out most of the QSOs below $z=3.5$, and these are covered by FBQS-2, our survey is better at rejecting galaxies, since we take the morphological information from SDSS whereas FBQS-2 used the less reliable APM discrimination."
 Although our sample has fewer sources. the ellicieney for moderate to high redshift QSOs is high: 20/78 = 26 per cent for 2xz<3.7 compared to 65/1130 = 6 per cent for FBQS-2.," Although our sample has fewer sources, the efficiency for moderate to high redshift QSOs is high: 20/78 = 26 per cent for $2 \le z \le 3.7$ compared to 65/1130 = 6 per cent for FBQS-2."
 M 33.7 the cllicieney of our selection is 10/78 = 13 per cent., At $z \ge 3.7$ the efficiency of our selection is 10/78 = 13 per cent.
 The maximum redshift in EFDOS-2 is 2=3.4., The maximum redshift in FBQS-2 is $z=3.4$.
 Table 4 compares the sizes and. search efficiencies of various radio-selected. high-redshift QSO s;mples., Table 4 compares the sizes and search efficiencies of various radio-selected high-redshift QSO samples.
 The success-rate for high-z QSOs and the number of QSOs in this work are comparable to or larger than those of previous searches for racdio-selected QSOs at much.higher radio Lux densities., The success-rate for $z$ QSOs and the number of QSOs in this work are comparable to or larger than those of previous searches for radio-selected QSOs at much radio flux densities.
 Twenty-six of the QSOs with spectra have redshifts 2< 4. πο that Siiv 1397 and ο 1549 are included in the observed wavelength range.," Twenty-six of the QSOs with spectra have redshifts $ 2.0 \le z \le
4.4$ , so that Si 1397 and C 1549 are included in the observed wavelength range."
" Seven of these QSOs show strong broad. absorption lines (BALS) with velocity widths above 2000 kim "" . ⋜⋯∠⇂↓⋅∢⋅⋯∼↓⊔⊔⋏∙≟∖⇁∢⋅↓⋯∙↓↿⊓⋅≱∖⋜↧∣⋯∖⇁⋖⋅∆∫≻∪∪∪↳⊔↓τον ⋅ + and were classified as BALs."," Seven of these QSOs show strong broad absorption lines (BALs) with velocity widths above 2000 km $^{-1}$ and reaching velocities above 3000 km $^{-1}$, and were classified as BALs."
 The balnicity index (BL) of cach BAL was computed. following the prescription. of Wevmann ct al. (, The balnicity index (BI) of each BAL was computed following the prescription of Weymann et al. (
1991) ancl is presented in Table 5.,1991) and is presented in Table 5.
 The fraction of BALs in our sample is therefore 7/26 = 217x10 per cent., The fraction of BALs in our sample is therefore 7/26 = $27 \pm 10$ per cent.
 For comparison. from the study of the BAL QSOs in FBOS-2 (Becker et al.," For comparison, from the study of the BAL QSOs in FBQS-2 (Becker et al."
 2000) we derive an observed raction of BALs for 1.5xz3 and 15<ExITS of 14/134 = 1043 per cent (BALs with DIO excluded)., 2000) we derive an observed fraction of BALs for $1.5 \le z \le 3$ and $15 \le E \le 17.8$ of 14/134 = $10 \pm 3$ per cent (BALs with BI=0 excluded).
 In view of the redder colours found for BALS. compared to non-DALs. and the adopted colour selection. O—Ex2. Decker et al.," In view of the redder colours found for BALs, compared to non-BALs, and the adopted colour selection, $O-E \le 2$, Becker et al."
 conclude that the BAL fractions derived from their work should be considered as lower limits., conclude that the BAL fractions derived from their work should be considered as lower limits.
 Our result. 2743:10 »er cent. confirms a larger fraction of BALs in a red-selected sample (Off> 2). and following their argument this value should be considered as an upper limit to the actual fraction.," Our result, $27 \pm 10$ per cent, confirms a larger fraction of BALs in a red-selected sample $O-E \ge 2$ ), and following their argument this value should be considered as an upper limit to the actual fraction."
satellites that would achieve these very different. goals. and (he characteristics derived from fairly mature engineering work were almost identical: same aperture. same IR pixel scale. same IR camera size. same orbit.,"satellites that would achieve these very different goals, and the characteristics derived from fairly mature engineering work were almost identical: same aperture, same IR pixel scale, same IR camera size, same orbit."
 The only significant difference is that the WL satellite requires an oplical camera in addition to IR., The only significant difference is that the WL satellite requires an optical camera in addition to IR.
 Even the fields are complementary: WL looks ab high-latitude fields while microlensing looks at the Galactic bulge., Even the fields are complementary: WL looks at high-latitude fields while microlensing looks at the Galactic bulge.
 So thev could amicably share (ime on the same satellite. with microlensing observations taking place during the 3J) mouths per vear (hat the camera cannot view high-latitude fields and so must observe Galactic-plane targets.," So they could amicably share time on the same satellite, with microlensing observations taking place during the 3 months per year that the camera cannot view high-latitude fields and so must observe Galactic-plane targets."
 When I sav. “completely independent. I mean (hat there was absolutely no contact between the science/engineering teams thal developed these designs.," When I say “completely independent”, I mean that there was absolutely no contact between the science/engineering teams that developed these designs."
 Indeed they did not even know of each others’ existence until a French astronomer in contact with both DUNE (WL) and MIPF (microlensing) put them in contact with each other., Indeed they did not even know of each others' existence until a French astronomer in contact with both DUNE (WL) and MPF (microlensing) put them in contact with each other.
 Twill not go into detail about the designs., I will not go into detail about the designs.
 This is properly the subject of an REI paper., This is properly the subject of an RFI paper.
 llere I am just locusine on the proper scientific approach to (wo big problems., Here I am just focusing on the proper scientific approach to two big problems.
 AMicrolensing is potentially the most powerful method of finding planets., Microlensing is potentially the most powerful method of finding planets.
 It is the only method sensitive to analogs of all M20.1AL solar-svstem planets (no method is sensitive to Mercuries): the only method sensitive to. Mars-mnass planets in the habitable zone: the onlv method sensitive to old free-floating planets: the only method (hat is sensitive to planets independent of (the mass of their hosts: and the only method that has good sensitivity (o planets in (wo major Galaxy environments (disk and bulge).," It is the only method sensitive to analogs of all $M>0.1\,M_\oplus$ solar-system planets (no method is sensitive to Mercuries); the only method sensitive to Mars-mass planets in the habitable zone; the only method sensitive to old free-floating planets; the only method that is sensitive to planets independent of the mass of their hosts; and the only method that has good sensitivity to planets in two major Galaxy environments (disk and bulge)."
 That is. while microlensine certainlv does not supersede all other methods. it is the best single method for conducting a svslenmatic survey of planets as a Function of planet mass. host mass. host-planet separatioΕν and position in the Galaxy.," That is, while microlensing certainly does not supersede all other methods, it is the best single method for conducting a systematic survey of planets as a function of planet mass, host mass, host-planet separation, and position in the Galaxy."
 Given the strong claims just made. why has microlensing so far discovered. only 14 planets (8 published + 6 in prep). while (ransits have discovered dozens aud RV has discovered hundreds?," Given the strong claims just made, why has microlensing so far discovered only 14 planets (8 published + 6 in prep), while transits have discovered dozens and RV has discovered hundreds?"
 The short answer is that to fully achieve (he above potential requires a wide-field Wager in space., The short answer is that to fully achieve the above potential requires a wide-field imager in space.
" Today. microlensing is already. making some remarkable discoveries about. planets: first detection of ""cold Neptunes’. [ist Sun/Jupiter/Saturn analog. aud lowest-mass planet around a normal star."," Today, microlensing is already making some remarkable discoveries about planets: first detection of “cold Neptunes”, first Sun/Jupiter/Saturn analog, and lowest-mass planet around a “normal” star."
 But when considering these discoveries. it is important to keep in mind (hat they are being made by lin class telescopessmaller.," But when considering these discoveries, it is important to keep in mind that they are being made by 1m class telescopes."
 Indeed. amateurs (equipped with 2540cm telescopes) made major contributions to the detection of 8 out of 14 microlensing," Indeed, amateurs (equipped with 25–40cm telescopes) made major contributions to the detection of 8 out of 14 microlensing"
"between J, and J,- found using this method is similar to the average ratio found from spectral fitting. about3%.","between $I_\mathrm{A}$ and $I_\mathrm{U}$ found using this method is similar to the average ratio found from spectral fitting, about."
. Using this best fit model. Nyy as a function of time is found by projecting the measured colors onto the track.," Using this best fit model, $N_\mathrm{H}$ as a function of time is found by projecting the measured colors onto the track."
 The projection provides slightly different values of Ny for each color-color diagram examined., The projection provides slightly different values of $N_\mathrm{H}$ for each color-color diagram examined.
 Major discrepancies are due to projection onto a wrong part of the model curve in the region where the curve overlaps with itself., Major discrepancies are due to projection onto a wrong part of the model curve in the region where the curve overlaps with itself.
 To even out these discrepancies wecalculated the median of Vy from all three color-color diagrams used in our analysis., To even out these discrepancies wecalculated the median of $N_\mathrm{H}$ from all three color-color diagrams used in our analysis.
 The time development of .Vjj found from the color-color diagramsis in good agreement with the Απ resulting from the spectral fits (Fig. 7))., The time development of $N_\mathrm{H}$ found from the color-color diagramsis in good agreement with the $N_\mathrm{H}$ resulting from the spectral fits (Fig. \ref{fig:nh}) ).
 The irregular variations observed in the lightcurve during the dips are generally thought to be due to the cloudy structure of an absorber extending above the disk surface., The irregular variations observed in the lightcurve during the dips are generally thought to be due to the cloudy structure of an absorber extending above the disk surface.
 The location of the material. r.e.. its distance from the neutron star. however. is unknown and depends on model assumptions.," The location of the material, i.e., its distance from the neutron star, however, is unknown and depends on model assumptions."
 According to Crosa&Boynton(1950).. the obscuring matter is. the temporarily thickened disk rim. or. in the modification of this model by Bochkarev(1989) and Bochkarev&Karitskaya (1989).. consists of “blobs” in an extended corona above the disk rim.," According to \citey{Crosa80}, the obscuring matter is the temporarily thickened disk rim, or, in the modification of this model by \citey{Bochkarev89.1} and \citey{Bochkarev89.2}, consists of “blobs” in an extended corona above the disk rim."
 On the other hand. in the coronal wind model of Schandletal.(1997) the dips are caused by a spray of matter at some inner disk radius where the accretion stream impacts on the disk.," On the other hand, in the coronal wind model of \citey{Schandl97.1} the dips are caused by a spray of matter at some inner disk radius where the accretion stream impacts on the disk."
 Since all models assume that the absorbing matter exists in the form of clumps of material. it is interesting to ask whether these blobs can be seen in the data.," Since all models assume that the absorbing matter exists in the form of clumps of material, it is interesting to ask whether these blobs can be seen in the data."
 The symmetry of the structures in the first half of the lighteurve shown in Fig., The symmetry of the structures in the first half of the lightcurve shown in Fig.
 4aa suggests that this is indeed the case., \ref{fig:timeevol}a a suggests that this is indeed the case.
 We therefore tried to model this part of the dip lighteurve by fitting Gaussians to each of the visibly identified structures., We therefore tried to model this part of the dip lightcurve by fitting Gaussians to each of the visibly identified structures.
 The best fit to the data is displayed in Fig. 8.., The best fit to the data is displayed in Fig. \ref{fig:lcfit}.
 Information about the parameters of each Gaussian is summarized in Table 2 (available in electronic form only)., Information about the parameters of each Gaussian is summarized in Table 2 (available in electronic form only).
 Note that the structures appear to be quasi-pertodic on a time-scale of about 10.? dd ss) which could provide a hint on their nature., Note that the structures appear to be quasi-periodic on a time-scale of about $10^{-3}$ d s) which could provide a hint on their nature.
 A similar ss periodicity has previously been reported by Leahyetal.(1992)., A similar s periodicity has previously been reported by \citey{Leahy92.1}.
 A study of a larger sample of dips is necessary. however. before any claim on the existence of periodic structures in the dip data can be settled.," A study of a larger sample of dips is necessary, however, before any claim on the existence of periodic structures in the dip data can be settled."
 Assuming that each structure in the lightcurve represents a single cloud of matter which is absorbing X-rays while crossing the line of sight. the width of the Gaussian represents a direct measure of the crossing time for such a cloud.," Assuming that each structure in the lightcurve represents a single cloud of matter which is absorbing X-rays while crossing the line of sight, the width of the Gaussian represents a direct measure of the crossing time for such a cloud."
 The horizontal extent of the cloud could be derived if the radial distance of these clouds from the neutron star and their velocity were known., The horizontal extent of the cloud could be derived if the radial distance of these clouds from the neutron star and their velocity were known.
 To give a rough estimate for the cloud size. we assume that the matter moves on Keplerian orbits close to the outer disk rim (as is favored by the model ofCrosa&Boynton 1980)). at a distance of 2104+ em from the neutron star (Chengetal. 1995)).," To give a rough estimate for the cloud size, we assume that the matter moves on Keplerian orbits close to the outer disk rim (as is favored by the model of\cite{Crosa80}) ), at a distance of $2\,10^{11}\,$ cm from the neutron star \cite{Cheng95}) )."
" The Keplerian velocity at this radius is approximately 310* cem/s. Assuming that the clouds are spherical. re.. assuming that their angular size estimated from the FWHM of the fitted Gaussian corresponds to their radial size. we derive typical cloud diameters on the order of 105:72"" cem."," The Keplerian velocity at this radius is approximately $3\,10^7$ cm/s. Assuming that the clouds are spherical, i.e., assuming that their angular size estimated from the FWHM of the fitted Gaussian corresponds to their radial size, we derive typical cloud diameters on the order of $10^{8\ldots 10}$ cm."
 Making use of the observed column densities we find a proton density of about 610!em..," Making use of the observed column densities we find a proton density of about $6\,10^{14}\,\mbox{cm}^{-3}$."
 On the other hand. the apparent periodicity mentionedabove might also indicate that the blobs are very close to the neutron star.," On the other hand, the apparent periodicity mentionedabove might also indicate that the blobs are very close to the neutron star."
 Assuming Keplerian motion. the periodicity could indicate orbital radii as small as LO” cem (the inner radius of the accretion disk is at 10* cem: Horn 1992)) and their densities would be accordingly higher.," Assuming Keplerian motion, the periodicity could indicate orbital radii as small as $10^9$ cm (the inner radius of the accretion disk is at $10^8$ cm; \cite{Horn92.1}) ) and their densities would be accordingly higher."
 In this paper we have presented the temporal evolution of spectral parameters of during a pre-eclipse dip., In this paper we have presented the temporal evolution of spectral parameters of during a pre-eclipse dip.
 Due to the large effective area of the PCA. the temporal resolution of our data is higher than that of earlier data.," Due to the large effective area of the PCA, the temporal resolution of our data is higher than that of earlier data."
 Using two different methods we were able to show that a partial covering model in which the column density η and the normalization of the unabsorbed continuum /{- are the only time variable spectral components is sufficient to explain the observed spectral and temporal variability., Using two different methods we were able to show that a partial covering model in which the column density $N_\mathrm{H}$ and the normalization of the unabsorbed continuum $I_\mathrm{U}$ are the only time variable spectral components is sufficient to explain the observed spectral and temporal variability.
 This result is 1n. qualitative agreement with the previous analyses for pre-eclipse dips presented by Ushimaruetal. (1989).. Choietal. (1994).. Leahyetal. (1994).. and Reynolds&Parmar (1995).," This result is in qualitative agreement with the previous analyses for pre-eclipse dips presented by \citey{Ushimaru89}, , \citey{Choi94.1}, , \citey{Leahy94.1}, , and \citey{Reynolds95.1}. ."
. Note that models without an unabsorbed component were not able to result in satisfactory fits. which is also consistent with the low state," Note that models without an unabsorbed component were not able to result in satisfactory fits, which is also consistent with the low state"
according to equation (8)) we may think of the free paranieter X as a direct measure of planctesimal AL. and Afaag scales roughly as 4979.,"according to equation \ref{eq:mdot}) ) we may think of the free parameter $\chi$ as a direct measure of planetesimal $\dot M$, and $M_{crit}$ scales roughly as $\chi^{0.35}$."
 Just for illustration. if (as has been done in Dodson-Robinson 2009) planctesimal random velocitics were kept at the level 87OR (rather than being esseutiallv zero as assumed in deriving equation (8))) the equation (8)) could still be used but with 4zxp?a| (DT93: ROG).," Just for illustration, if (as has been done in Dodson-Robinson 2009) planetesimal random velocities were kept at the level $\sigma\sim \Omega R_H$ (rather than being essentially zero as assumed in deriving equation \ref{eq:mdot}) )) the equation \ref{eq:mdot}) ) could still be used but with $\chi\approx p^{1/2}\ll 1$ (DT93; R06)."
 Ata =10 AU where p=10 one would then find. i; naller by a factor of 5 compared to what equation (191) predicts for \= Lat the same location., At $a=40$ AU where $p\approx 10^{-4}$ one would then find $M_{crit}$ smaller by a factor of $5$ compared to what equation \ref{eq:Mcrit}) ) predicts for $\chi=1$ at the same location.
" With equation (19)) we also show for our chosen. AL behavior that A/.,;, varies with the distance from the central object: for \y=coust the critical mass scales as xa977 in the MMSN (a= 3/2).", With equation \ref{eq:Mcrit}) ) we also show for our chosen $\dot M$ behavior that $M_{crit}$ varies with the distance from the central object: for $\chi$ =const the critical mass scales as $\approx a^{-0.53}$ in the MMSN $\alpha=3/2$ ).
 This makes Mis at 10 AU twice as large as Mj; at LO AU. everything else being equal.," This makes $M_{crit}$ at 10 AU twice as large as $M_{crit}$ at 40 AU, everything else being equal."
 Thus. taking με to be independent of the reeime of planctesimal accretion and distance from the central object 1s ecuerally not justified.," Thus, taking $M_{crit}$ to be independent of the regime of planetesimal accretion and distance from the central object is generally not justified."
 For the CT to happen before the nebular gas dispersal the erowth time of the core with mass Αν nmst be less than the lifetime of the protoplanetary nebula Τις. be. TOAit)<Tren," For the CI to happen before the nebular gas dispersal the growth time of the core with mass $M_{crit}$ must be less than the lifetime of the protoplanetary nebula $\tau_{neb}$, i.e. $\tau(M_{crit})< \tau_{neb}$."
 Plugeine in our result (19)) iuto equation (11)) this constraint cau be replirased in ters of theμέ ou the surface density of solids: SigmaGon](33) )↽⋡ This inequality represcuts the main result of this work a robust lower lit on theplenetestmeal surface density at which CT is capable of producing eiaut plaucts within a protoplanetary nebula lifetime., Plugging in our result \ref{eq:Mcrit}) ) into equation \ref{eq:timescale}) ) this constraint can be rephrased in terms of the on the surface density of solids: _s 0.1 This inequality represents the main result of this work — a robust lower limit on the surface density at which CI is capable of producing giant planets within a protoplanetary nebula lifetime.
" Note that “y;,, is a sensitive function of the nebula lifetime των. while it depeuds rather weakly on the bulk density of the core po. atmospheric dust opacity ay. aud the mean molecular weight of the atinospheric gas jv."," Note that $\Sigma_{lim}$ is a sensitive function of the nebula lifetime $\tau_{neb}$, while it depends rather weakly on the bulk density of the core $\rho_c$, atmospheric dust opacity $\kappa_0$, and the mean molecular weight of the atmospheric gas $\mu$."
" Now we cau also determine the lamiting core mass η defined as the critical core mass for 4X,=Xii. lo. at the very extreme of the region where CI cau still occur within Τιμ."," Now we can also determine the limiting core mass $M_{lim}$ defined as the critical core mass for $\chi\Sigma_s=\Sigma_{lim}$, i.e. at the very extreme of the region where CI can still occur within $\tau_{neb}$."
" It is fouud by substituting (23)) iuto (19)):501(2 By construction Mi;,, is indepeudent of the surface deusity profile in the nebula.", It is found by substituting \ref{eq:Sig_lim}) ) into \ref{eq:Mcrit}) ): 15 By construction $M_{lim}$ is independent of the surface density profile in the nebula.
 This mass is to be compared with the isolation mass Ανω (the core mass at which it has accreted all solid mass within its feeding zone) an aunuulus ceutered ou cores orbit and having a full width equal to CRyy: (an MMSN-like deusitv profile was used iu mniakiug uunerical estimate)., This mass is to be compared with the isolation mass $M_{iso}$ (the core mass at which it has accreted all solid mass within its feeding zone) — an annulus centered on core's orbit and having a full width equal to $\xi R_H$: (an MMSN-like density profile was used in making numerical estimate).
" At LL AU one finds Mj,26 Mj. which is smaller than Mj, ]"," At $44$ AU one finds $M_{iso}\approx 6$ $_\oplus$, which is smaller than $M_{lim}$."
ILowever. a modest radial displacement of the core due to some type of mieration can casily expose it to additional fresh material allowing AI. to reach. ο (Alibert 2005).," However, a modest radial displacement of the core due to some type of migration can easily expose it to additional fresh material allowing $M_c$ to reach $M_{lim}$ (Alibert 2005)."
" Alternatively. increasing X, (boosting up X4) bv a factor of 2 makes MigzMii at HE AU."," Alternatively, increasing $\Sigma_s$ (boosting up $\Sigma_1$ ) by a factor of 2 makes $M_{iso}\approx M_{lim}$ at 44 AU."
" Quite interestingly. the value of Aj, is not too far from LO M. conuuoulv accepted as the core nass throughout the whole protoplauctary disk."," Quite interestingly, the value of $M_{lim}$ is not too far from $10$ $_\oplus$ commonly accepted as the core mass throughout the whole protoplanetary disk."
 This coincidence 18 accidental since our estimate of τε was derived selt-cousisteutlv rather than postulated iu an ad hoc fashion., This coincidence is accidental since our estimate of $M_{lim}$ was derived self-consistently rather than postulated in an ad hoc fashion.
" Given a coustraiut (23)) and having a particular model of the radial distribution. of X, one cau determine the inaxiual radial extent ο of the region in protoplauctary disk in which the CT can produce eiaut planets within the nebula life time Τμ", Given a constraint \ref{eq:Sig_lim}) ) and having a particular model of the radial distribution of $\Sigma_s$ one can determine the maximal radial extent $a_{lim}$ of the region in protoplanetary disk in which the CI can produce giant planets within the nebula life time $\tau_{neb}$.
" Since X, is expected to be a decreasing function of α this would only be possible for ακ αμ ", Since $\Sigma_s$ is expected to be a decreasing function of $a$ this would only be possible for $a<a_{lim}$ .
"Usiug our power-law paraluctrization (5)) of X, we fiud that in à AIAISN-like disk with a=3/2 LL", Using our power-law parametrization \ref{eq:Sig_s}) ) of $\Sigma_s$ we find that in a MMSN-like disk with $\alpha=3/2$ 44.
" Accretion at AZ=M, corresponds to y=1 in this formmla.", Accretion at $\dot M=\dot M_{max}$ corresponds to $\chi=1$ in this formula.
" The limiting distance fouud iu equation (30)) Is shnudlar to previous estimates (6.8. Dodsou-Robinsou 2009) obtained for less vigorous planctesimal accretion (i.e. for sualler M) aud fixed A,;;=10 AD but this is just a coincidence.", The limiting distance found in equation \ref{eq:a_lim}) ) is similar to previous estimates (e.g. Dodson-Robinson 2009) obtained for less vigorous planetesimal accretion (i.e. for smaller $\dot M$ ) and fixed $\dot M_{crit}=10$ $_\oplus$ but this is just a coincidence.
" Equation (30)) shows that as X, (or X4) ducreases. tle extent of CL-capable part of the protoplanctary disk also erows."," Equation \ref{eq:a_lim}) ) shows that as $\Sigma_s$ (or $\Sigma_1$ ) increases, the extent of CI-capable part of the protoplanetary disk also grows."
" But X, cannot be increased without mit at sole point the gascous colmpouent of the protoplanetary nebula would become self-eraxitatiug. aud. depending ou its cooling time (CGamuiue 2000). would either fragment or evolve quasi-viscously while maintaining the mareiuallyeravitationally unstable state (Bafkov 2009: Clarke 2009)."," But $\Sigma_s$ cannot be increased without limit — at some point the gaseous component of the protoplanetary nebula would become self-gravitating, and, depending on its cooling time (Gammie 2000), would either fragment or evolve quasi-viscously while maintaining the marginallygravitationally unstable state (Rafikov 2009; Clarke 2009)."
 Thus the inost extreme value of απ can be, Thus the most extreme value of $a_{lim}$ can be
 , 
We asstune two SN survey models: a shallow one (S20k) that covers 20.000 dee? to 2=(8 for 10 years. aud a deep one (D2k) that covers2000 dee? to 2=12 for 5 vears.,"We assume two SN survey models: a shallow one (S20k) that covers 20,000 $^2$ to $z = 0.8$ for 10 years, and a deep one (D2k) that covers2000 $^2$ to $z = 1.2$ for 5 years."
 The S20k data may be extracted as one of the may products frou the proposed LSST. and a dedicated ground-based SN survey will likely be sufficient to produce the D2k data.," The S20k data may be extracted as one of the many products from the proposed LSST, and a dedicated ground-based SN survey will likely be sufficient to produce the D2k data."
 Since there is no precedent for such SN BAO surveys. we have to make crude estimates for the fiducial survey parameters and vary them to cover a wider rauge.," Since there is no precedent for such SN BAO surveys, we have to make crude estimates for the fiducial survey parameters and vary them to cover a wider range."
 We calculate the observer-framic SN rate based on the rest-frame SN rates (Cappellareetal.1999:Tardin 2006).," We calculate the observer-frame SN rate based on the rest-frame SN rates \citep{cappellaro99,hardin00,pain02,madgwick03,tonry03,blanc04,
dahlen04,neill06}."
. The resulting wmmber of SNe per steradiau per nuit redshift per vear (observer frame) is roughly where the first terii on the r.lis., The resulting number of SNe per steradian per unit redshift per year (observer frame) is roughly where the first term on the r.h.s.
" fits the observed SN vate af i<0.55, ancl the additional exponeutial termwith ο=0.125 (0.21) cuts off the distribution at z0.8 (1.2) for S20k (D2k)."," fits the observed SN rate at $z \le 0.55$, and the additional exponential termwith $a = 0.125$ $0.21$ ) cuts off the distribution at $z \sim 0.8$ $1.2$ ) for S20k (D2k)."
 We take the efficiency with which S20k (D2k) will produce well-observed SN light-curves suitable for accurate phioto-zss from this distributionas (10043)., We take the efficiency with which S20k (D2k) will produce well-observed SN light-curves suitable for accurate s from this distributionas ).
" The accuumated SN surface deusitv is then X=370 and 980 7 for S20x (7. Linillion SNe in 10 vears) aud D2k (2 million SNe im 5 vears}. respectively,"," The accumulated SN surface density is then $\Sigma = 370$ and 980 $^{-2}$ for S20k (7.4 million SNe in 10 years) and D2k (2 million SNe in 5 years), respectively."
 We adopt a pedagogicalOa convention and model the eror as a Gaussian with rum σ.=o.y(l|2) and zero ὃν= 0., We adopt a pedagogical convention and model the error as a Gaussian with rms $\sigma_z=\sigma_{z{\rm0}}(1+z)$ and zero $\delta z = 0$ .
" We assien for S20k 6,4=0.02 and for D2k o9=0.01. which are achievable with simple algorithms (Pintoctal.200L:Wang2007).."," We assign for S20k $\sigma_{z{\rm 0}}=0.02$ and for D2k $\sigma_{z{\rm 0}}=0.01$, which are achievable with simple algorithms \citep{pinto04,wang07b}."
 Table Po stuumarizes the SN surveys., Table \ref{tab:model} summarizes the SN surveys.
 It tuchides two additional quantities: the SN clusterimg bias 5. aud the cut-off wavenuuber Aga for the BAO analysis. which is set to reduce the impact of nonlinear erowth (see Eisensteinetal.2007. for recovering hieh-A information from spectroscopic surveys}.," It includes two additional quantities: the SN clustering bias $b$ , and the cut-off wavenumber $k_{\rm max}$ for the BAO analysis, which is set to reduce the impact of nonlinear growth (see \citealt{eisenstein06}
 for recovering $k$ information from spectroscopic surveys)."
 We use a modified forecast tool (Zhan2006) to assess the cosmological constraints from the clustering of SNe., We use a modified forecast tool \citep{zhan06d} to assess the cosmological constraints from the clustering of SNe.
 Iu sumunary. the SN power spectrum in the th redshift bin (Κε) reads Eiseustein2003) hey) ο URAL) y) | sj. where the subscript f denotes quantities in a reference cosimological model (the same as the fiducial model in the forecast). D is the aneular diameter distance. Z7 is the IIubble parameter. jis the redshift distortion paramicter. =hyk. G ds the linear growth factor. P(k} is the latter power spectrum at 2.=O0. aud s=nlods the shot noise.," In summary, the SN power spectrum in the $i$ th redshift bin $P_i(\bit{k}_{\rm f})$ reads \citep{seo03}
 ) )^2 b_i^2 G_i^2 (k) ] + s_i, where the subscript f denotes quantities in a reference cosmological model (the same as the fiducial model in the forecast), $D$ is the angular diameter distance, $H$ is the Hubble parameter, $\beta$ is the redshift distortion parameter, $\mu = k_\parallel/k$, $G$ is the linear growth factor, $\mathcal{P}(k)$ is the matter power spectrum at $z = 0$, and $s=n^{-1}$ is the shot noise."
 The true saveuiuubers Ayaud Ay are related to the references by Ay=hepDgfD aud hy=T/T., The true wavenumbers $k_\perp$and $k_\parallel$ are related to the references by $k_\perp = k_{\rm f\perp} D_{{\rm f}} / D$ and $k_\parallel = k_{\rm f\parallel} H / H_{{\rm f}}$.
 Note that for a survey with o.z0.0112). the radial BAO information is esseutiallv lost. c.g.. the power spectrin (excluding the shot noise) at the fiudamental mode of BAOs ~27/150XIpe is suppressed by a factor of SU or moreky at +=0.6. though a rius of O.O03€1|2) would recover thatinforxiuatiou.," Note that for a survey with $\sigma_z \gtrsim 0.01(1+z)$, the radial BAO information is essentially lost, e.g., the power spectrum (excluding the shot noise) at the fundamental mode of BAOs $k_\parallel\sim 2\pi/150\,\mbox{Mpc}^{-1}$ is suppressed by a factor of 80 or more at $z = 0.6$, though a rms of $0.003(1+z)$ would recover thatinformation."
" uses the Fisher matrix to estimate the error bounds of the parameters (TeemarkL997): — where Vy is the comoving survey volume and [4,1 is the parameter set. which includes );. σι, ὅτι and s; of cach bin aud 8 cosmological parameters: Wy. Wy. ey, (the matter deusitv). wy, (the barvon deusitv). 0, (the angular size of the sound horizon at the last scattering surface). Ok (the curvature parameter). às, (the scalar spectral index). and AT (the normalization of the prunordial curvature power spectrum)."," uses the Fisher matrix to estimate the error bounds of the parameters \citep{tegmark97b}: = where $V_{\rm f}$ is the comoving survey volume and $\{q_\alpha\}$ is the parameter set, which includes $b_i$, $\sigma_{z,i}$, $\delta z_i$ , and $s_i$ of each bin and 8 cosmological parameters: $w_0$ , $w_{\rm a}$, $\omega_{\rm m}$ (the matter density), $\omega_{\rm b}$ (the baryon density), $\theta_{\rm s}$ (the angular size of the sound horizon at the last scattering surface), $\Omega_{\rm k}$ (the curvature parameter), $n_{\rm s}$ (the scalar spectral index), and $\Delta_{\rm R}^2$ (the normalization of the primordial curvature power spectrum)."
" The minima mareinalized error of q, is σι)=(F 11/2 "," The minimum marginalized error of $q_\alpha$ is $\sigma(q_\alpha) =
(F^{-1})_{\alpha\alpha}^{1/2}$ ."
"ludepeudeut Fisher matrices are additive: a prior on μι. 9p(q,). can be introduced via BE’=Fi,|Ty)."," Independent Fisher matrices are additive; a prior on $q_\alpha$, $\sigma_{\rm P}(q_\alpha)$ can be introduced via $F_{\alpha\alpha}^{\rm new} = F_{\alpha\alpha}
+\sigma_{\rm P}^{-2}(q_\alpha)$."
" The fiducial values of hj. σε, and s; lore listed in Table 1.. ὅτι= 0. and (ye Wye eine why Oo. Oi ns AR)=O 1.0. 0.127. 0.0223. 0.5967. 0. 0.951. 2.0«10 9?) from the 3-year results (Spereeletal.2007)."," The fiducial values of $b_i$, $\sigma_{z,i}$, and $s_i^{-1}$ are listed in Table \ref{tab:model}, $\delta z_i = 0$ , and $w_0$, $w_{\rm a}$, $\omega_{\rm m}$, $\omega_{\rm b}$, $\theta_{\rm s}$, $\Omega_{\rm k}$, $n_{\rm s}$, $\Delta_{\rm R}^2) = (-1$, 0, 0.127, 0.0223, $^\circ$, 0, 0.951, $2.0\times 10^{-9}$ ) from the 3-year results \citep{spergel07}."
". Unless stated otherwise. we always include fairly weak priors op(lub;)= 0.3. op(lnAR)= 0.2. op(lu0.)=cp(n,)0.05. aud συ(ὅτ)92U?op(o,= 0250, "," Unless stated otherwise, we always include fairly weak priors $\sigma_{\rm P}(\ln b_i) = 0.3$ , $\sigma_{\rm P}(\ln \Delta_{\rm R}^2) = 0.2$ , $\sigma_{\rm P}(\ln \theta_{\rm s}) = 
\sigma_{\rm P}(n_{\rm s}) = 0.05$, and $\sigma_{\rm P}(\delta z_i) = 2^{-1/2} \sigma_{\rm P}(\sigma_{z,i}) =
0.25\sigma_{z,i}$ ."
To show the complementarity;) between the SN BAO and Di techniques. we include the Dp coustraimts at the end of Section L..," To show the complementarity between the SN BAO and $D_{\rm L}$ techniques, we include the $D_{\rm L}$ constraints at the end of Section \ref{sec:cons}. ."
" The Fisher matrix for SN Dy is fau)ety no DLU) |AL | eye | ο. ο where σιlog0.15 is the scatter of the ον apparent magnitude. the subscript psignifies space.AL is the SN absolute maeguitucde. ο and οὐ account forpossibleSN evolution. aud p(z[:,) is the probability"," The Fisher matrix for SN $D_{\rm L}$ is = ) = (z) +M + e_1z + e_2] p(z | ) dz, where $\sigma_{\rm m} = 0.15$ is the scatter of the SN apparent magnitude, the subscript psignifies space,$M$ is the SN absolute magnitude, $e_1$ and $e_2$ account forpossibleSN evolution, and $p(z|z_{\rm p})$ is the probability"
clusters also contributes to the growth of a stellar cluster.,clusters also contributes to the growth of a stellar cluster.
 shows an example of this process whereby accretion of gas FigureB]and stars occurs along filaments flowing into the cluster., Figure \ref{OAclusform} shows an example of this process whereby accretion of gas and stars occurs along filaments flowing into the cluster.
 An important result from this work is that the clustering depends strongly on the local gravitational binding of the gas prior to star formation., An important result from this work is that the clustering depends strongly on the local gravitational binding of the gas prior to star formation.
" Figure H] shows the resulting stellar densities as a function of how bound the cloud was initially in terms of the critical mass per unit length, M/L, to be bound."," Figure \ref{clusvsMoL} shows the resulting stellar densities as a function of how bound the cloud was initially in terms of the critical mass per unit length, $M/L$, to be bound."
 Two measures of the stellar densities are plotted., Two measures of the stellar densities are plotted.
 The blue crosses show the median stellar density determined by the volume needed to contain the ten nearest stellar neighbours., The blue crosses show the median stellar density determined by the volume needed to contain the ten nearest stellar neighbours.
 The black triangles show the density of stars contained in a fixed volume of size 0.5pc., The black triangles show the density of stars contained in a fixed volume of size 0.5pc.
 The density determined by the first method is significantly higher as it typically is based on much smaller volumes., The density determined by the first method is significantly higher as it typically is based on much smaller volumes.
" In both cases, the stellar density is low in regions that were initially unbound and is much higher in the bound parts of the cloud with M/L>1."," In both cases, the stellar density is low in regions that were initially unbound and is much higher in the bound parts of the cloud with $M/L>1$."
 This result is understandable in that clusters are (at least temporarally) bound objects and their formation requires that the pre-star formation gas is also bound., This result is understandable in that clusters are (at least temporarally) bound objects and their formation requires that the pre-star formation gas is also bound.
" In locally unbound regions, it is still possible to form small stellar clusters in regions where turbulent compression and shocks result in a locally bound region."," In locally unbound regions, it is still possible to form small stellar clusters in regions where turbulent compression and shocks result in a locally bound region."
 This process is more efficient in regions that are globally bound., This process is more efficient in regions that are globally bound.
 Larger scale regions containing many subsystems are bound even before any turbulent support is dissipated., Larger scale regions containing many subsystems are bound even before any turbulent support is dissipated.
 The smaller systems, The smaller systems
explaining the systematic effect seeu in Figure &..,explaining the systematic effect seen in Figure \ref{fig:gam_diff}.
 The departures of the VCS mioclels Irou the WEPC2 data appear to be c«npletely. consistent their residuals from the original ACS/WEC data as well., The departures of the VCS models from the WFPC2 data appear to be completely consistent their residuals from the original ACS/WFC data as well.
 The patterns produced by assuming the inuer cusp to have coustaut slope are reacliy evident iu the residual traces presented in Figure 103 ol Ferrareseetal.(2006)., The patterns produced by assuming the inner cusp to have constant slope are readily evident in the residual traces presented in Figure 103 of \citet{lf}.
. Because that figure ouly presents residuals in the observed domain. the systelmatic trends are more subtle than tlose evident in Figure 11: however. one always sees a sharp bump at r=ry. with a nonzero reskual slope at sinaller raclii.," Because that figure only presents residuals in the observed domain, the systematic trends are more subtle than those evident in Figure \ref{fig:acscomp1}; however, one always sees a sharp bump at $r=r_b,$ with a nonzero residual slope at smaller radii."
 I1 the cases of NGC. [173. LITS. arr LISGB. Ferrareseοἱal.(2006). fited models to £alaxies tha we classified as core systems.," In the cases of NGC 4473, 4478, and 4486B, \cite{lf} fitted models to galaxies that we classified as core systems."
 GC LIT3 ts well described by à Niker law with a sinall a parameter., NGC 4473 is well described by a Nuker law with a small $\alpha$ parameter.
 The model in contras clearly over estimates the central 5': again the pattern of resicluals hat we see here is obvious in VCS Figure 103 as well., The model in contrast clearly over estimates the central $\gamma';$ again the pattern of residuals that we see here is obvious in VCS Figure 103 as well.
 Tlits is also rue for NGC LISGB. The extremely large residuals from the fit for NGC [178 seen i1 Figure LL are again echoed in VC5 Figure 103., This is also true for NGC 4486B. The extremely large residuals from the fit for NGC 4478 seen in Figure \ref{fig:acscomp1} are again echoed in VCS Figure 103.
 While we concluded t1at NGC E178 bas a snall nileus (Laueretal.2005).. no πα‘omponent was included in the VCS moclel fits.," While we concluded that NGC 4478 has a small nucleus \citep{l05}, no such component was included in the VCS model fits."
 Lastly. one low-luninosity core galaxy. NGC εως. was [itec with a highly IunOlls àid spatially extended nucleis by Ferrareseetal.(2006):: it will be cliscussecl in the context o “other galaxies with large assumed VCS uuclei below.," Lastly, one low-luminosity core galaxy, NGC 4458, was fitted with a highly luminous and spatially extended nucleus by \citet{lf}; it will be discussed in the context of other galaxies with large assumed VCS nuclei below."
 We couclude that uoue of the presen core galaxies observed in he VCS are well matched by the VC'S models as the resolution liuit is approached., We conclude that none of the present core galaxies observed in the VCS are well matched by the VCS models as the resolution limit is approached.
 The VCS choice of a cusp with a coustaut inuer slope or fitting a mocle to galaxies with small cores always results in higher +! valttes than result (roin the Nuker parameers., The VCS choice of a cusp with a constant inner slope or fitting a model to galaxies with small cores always results in higher $\gamma'$ values than result from the Nuker parameters.
 NGC οί. H61.. 1621. and. L660 are galaxies that we have classified as power-laws that Ferrareseetal.(2006) fitted. with pu'e laws.," NGC 4434, 4464, 4621, and 4660 are galaxies that we have classified as power-laws that \citet{lf} fitted with pure laws."
 The VCS inodels for a| four galaxies [all below the data as r—0 (Figure 12 2))., The VCS models for all four galaxies fall below the data as $r\rightarrow0$ (Figure \ref{fig:acscomp2}) ).
 Ágain. this behavior is echoed in VCS Figure 103. which shows that the fits [all beloyv the ACS VCS data for r<073 in NGC LIGL 1621. ancl 166Q. and r«OL:5 iu NGC 1131.," Again, this behavior is echoed in VCS Figure 103, which shows that the fits fall below the ACS VCS data for $r<0\asec3$ in NGC 4464, 4621, and 4660, and $r<0\asec15$ in NGC 4434."
 Ferrareseetal.(2006) note that the ACS F172W images of boh NGC 1621 αιd 1660 are actually centrally saturated., \citet{lf} note that the ACS F475W images of both NGC 4621 and 4660 are actually centrally saturated.
 This is problematic as intormation about the central sructwe is lost once saturatiou oceurs. but it still must be accounted for in the profile moclels. given that lielit from the center is also conveyed to larger radii by the PSF.," This is problematic as information about the central structure is lost once saturation occurs, but it still must be accounted for in the profile models, given that light from the center is also conveyed to larger radii by the PSF."
 Iu contrast the Nuker models nicely uateli all four galaxies. In the case of NGC [161 it is especially uotable that the Nuker profile. whic1 was derived from WEDPCI imagery. bas better fidelity to the ACS prolile than the VCS mocel de‘ived directly from the same data.," In contrast the Nuker models nicely match all four galaxies, In the case of NGC 4464 it is especially notable that the Nuker profile, which was derived from WFPC1 imagery, has better fidelity to the ACS profile than the VCS model derived directly from the same data."
 The plot of Nuker ο) VELsts VCS ο in Figure 8 contains two more power-law galaxies. NGC L117 aud 1561. for which we do not have our own WEPC? or WEPCT profiles. but again the residuals in VCS Figure 103 show that their models fall well below the ACS data as --0. This appears to be a generic feature of nearly all the VCS galaxies modelled as non-nucleated," The plot of Nuker $\gamma'$ versus VCS $\gamma'$ in Figure \ref{fig:gam_diff}
 contains two more power-law galaxies, NGC 4417 and 4564, for which we do not have our own WFPC2 or WFPC1 profiles, but again the residuals in VCS Figure 103 show that their models fall well below the ACS data as $r\rightarrow0.$ This appears to be a generic feature of nearly all the VCS galaxies modelled as non-nucleated"
most stable members of its class.,most stable members of its class.
 This characteristic allows us to sum up all the oobservations taken with the same instrumental settings., This characteristic allows us to sum up all the observations taken with the same instrumental settings.
 The cumulative PN spectrum of ccannot be adequately fit by a power-law plus blackbody model., The cumulative PN spectrum of cannot be adequately fit by a power-law plus blackbody model.
 This ts the first magnetar for which such a conclusion can be derived based only on the X-ray data. thanks to the very low interstellar absorption.," This is the first magnetar for which such a conclusion can be derived based only on the X–ray data, thanks to the very low interstellar absorption."
 A good fit is instead obtained with a double blackbody model., A good fit is instead obtained with a double blackbody model.
 The known distance of the SMC (60kpe.?) allows a precise measure of the size of the emitting region of the two blackbodies.," The known distance of the SMC \citep[60 kpc,][]{keller06} allows a precise measure of the size of the emitting region of the two blackbodies."
 The ~2 km radius of the region associated to the hotter component is compatible with a hot spot on the neutron star surface., The $\sim$ 2 km radius of the region associated to the hotter component is compatible with a hot spot on the neutron star surface.
 The radius of the emitting region of the cooler blackbody. km. 15 consistent with a large fraction of the magnetar 12.171surface.," The radius of the emitting region of the cooler blackbody, $^{+2.1}_{-1.4}$ km, is consistent with a large fraction of the magnetar surface."
 However. the strong pulsation below | keV (see Fig. 3)). 1)). (?) (3)," However, the strong pulsation below 1 keV (see Fig. \ref{profile}) \ref{spec})"
 (?).. (see.," \citep{halpern05} \citep{israel08}
 \citep{tmt05}."
 (?).. (22)... (?).. (?).. (?).. (1)," \citep[see,
e.g.,][]{kuiper06,gotz06,leyder08} \citep{nakagawa08}. \citep{lyutikov06,fernandez07}. \citep{RCSnanda}. \citep{juett02}, \citep{tmt05},"
 ? 2.. (see.e.g..?)..," \citep{woods04} \cite{durant06} \ref{lines}, \citep[see, e.g.,][]{ho03},"
the reference stars are insignificantly variable themselves.,the reference stars are insignificantly variable themselves.
 We therefore adopted the ollowiug procedure to test for variability in the reference stars., We therefore adopted the following procedure to test for variability in the reference stars.
 A refereuce stars s. is chosen. and the average flux of all other reference stars ormed.," A reference star, $s$, is chosen, and the average flux of all other reference stars formed."
" The relative magnitude of s is hen evaluated with respect to this average,", The relative magnitude of $s$ is then evaluated with respect to this average.
 This is done at every epoch. producing a fune series Hy[etydeenaMlapendee Toni which the mean is subtracted so that M.Hg=0.," This is done at every epoch, producing a time series $m_{s,1},m_{s,2},\ldots,m_{s,t},\ldots$, from which the mean is subtracted so that $\sum_t m_{s,t} = 0$."
 At cach epoch the tota uncertaintv iu he relative niaguitude. εκ. 18. also estimated. (see below).," At each epoch the total uncertainty in the relative magnitude, $\epsilon_{s,t}$, is also estimated (see below)."
 WeN then use a \>5 test o decide whether the scatter of these relative maguitudes is consistent with their measurement uucertaiufies. using he statistic The nul hypothesis for this test is that there is no variability iu the time series.," We then use a $\chi^2$ test to decide whether the scatter of these relative magnitudes is consistent with their measurement uncertainties, using the statistic The null hypothesis for this test is that there is no variability in the time series."
 If P(42) is the probability hat the null liypothesis is true. then the larger 42 the stnaller Pi\2).," If ${\rm P}(\chi^2_s)$ is the probability that the null hypothesis is true, then the larger $\chi^2_s$ the smaller ${\rm
P}(\chi^2_s)$."
 A star will fail the test (show evidence or variability) if P(42)«0.01. a confidence interval. (," A star will fail the test (show evidence for variability) if ${\rm P}(\chi^2_s)<0.01$, a confidence interval. ("
This would correspoud to 2.50 significance if he distribution were Gaussian.),This would correspond to $\sigma$ significance if the distribution were Gaussian.)
 (2 is micasured for every reference star. aud if auv fails the test. the star with the lighest 4? is rejected from the set of reference stars.," $\chi^2_s$ is measured for every reference star, and if any fails the test, the star with the highest $\chi^2$ is rejected from the set of reference stars."
 This whole process (starting from formine the time series for cach star) is iterated until all remaining reference stars oes the \? test., This whole process (starting from forming the time series for each star) is iterated until all remaining reference stars pass the $\chi^2$ test.
 These then define the reference systeni roni which to determine the relative maeuitude time series of the science star., These then define the reference system from which to determine the relative magnitude time series of the science star.
 The point of using more than oue reference star to eternune the time series for the scicuce star is that it increases the photometric precision of the relative magnitudes., The point of using more than one reference star to determine the time series for the science star is that it increases the photometric precision of the relative magnitudes.
 Furthermore. auv small residual variability In any one reference star will have a smaller effect on the verage.," Furthermore, any small residual variability in any one reference star will have a smaller effect on the average."
 The X test is a good initial test for variability. as --- js sensitive to any type of variability aud not just periodic variations.," The $\chi^2$ test is a good initial test for variability, as it is sensitive to any type of variability and not just periodic variations."
 However. for the test to be successful. the errors. εκ. nmst be accurately deteruuned.," However, for the test to be successful, the errors, $\epsilon_{s,t}$, must be accurately determined."
 Usine photon statistics we can evaluate the “formal” errors in auv magnitude determination., Using photon statistics we can evaluate the “formal” errors in any magnitude determination.
 These are (1) photou arrival noise from the object. (2) photon arrival noise from). the sky. and (3) imperfect determination of the subtracted sla level.," These are (1) photon arrival noise from the object, (2) photon arrival noise from the sky, and (3) imperfect determination of the subtracted sky level."
" There are. however. additional ""óufonuual errors which need to be accounted for if we are not to overestimate: DX7."," There are, however, additional “informal” errors which need to be accounted for if we are not to overestimate $\chi^2$."
 Thed most iuportaut- of these results frou muporfect flat fielding., The most important of these results from imperfect flat fielding.
 We have estimated. (from the construction o| different flats) the accuracv of the flat to be (0.54. aud have combined this in quadrature with the forma error sources to give €4.," We have estimated (from the construction of different flats) the accuracy of the flat to be $0.5\%$, and have combined this in quadrature with the formal error sources to give $\epsilon_{s,t}$."
 Other informal eror sources include (1) imperfect ceutering of the aperture aud (5) PSF varlatious ovor the frame., Other informal error sources include (4) imperfect centering of the aperture and (5) PSF variations over the frame.
 These were πίσσα wousing a sufBcieutlv large aperture., These were minimised by using a sufficiently large aperture.
 To maxiuise the sienalto-noise ratio 1u aperture photometry. one usually chooses au aperture of size approximately equal to the ΠΛΗΝ (haltawidth at ιαπακια) of the PSF (ITowell 1989)).," To maximise the signal-to-noise ratio in aperture photometry, one usually chooses an aperture of size approximately equal to the HWHM (half-width at half-maximum) of the PSF (Howell \cite{howell89}) )."
 It is not necessary to get all of the light | or apply an aperture correction with relative photometry. provided that in a eiven frame a sinele-sized aperture is used and all stars have the same PSF.," It is not necessary to get all of the light – or apply an aperture correction – with relative photometry, provided that in a given frame a single-sized aperture is used and all stars have the same PSF."
 When there are PSF variations (as was the case for our data). a larger aperture must be used.," When there are PSF variations (as was the case for our data), a larger aperture must be used."
 This reduces the contribution from the informal errors 55) at the expense of iucreasing the formal errors (13)., This reduces the contribution from the informal errors 5) at the expense of increasing the formal errors (1–3).
 Thus the threshold for variability detections n auv science star will be at a higher magnitude limit but more reliable — than when using a simaller aperture., Thus the threshold for variability detections in any science star will be at a higher magnitude limit – but more reliable – than when using a smaller aperture.
 We found that au aperture size of about 1&IIWITME reduced the informal errors to well below the 0.5( level., We found that an aperture size of about $4\times$ HWHM reduced the informal errors to well below the $0.5\%$ level.
 Yot other sources of error (such as imperfect clharec transfer eficicney aud ronlinearity of the CCD flux seusitivitv) are also helow lis level., Yet other sources of error (such as imperfect charge transfer efficiency and nonlinearity of the CCD flux sensitivity) are also below this level.
 After establishing reliable reference stars. the 4? test was applied to the science star.," After establishing reliable reference stars, the $\chi^2$ test was applied to the science star."
 Tf it failed his test (eave evidence for variability). the next stage was to determine whether there was evidence for this variability cine periodic.," If it failed this test (gave evidence for variability), the next stage was to determine whether there was evidence for this variability being periodic."
 For this purpose we evaluated the described by Scargle (1982))., For this purpose we evaluated the described by Scargle \cite{scargle82}) ).
 Iu contrast o the classical periodograim. if can be shown that the nodified periodogram ds equivalent to a least-squares fit (iu the time domain) of sinusoids to the data (Lomb 1976)).," In contrast to the classical periodogram, it can be shown that the modified periodogram is equivalent to a least-squares fit (in the time domain) of sinusoids to the data (Lomb \cite{lomb76}) )."
 Furthermore. the method allows a simple determination of the significance of any peak iu the xriodograim based ou the known noise in the We consider a peak to he siguificaut only at the leve wwheu the probability. p. that a peal is due to noise is < o0.01).," Furthermore, the method allows a simple determination of the significance of any peak in the periodogram based on the known noise in the We consider a peak to be significant only at the level when the probability, $p$, that a peak is due to noise is $< 0.01$ )."
 If a significant peak is found. the time series is phased o the detected period aud examined to see if it shows evelie variability.," If a significant peak is found, the time series is phased to the detected period and examined to see if it shows cyclic variability."
 We now have four criteria which the data wast satisty before a science star is flageccl as showing evidence for periodic variations:, We now have four criteria which the data must satisfy before a science star is flagged as showing evidence for periodic variations:
adopting a constant [a/Fe] over the entire range of metallicity.,adopting a constant $\rm [{\alpha}/Fe]$ over the entire range of metallicity.
 The computational routine requires as input parameters: Fie., The computational routine requires as input parameters: Fig.
 9. shows the flow-diagram of the computational routine., \ref{diagram} shows the flow-diagram of the computational routine.
 In the following. references to eqs.," In the following, references to eqs."
 1--3 either equations in Sect., -8 either equations in Sect.
 3.1 or in Sect. 3.2..," \ref{diskeq} or in Sect. \ref{bulge},"
 depending on the adopted enrichment scenario., depending on the adopted enrichment scenario.
(n). nj. c Kao Ko J1650—500 J1650—500 J1650—500 Ka ,$\dot{m}$ $\dot{m}$ $c$ $\alpha$ $\alpha$ $-$ $-$ $-$ $\alpha$ 
Or Nas=512.,for $N_{\rm side}=512$.
 The second term on the simulates he noise contribution on each pixel with Caussian random number 5gNCO.1).," The second term on the simulates the noise contribution on each pixel with Gaussian random number $\eta \sim N(0,1)$."
 In this work. we perform both a Gaussian [requoentist est ancl fyp--estimations.," In this work, we perform both a Gaussian frequentist test and -estimations."
 In the former. the Gaussian simulations are processed in the same wav as Ίσα.," In the former, the Gaussian simulations are processed in the same way as Eq."
 13. but ree of the term., \ref{eq_fnl_map} but free of the term.
 In this section. we introduce the data processing methods applied to both the observed and the sipulated realisations for studies of the skeleton leneth distribution.," In this section, we introduce the data processing methods applied to both the observed and the simulated realisations for studies of the skeleton length distribution."
 The processing steps presented here follow the strategy detailed in Section 4 of Eriksenetal.(2004)., The processing steps presented here follow the strategy detailed in Section 4 of \citet{Eriksen_etal_2004}.
. The base-mask is applied to the map to avoid. Galactic foreground contamination., The base-mask is applied to the map to avoid Galactic foreground contamination.
 Following the methodology. of ]iriksenetal.(2004).. we do not exclude point sources. in particular because any additional smoothing applied to the mask reduces the sky. coverage available for analysis clramatically.," Following the methodology of \citet{Eriksen_etal_2004}, we do not exclude point sources, in particular because any additional smoothing applied to the mask reduces the sky coverage available for analysis dramatically."
 This approach is supported by studies of the spectral parameter. 5. by Eriksenetal.(2004).. which indicates that smoothing of the data renders the skeleton less sensitive to point source signal contributions for larger ΕΝΝΕ.," This approach is supported by studies of the spectral parameter, $\gamma$, by \citet{Eriksen_etal_2004}, which indicates that smoothing of the data renders the skeleton less sensitive to point source signal contributions for larger FWHMs."
 Moreover. a mecian-filter technique is applied to the point sources to investigate their impact on the skeleton statistics for smaller smoothing ΗΕ.," Moreover, a median-filter technique is applied to the point sources to investigate their impact on the skeleton statistics for smaller smoothing FWHMs."
 pecifically. for a given pixel 7 that would be eliminated by the point-source mask. we consider all other unmaskecl pixels within a 1 radius and determine the median temperature for this set of pixels.," Specifically, for a given pixel $i$ that would be eliminated by the point-source mask, we consider all other unmasked pixels within a $1^\circ$ radius and determine the median temperature for this set of pixels."
 The temperature at. pixel 7 is then replaced. by this median. value. and the process repeated. for all. pixels specified in the point-source mask.," The temperature at pixel $i$ is then replaced by this median value, and the process repeated for all pixels specified in the point-source mask."
 The median-filtered map is then analysed in the same manner as the undiltered data sel., The median-filtered map is then analysed in the same manner as the unfiltered data set.
 Following standard procedure in CIB data analysis. the monopole and dipole components are fitted and removed from cach map outside the masked. region.," Following standard procedure in CMB data analysis, the monopole and dipole components are fitted and removed from each map outside the masked region."
 This step is achieved using the LIEALPix E90 subroutine.dipole.," This step is achieved using the HEALPix F90 subroutine,."
 The resulting map is then smoothed with a Gaussian beam in harmonic space. again using ΤΗΛΟν tools.," The resulting map is then smoothed with a Gaussian beam in harmonic space, again using HEALPix tools."
 For easy comparison with the previous analysis of Eriksenetal.(2004).. the ENIINE widths. @pyHar. selected. for. these Gaussian smoothingbeams are taken to be 0753. 0764. 0285. 1728. 1770. 2713. 2755. 2798 and 3:40 (we abandon the larger angular scales used in the former analysis to ensure good convergence of the fyp--likelihood).," For easy comparison with the previous analysis of \citet{Eriksen_etal_2004}, the FWHM widths, $\theta_{\rm FWHM}$ , selected for these Gaussian smoothingbeams are taken to be $0\fdg53$, $0\fdg64$, $0\fdg85$, $1\fdg28$, $1\fdg70$, $2\fdg13$, $2\fdg55$, $2\fdg98$ and $3\fdg40$ (we abandon the larger angular scales used in the former analysis to ensure good convergence of the -likelihood)."
 Since a higher resolution map is necessary for more accurate estimation of the skeleton statistics (see Appendix A2). the resolution xwameter of the resulting smoothed map is set to Ia=1024.," Since a higher resolution map is necessary for more accurate estimation of the skeleton statistics (see Appendix \ref{sec_app_1}) ), the resolution parameter of the resulting smoothed map is set to $N_{\rm side}=1024$ ."
 After the smoothing of the data. the base-mask must also be expanded and the same processing method. is ollowed.," After the smoothing of the data, the base-mask must also be expanded and the same processing method is followed."
 For each smoothing scale. only those base-mask Xxels with values larger than 0.99 are defined to be valid jxixels on the smoothed mask.," For each smoothing scale, only those base-mask pixels with values larger than 0.99 are defined to be valid pixels on the smoothed mask."
 Finally. on the valid region defined by the smootheel owe-nnsk. the map 1(ρ). cither from. observation or simulation. is renormalised το temperature thresholds. wv (σα) while the invalid. pixels are abandoned: for computing the standard deviation.," Finally, on the valid region defined by the smoothed base-mask, the map $\tilde T(p)$, either from observation or simulation, is renormalised to temperature thresholds, $\nu$ \ref{eqn_renorm}) ), while the invalid pixels are abandoned for computing the standard deviation."
 Using the method. discussed in Appendix A.. the skeleton length distributions. Dato) and ον]. can be estimated for each set of smootheel samples.," Using the method discussed in Appendix \ref{sec_app_1}, the skeleton length distributions, $\mathcal{L}_{d}(\nu)$ and $\mathcal{L}_{a}(\nu)$, can be estimated for each set of smoothed samples."
 The original distribution £07) in Eq., The original distribution $L(\nu)$ in Eq.
 6. is divided into 200 bins with vyC—4.0.4.0] during skeleton tracing.," \ref{eq_Ld_nu} is divided into 200 bins with $\nu \in [-4.0, 4.0]$ during skeleton tracing."
 From the processed Gaussian simulations. we compute the Gaussian expectation of the skeleton statistics for each smoothing scale. CUGhpwaY.," From the processed Gaussian simulations, we compute the Gaussian expectation of the skeleton statistics for each smoothing scale, $\langle \mathcal{L}^{\rm G}(\nu, \theta_{\rm FWHM})
\rangle$."
 The departure from these expectation values is then obtained for both the observed data and cach Gaussian sample as and the corresponding x7 value is then computed where we omit the £N term since it is definitelv zero., The departure from these expectation values is then obtained for both the observed data and each Gaussian sample as and the corresponding $\chi^2$ value is then computed where we omit the $\langle\Delta \mathcal{L}\rangle$ term since it is definitely zero.
" ln the aanalvsis. the non-Ciaussian departure. LN(p,ΓκιPpar). and the 47 statisties are estimated. by ]5q.10. and 11.."," In the analysis, the non-Gaussian departure, $\mathcal{L}^{\rm NG}(\nu,f_{\rm NL},\theta_{\rm FWHM})$, and the $\chi^2$ statistics are estimated by \ref{eq_ske_NG} and \ref{eq_chi2}."
 Phe best-fit value and error of ccan then be obtained by analysing the likelihood functionas discussed in Section 2.2.., The best-fit value and error of can then be obtained by analysing the likelihood functionas discussed in Section \ref{subsec_nulltest}. .
 Aefore we provide final estimates of. {from the cillerent smoothing scales. we combined the estimators. AL(v) of the data and 2S(v.fre) of cach set of ssaniple. to and," Before we provide final estimates of from the different smoothing scales, we combined the estimators, $\Delta
\mathcal{L}(\nu)$ of the data and $\mathcal{L}^{\rm NG}(\nu, f_{\rm
NL})$ of each set of sample, to and"
Despite these complicated parameter dependencies. we have shown that the OI]62. μπα line flux depends basically on two quantities. namely the total dise gas mass and the mean disc temperature.,"Despite these complicated parameter dependencies, we have shown that the $\rm[OI]\,63.2\mu$ m line flux depends basically on two quantities, namely the total disc gas mass and the mean disc temperature."
 We will continue this work by two follow-up papers (Kamp et al., We will continue this work by two follow-up papers (Kamp et al.
 2010. in prep..," 2010, in prep.,"
 Ménnard et al., Ménnard et al.
 2010. in prep.)," 2010, in prep.)"
 that will provide more insight into the statistical behaviour of gas line and dust continuum predictions. respectively. to identify trends and robust correlations with dise mass.," that will provide more insight into the statistical behaviour of gas line and dust continuum predictions, respectively, to identify trends and robust correlations with disc mass."
 In summary. the DENT grid allows We intend to make the ealeulated DENT grid available to the scientific community.," In summary, the DENT grid allows We intend to make the calculated DENT grid available to the scientific community."
 A graphical user interface called has been developed to allow researchers to visualise the DENT results. to make plots as presented in this letter. and to search for models for a given set of continuum and line flux data.," A graphical user interface called has been developed to allow researchers to visualise the DENT results, to make plots as presented in this letter, and to search for best-fitting models for a given set of continuum and line flux data."
 We emphasise. however. that the DENT grid has not been developed for detailed fitting of individual objects.," We emphasise, however, that the DENT grid has not been developed for detailed fitting of individual objects."
 The coarse sampling of the [2-dimensional parameter space can mostly be used to narrow down the parameter range for individual objects. for example to design a finer sub-grid. especially for not so well-known objects.," The coarse sampling of the 12-dimensional parameter space can mostly be used to narrow down the parameter range for individual objects, for example to design a finer sub-grid, especially for not so well-known objects."
 With a comprehensive data set of far IR. gas emission lines to be obtained by Herschel/GASPS very soon. we aim at breaking the degeneracy of SED fitting and make possible a more profound analysis of the physical. chemical and temperature structures of discs around young stars.," With a comprehensive data set of far IR gas emission lines to be obtained by Herschel/GASPS very soon, we aim at breaking the degeneracy of SED fitting and make possible a more profound analysis of the physical, chemical and temperature structures of discs around young stars."
 Let us assume that an emission line is optically thin and that the emitting species is populated with a uniform excitation temperature Tix., Let us assume that an emission line is optically thin and that the emitting species is populated with a uniform excitation temperature $T_{\rm exc}$.
" The line luminosity is then given by where v is the line centre frequency. .1,; is the Einstein coefficient of the line transition from level i to level /. 6) is the partition function. and g, and ££, are the statistical weight and energy [K] of the upper level."," The line luminosity is then given by where $\nu$ is the line centre frequency, $A_{ul}$ is the Einstein coefficient of the line transition from level $u$ to level $l$, $Q$ is the partition function, and $g_u$ and $E_u$ are the statistical weight and energy [K] of the upper level."
 The total number of line emitting particles Ay.) is related to the total dise gas mass by, The total number of line emitting particles $N_{\rm tot}$ is related to the total disc gas mass by
(Lalavοἱal.1991).,\citep{Lah91}.
. This implies that clusters even in (his model should be dynamically older (han in an Ed$ model (Richstoneetal.1992:EvrardBeishart.Valdarnini&Buehert2002:Suwaοἱal.," This implies that clusters even in this model should be dynamically older than in an EdS model \citep{Ri,Ev,Lac,Bei,Suwa}."
 2002).. Therefore one should be able to put constraints on O [rom the evolution of the cluster cdvnamical state., Therefore one should be able to put constraints on $\Omega$ from the evolution of the cluster dynamical state.
 One such indicator. the cluster ellipticity. was recently proposed in (2001): the more elliptical a cluster the vounger it is. since ellipticity should be correlated with the anisotropic accretion of matter and merging (e.g..Duote1998:lolokotronis 2001).," One such indicator, the cluster ellipticity, was recently proposed in \citet{MelC}; the more elliptical a cluster the younger it is, since ellipticity should be correlated with the anisotropic accretion of matter and merging \citep{Buote, KBPG}."
. Melottetal.(2001) presented. evidence for a recent evolution (2<0.1) of cluster ellipticity. defined in both optical aud X-ray. bands. based on various cluster samples. containing in total ~160 rich Abell clusters.," \citet{MelC} presented evidence for a recent evolution $z\lesssim 0.1$ ) of cluster ellipticity, defined in both optical and X-ray bands, based on various cluster samples, containing in total $\sim 160$ rich Abell clusters."
 An earlier study. using clusters identified in the Lick map (Plionisetal.1991).. had also Found that cluster ellipticity decreases with reclshilt [or zS0.1. due however to possible svstematic effects involved in the construction of the data. the authors did not attach much weight to this result.," An earlier study, using clusters identified in the Lick map \citep{PBF91}, had also found that cluster ellipticity decreases with redshift for $z\lesssim 0.1$, due however to possible systematic effects involved in the construction of the data, the authors did not attach much weight to this result."
" This interesting finding was interpreted in Melottetal.(2001). às an indication of a low-O,, universe. because in such a universe and under (he hierarchical clustering scenario. one expects that mereine and anisotropic accretion of matter along filaments Klvpin1984). will have stopped long ago."," This interesting finding was interpreted in \cite{MelC} as an indication of a $\Omega_{m}$ universe, because in such a universe and under the hierarchical clustering scenario, one expects that merging and anisotropic accretion of matter along filaments \citep{Shan} will have stopped long ago."
 Thus the clusters should be relatively isolated ancl eravitational relaxation will tend (to isotropize the clusters reducing their ellipticity. more so in the recent times.," Thus the clusters should be relatively isolated and gravitational relaxation will tend to isotropize the clusters reducing their ellipticity, more so in the recent times."
 A variety of indications do exist supporting the formation of clusters by hierarchical aggregation of smaller units along filamentary large-scale structures Jones&Forman1995:PlionisBasilakos 2002).," A variety of indications do exist supporting the formation of clusters by hierarchical aggregation of smaller units along filamentary large-scale structures \citep{West, PB02}."
". Although the cluster ellipticity is a relatively well defined. quantity. svstematic effects due to projections in (he optical. that could increase the ellipticity of spherical clusters or decrease (he ellipticity of flat ones. or the strong central concentration of (le X-ray emitting eas (since L,x n2) in the N-ray band. which in conjuction with resolution effects could reduce the ellipticity of distant clusters. should be taken into account."," Although the cluster ellipticity is a relatively well defined quantity, systematic effects due to projections in the optical, that could increase the ellipticity of spherical clusters or decrease the ellipticity of flat ones, or the strong central concentration of the X-ray emitting gas (since $L_{x} \propto n^{2}_{e}$ ) in the X-ray band, which in conjuction with resolution effects could reduce the ellipticity of distant clusters, should be taken into account."
 One should therefore use large and well constructed samples of galaxy. clusters (ο verily such an ellipticity. trend., One should therefore use large and well constructed samples of galaxy clusters to verify such an ellipticity trend.
 Furthermore. other indicators of recent cluster evolution should be used as well.," Furthermore, other indicators of recent cluster evolution should be used as well."
 For example. if the ellipticity decrease with z is indeed due (o recent clvnamical relaxation processes. then one should expect an evolution of the temperature of the N-rav emitüng gas as well as the X-ray cluster luminosity which should follow the same trend as (he cluster ellipticity. decreasing al recent times. since the violent merging events. at relatively higher recshifts. will compress and shock heat the dilfuse ICM eas (e.g..Ritchie&Thomas2002).," For example, if the ellipticity decrease with $z$ is indeed due to recent dynamical relaxation processes, then one should expect an evolution of the temperature of the X-ray emitting gas as well as the X-ray cluster luminosity which should follow the same trend as the cluster ellipticity, decreasing at recent times, since the violent merging events, at relatively higher redshifts, will compress and shock heat the diffuse ICM gas \citep{Ritchi}."
. Another possible indicator could be the cluster velocity dispersion. which naively one would expect {ο increase al lower redshifts. since virializalion will tend to increase the cluster ‘thermal’ velocity dispersion.," Another possible indicator could be the cluster velocity dispersion, which naively one would expect to increase at lower redshifts, since virialization will tend to increase the cluster `thermal' velocity dispersion."
" Although there is a quite well defined £,—o relationship", Although there is a quite well defined $L_x-\sigma$ relationship
"where we apply a prescription for QHE, and under-predict the observed line strength.","where we apply a prescription for QHE, and under-predict the observed line strength."
" This indicates that accounting for binary stellar evolution alone cannot entirely explain the observed profile, but rotation in binaries may play an important role in producing the observed spectra."," This indicates that accounting for binary stellar evolution alone cannot entirely explain the observed profile, but rotation in binaries may play an important role in producing the observed spectra."
 Because the spectrum analysed here is a composite of a number of different galaxies we can only attempt to determine a mean or range of metallicities., Because the spectrum analysed here is a composite of a number of different galaxies we can only attempt to determine a mean or range of metallicities.
 A qualitative analysis suggests an allowed metallicity range of 0.001«Z«0.008 for the ‘typical’ Lyman-break galaxy at z~3., A qualitative analysis suggests an allowed metallicity range of $0.001 < Z < 0.008$ for the `typical' Lyman-break galaxy at $z\sim3$.
" This corresponds to a 7.6«[O/H]8.5 (or 0.1<Z/Zo«0.6 where [O/H]= 8.7), and is consistent with the previous measurements determined by Shapleyetal.(2003) and etal. (2004).."," This corresponds to a $ 7.6 < [O/H] < 8.5$ (or $ 0.1 <
Z/Z_{\odot} < 0.6$ where $[O/H]_{\odot}=8.7$ ), and is consistent with the previous measurements determined by \citet{shapley} and \citet{2004ApJ...615...98R}."
" We address the spectra of the individual lensed galaxies compiled here starting with the sources with the lowest inferred metallicity, based on our model comparisons."," We address the spectra of the individual lensed galaxies compiled here starting with the sources with the lowest inferred metallicity, based on our model comparisons."
" In each case, we make the comparison at a constant log(age/years)=8.5 MMyr), based on the median star-formation derived age of the z~3 Lyman-break galaxy population as a whole."," In each case, we make the comparison at a constant log(age/years)=8.5 Myr), based on the median star-formation derived age of the $z\sim3$ Lyman-break galaxy population as a whole."
" We note that, as figure 3 indicates, there is little evolution in the model strengths of our lines of interest at stellar population ages above ~100 MMyr."," We note that, as figure \ref{modelew}
 indicates, there is little evolution in the model strengths of our lines of interest at stellar population ages above $\sim$ Myr."
spin distributions found using these independent methods are formally consistent.,spin distributions found using these independent methods are formally consistent.
 The sample of natal neutron star periods is also small. with only nine estimates.," The sample of natal neutron star periods is also small, with only nine estimates."
 Assuming the extreme stellar parameters. a Gaussian fit to this distribution gives dea=2.1\107. and o=1.9«1077.," Assuming the extreme stellar parameters, a Gaussian fit to this distribution gives $a_{\rm cent} = 2.1\times
10^{-2}$, and $\sigma = 1.9\times 10^{-2}$."
 Assuming canonical stellar parameters. a Gaussian fit gives au=9.5«107. and g284.107.," Assuming canonical stellar parameters, a Gaussian fit gives $a_{\rm cent} = 9.5\times
10^{-3}$, and $\sigma = 8.4\times 10^{-3}$."
 Table 3 lists the results of running two-sided Kolmogorov-Smirnov tests on different angular momentum distributions., Table 3 lists the results of running two-sided Kolmogorov-Smirnov tests on different angular momentum distributions.
 This statistical test measures the probability that two samples are drawn from the same parent distribution., This statistical test measures the probability that two samples are drawn from the same parent distribution.
 For the samples and assumed stellar parameters considered in. this. work. the highest probability that the neutron star and black hole distributions are drawn from the same parent is just 3.6«1077.," For the samples and assumed stellar parameters considered in this work, the highest probability that the neutron star and black hole distributions are drawn from the same parent is just $3.6\times 10^{-4}$."
 Despite the fact that the sample of black hole spin parameters and natal neutron star spin periods is small. it is clear that they belong to distinct distributions.," Despite the fact that the sample of black hole spin parameters and natal neutron star spin periods is small, it is clear that they belong to distinct distributions."
 Black holes are likely born with high spin parameters. and neutron stars may be born with very low spin parameters.," Black holes are likely born with high spin parameters, and neutron stars may be born with very low spin parameters."
 At the time of writing. the Australia Telescope National Facility Pulsar Catalog (Manchester et 22005) contains 1875 rotation-powered pulsars.," At the time of writing, the Australia Telescope National Facility Pulsar Catalog (Manchester et 2005) contains 1875 rotation-powered pulsars."
 The spins of the pulsars in this compilation do not reflect birth conditions., The spins of the pulsars in this compilation do not reflect birth conditions.
 We simply note that no pulsar in this sample has a dimensionless spin parameter in excess of ¢=0.3. neither for typical nor extreme stellar parameters. and the sample is statistically distinct from the black hole distributions at extremely high confidence.," We simply note that no pulsar in this sample has a dimensionless spin parameter in excess of $a = 0.3$, neither for typical nor extreme stellar parameters, and the sample is statistically distinct from the black hole distributions at extremely high confidence."
 It is difficult to assess the degree to which current black hole spin measurements may be affected by systematic errors and skew comparisons to the sample of neutron star angular momenta., It is difficult to assess the degree to which current black hole spin measurements may be affected by systematic errors and skew comparisons to the sample of neutron star angular momenta.
 The thermal continuum emission from stellar-mass black hole accretion disks can be fit well with simple. two-parameter models that lack inner torque conditions and provide no direct spin. constraint.," The thermal continuum emission from stellar-mass black hole accretion disks can be fit well with simple, two-parameter models that lack inner torque conditions and provide no direct spin constraint."
 New models that include spin às a variable parameter are more complex. and rely on accurate measures of the mass and distance of the black hole. as well as reasonable estimates of the mass accretion rate and scattering in the disk atmosphere (see. e.g.. McClintock et 22006: Nowak et 22008).," New models that include spin as a variable parameter are more complex, and rely on accurate measures of the mass and distance of the black hole, as well as reasonable estimates of the mass accretion rate and scattering in the disk atmosphere (see, e.g., McClintock et 2006; Nowak et 2008)."
 Spin constraints derived based on fits to disk reflection spectra have primarily made use of relatively simple models which assume a constant disk density (see. e.g.. Miller et 22009).," Spin constraints derived based on fits to disk reflection spectra have primarily made use of relatively simple models which assume a constant disk density (see, e.g., Miller et 2009)."
 A few spectral fits have been made with more physical reflection models: these fits appear to give consistent results (see. e.g.. Rets et 22008). but more work is needed.," A few spectral fits have been made with more physical reflection models; these fits appear to give consistent results (see, e.g., Reis et 2008), but more work is needed."
 For three black holes in Table I. spins have been measured using both techniques. and the values do not agree in detail.," For three black holes in Table 1, spins have been measured using both techniques, and the values do not agree in detail."
 All of the spin values and their associated errors were derived based on spectral fitting that minimized a  statistic., All of the spin values and their associated errors were derived based on spectral fitting that minimized a $\chi^{2}$ statistic.
 The resulting \7 space can be complex. and using a Io error to evaluate the degree to which measurements disagree can give an over-estimate (see Miller et 22009).," The resulting $\chi^{2}$ space can be complex, and using a $\sigma$ error to evaluate the degree to which measurements disagree can give an over-estimate (see Miller et 2009)."
 The measurements made for 4U 1543-475 are most strongly at odds., The measurements made for 4U $-$ 475 are most strongly at odds.
 In this case. the low resolution and calibration of the gas proportional counter spectrum may serve to under-estimate the extent of the iron line: a single spectral bin could account for 0.1 —0.2.," In this case, the low resolution and calibration of the gas proportional counter spectrum may serve to under-estimate the extent of the iron line; a single spectral bin could account for $\Delta(a) \simeq 0.1-0.2$ ."
 The low spin value for Cygnus X-1 may be another good example of a suspicious value., The low spin value for Cygnus X-1 may be another good example of a suspicious value.
" It derives from fits to the source in à ""low/hard"" state. wherein the disk may not always extend to the ISCO (see. e.g. Esin. McClintock. Narayan 1997)."," It derives from fits to the source in a “low/hard” state, wherein the disk may not always extend to the ISCO (see, e.g. Esin, McClintock, Narayan 1997)."
 Measurements of the inner disk radius based on relativistic reflection modeling in states where the disk is more certain to extend to the ISCO give values as small as ry22.241.0GM/c? (Miller et 22005). which translates into «20.9751.," Measurements of the inner disk radius based on relativistic reflection modeling in states where the disk is more certain to extend to the ISCO give values as small as $r_{in} = 2.2\pm
1.0~ GM/c^{2}$ (Miller et 2005), which translates into $a =
0.9^{+0.1}_{-0.2}$."
 In two cases. then. potential systematic errors. resulting from limited spectral resolution and source phenomenology may have served to under-estimate the degree to which black hole and neutron star angular momentum distributions differ.," In two cases, then, potential systematic errors resulting from limited spectral resolution and source phenomenology may have served to under-estimate the degree to which black hole and neutron star angular momentum distributions differ."
 An uncertainty in basic accretion disk physics. however. could potentially work in the opposite sense.," An uncertainty in basic accretion disk physics, however, could potentially work in the opposite sense."
 Both techniques assume that the inner edge of an actual accretion disk is truncated at the test-particle ISCO., Both techniques assume that the inner edge of an actual accretion disk is truncated at the test-particle ISCO.
 If real fluid aceretion disks push across the ISCO defined by test particle orbits. both spin. measurement techniques would give falsely high spin values.," If real fluid accretion disks push across the ISCO defined by test particle orbits, both spin measurement techniques would give falsely high spin values."
 New numerical simulations aimed at addressing this specific Issue appear to justify the underlying assumption of both techniques (see. e.g.. Shafee et 22008. Reynolds Fabian 2008).," New numerical simulations aimed at addressing this specific issue appear to justify the underlying assumption of both techniques (see, e.g., Shafee et 2008, Reynolds Fabian 2008)."
 However. this must be continually revisited as numerical techniques and computing power advance.," However, this must be continually revisited as numerical techniques and computing power advance."
 The dimensionless angular momenta of stellar-mass black holes and rotation-powered neutron stars are compared statistically., The dimensionless angular momenta of stellar-mass black holes and rotation-powered neutron stars are compared statistically.
 Two different methods of measuring black hole spin. parameters are considered. and both canonical and extreme stellar parameters were considered for neutron stars with estimated natal spin periods.," Two different methods of measuring black hole spin parameters are considered, and both canonical and extreme stellar parameters were considered for neutron stars with estimated natal spin periods."
 A strong statistical difference between the distributions signals that black holes could typically born with dramatically higher dimensionless angular momentum than neutron stars., A strong statistical difference between the distributions signals that black holes could typically born with dramatically higher dimensionless angular momentum than neutron stars.
 The observed dichotomy 1n their dimensionless angular momenta implies that there may be important differences in the core-collapse supernovae events that create these stellar remnants., The observed dichotomy in their dimensionless angular momenta implies that there may be important differences in the core-collapse supernovae events that create these stellar remnants.
 The collapse of a 35M.. star with a 14M.. He core is examined in detail by MacFadyen Woosley (1999).," The collapse of a $35~{\rm M}_{\odot}$ star with a $14~{\rm
 M}_{\odot}$ He core is examined in detail by MacFadyen Woosley (1999)."
 Reasonable pre-collapse angular momentum values for the eventual tron core are found to lead to the formation of a disk that rapidly aceretes onto a newly-formed black hole. quickly giving à spin parameter of «—0.9.," Reasonable pre-collapse angular momentum values for the eventual iron core are found to lead to the formation of a disk that rapidly accretes onto a newly-formed black hole, quickly giving a spin parameter of $a \simeq 0.9$."
 Our results are consistent with this prediction (see Tables | and 2. and Figure 1). lending support to the overall model described by MacFadyen Woosley (1999).," Our results are consistent with this prediction (see Tables 1 and 2, and Figure 1), lending support to the overall model described by MacFadyen Woosley (1999)."
" Of course. this model is the standard ""eollapsar"" picture for the production of gamma-ray bursts. and it is interesting to note that the same disk required for driving the new black hole to a high spin parameter is also responsible for producing a jet via MHD processes."," Of course, this model is the standard “collapsar” picture for the production of gamma-ray bursts, and it is interesting to note that the same disk required for driving the new black hole to a high spin parameter is also responsible for producing a jet via MHD processes."
 Shibata Shapiro (2002) used simulations to understand the spin of the black hole left by a uniformly rotating supermassive star. finding «~0.75.," Shibata Shapiro (2002) used simulations to understand the spin of the black hole left by a uniformly rotating supermassive star, finding $a \simeq 0.75$."
 Shapiro Shibata (2002) later showed that the same result can be derived analytically (also see Gammie. Shapiro. McKinney 2004).," Shapiro Shibata (2002) later showed that the same result can be derived analytically (also see Gammie, Shapiro, McKinney 2004)."
 These treatments are primarily concerned with extremely massive stars as candidate progenitors for the build-up of supermassive black holes in the early universe., These treatments are primarily concerned with extremely massive stars as candidate progenitors for the build-up of supermassive black holes in the early universe.
 Although our results are consistent with these predictions as, Although our results are consistent with these predictions as
also note that “curiously. the [observed] cool gas sccus to fori a shell around the radio lobes aud occupies preciscly the location where we would expect shock/compressionallv heated eas to reside”.,"also note that “curiously, the [observed] cool gas seems to form a shell around the radio lobes and occupies precisely the location where we would expect shock/compressionally heated gas to reside”."
 Regardless of the heating geometry assunied. by jets or ad hoc off-center bubbles. the global. long-term rate that hot gas cools bv radiative losses iu cooling flows is unlikely to change without altering the eas density aud temperature profiles in disagreement witli observations. (," Regardless of the heating geometry assumed, by jets or ad hoc off-center bubbles, the global, long-term rate that hot gas cools by radiative losses in cooling flows is unlikely to change without altering the gas density and temperature profiles in disagreement with observations. ("
Drigheuti Mathews 2002).,Brighenti Mathews 2002).
 We begin by constructing a simple cooling flow for a massive elliptical galaxy., We begin by constructing a simple cooling flow for a massive elliptical galaxy.
" The gravitational potential is a superposition of a de Vaucouleurs stellar mass profile of total mass Ma,=7.26«1013 AL. and a Navarro-Freuk-White dark halo with virial mass My,;,=|<lol? 1Η, and concentration ο=10. typical for the galaxw group environments in which massive E ealaxies are thought to form."," The gravitational potential is a superposition of a de Vaucouleurs stellar mass profile of total mass $M_{*t} = 7.26 \times 10^{11}$ $M_{\odot}$ and a Navarro-Frenk-White dark halo with virial mass $M_{h,vir} = 4 \times 10^{13}$ $M_{\odot}$ and concentration $c = 10$, typical for the galaxy group environments in which massive E galaxies are thought to form."
" The initial cooling flow that we consider is generated cutirely by mass loss from an old stellar population. Osp.=LTs1030(4t.)1s enis Lowhere pur) is the stellar density and f£,=13 Covrs is the assumed age of the stars."," The initial cooling flow that we consider is generated entirely by mass loss from an old stellar population, $\alpha_* \rho_* = 4.7 \times 10^{-20}(t/t_n)^{-1.3}$ gm $^{-1}$ where $\rho_*(r)$ is the stellar density and $t_n = 13$ Gyrs is the assumed age of the stars."
" Some heating is provided bv Type Ia supernovac. exploding at a rate of 0.06(t/t,,)! per 100 vears por 1019Z5. each releasing 1075 eres."," Some heating is provided by Type Ia supernovae, exploding at a rate of $0.06 (t/t_n)^{-1}$ per 100 years per $10^{10}L_B$, each releasing $10^{51}$ ergs."
 Supernova heating at us level is consistent with the iron abundance observed oei the hot eas but has a relatively minor influence ou re overall energeties of the hot eas., Supernova heating at this level is consistent with the iron abundance observed in the hot gas but has a relatively minor influence on the overall energetics of the hot gas.
 The οας-ανασα] equations that we solve are identical to those described in Brighenti Mathews (2002)., The gas-dynamical equations that we solve are identical to those described in Brighenti Mathews (2002).
 We use a 2D code modeled after ZEUS (Stone Nonuau 1992) with 190 exid zones in th exdliudrical coordinates., We use a 2D code modeled after ZEUS (Stone Norman 1992) with 490 grid zones in both cylindrical coordinates.
 For Πιτ«50 kpc the zones are uniform with size 125« pe.," For $R,z < 50$ kpc the zones are uniform with size $125 \times 125$ pc."
 Outside this reeion the size increases ecoletiically toward the outer boundary at 260900 kpc., Outside this region the size increases geometrically toward the outer boundary at $260 \times 260$ kpc.
 The quiescent cooling flow that results after slowly evolving for 13 Cis has a characteristic teiiperature of —1 keV aud a eas deusity profile pxc.|? within about LO kpc and slightly. steeper bevoud., The quiescent cooling flow that results after slowly evolving for 13 Gyrs has a characteristic temperature of $\sim 1$ keV and a gas density profile $\rho \propto r^{-1.2}$ within about 10 kpc and slightly steeper beyond.
" We assuue that eas within rj,=1l kpe of the ceuter is heated at a coustaut rate Lj=2.5«10% ere ο for fj=5«10* vrs.", We assume that gas within $r_h = 1$ kpc of the center is heated at a constant rate $L_h = 2.5 \times 10^{42}$ erg $^{-1}$ for $t_h = 5 \times 10^7$ yrs.
 The jeated gas reaches temperatures far im excess of 1 keV aud las a sound crossing time much less than fj. allowing au ultra-hot low-density region to develop ucar the galactic core.," The heated gas reaches temperatures far in excess of 1 keV and has a sound crossing time much less than $t_h$, allowing an ultra-hot low-density region to develop near the galactic core."
 The hot eas pushes outward. driving a shock iuto he cooling flow gas bevoud.," The hot gas pushes outward, driving a shock into the cooling flow gas beyond."
" The resulting evolution of he eas density and temperature are shown in Figure | at lice times ty=LOS108. #2=δις10%, and £4=100«109 vrs."," The resulting evolution of the gas density and temperature are shown in Figure 1 at three times $t_1 = 40 \times 10^6$, $t_2 = 80 \times 10^6$, and $t_3 = 100 \times 10^6$ yrs."
 The Ravleigh-Tavlor (RT) developiueut of a few larec phunes filled with heated eas produces regular N-rav xiteris that resemble. images of NGC 50LL aud other E galaxies., The Rayleigh-Taylor (RT) development of a few large plumes filled with heated gas produces irregular X-ray patterns that resemble images of NGC 5044 and other E galaxies.
 Suiall waveleneth RT unstable features are nof prominent even though they would be resolved by he ummerical exid., Small wavelength RT unstable features are not prominent even though they would be resolved by the numerical grid.
 The hot cavity expands ecutly aud is already subsonic CE500 kins +) at #=1 Myr., The hot cavity expands gently and is already subsonic $\lta 500$ km $^{-1}$ ) at $t = 1$ Myr.
 By time fq he weakened precursor shock has moved out to 20 kpe. noving aliiost. sonically.," By time $t_1$ the weakened precursor shock has moved out to $\sim 20$ kpc, moving almost sonically."
 Low temperature. bright rus Μποιο the hot phumes/Dubbles are clearly visible iu Figure 1.," Low temperature, bright rims surrounding the hot plumes/bubbles are clearly visible in Figure 1."
 The peak density in the rius at f4 is several times thatof the initial uushocked eas at this radius., The peak density in the rims at $t_1$ is several times thatof the initial unshocked gas at this radius.
" The minima cutropy oe(T.on25j in the sinulatiou at times t=0. fy. fo and ts are 7.15. 7.9. 8.1 aud 8.0 respectively,"," The minimum entropy $\log(T/n_e^{2/3})$ in the simulation at times $t =
0$, $t_1$, $t_2$ and $t_3$ are 7.15, 7.9, 8.1 and 8.0 respectively."
 Low eutropy eas imn the runs at f&f4 was originally at r~ry as verified with passively advected tracer particles.," Low entropy gas in the rims at $t \gta t_1$ was originally at $r \sim
r_h$ as verified with passively advected tracer particles."
 This eas cools by adiabatic expansion while rising iu the hot gas atmosphere (see also Nilsen ct al., This gas cools by adiabatic expansion while rising in the hot gas atmosphere (see also Nulsen et al.
 2002: Soker et al., 2002; Soker et al.
 2002)., 2002).
 The larger bubble at time f£» is penetrated by a disorganized S-shaped coluun of slowly rising cold eas., The larger bubble at time $t_2$ is penetrated by a disorganized S-shaped column of slowly rising cold gas.
 The upper panel of Figure 2 shows the X-ray surface brightuess of the central cooling flow at time to when the plies have separated into clearly visible buovaut N-rav bubbles., The upper panel of Figure 2 shows the X-ray surface brightness of the central cooling flow at time $t_2$ when the plumes have separated into clearly visible buoyant X-ray bubbles.
 It is important to stress that the physical size of the phunes that form is not determined by our conrutational resolution., It is important to stress that the physical size of the plumes that form is not determined by our computational resolution.
 C'aleulatious repeated at higher spatial resolution (ceutral grid size=50 pc) also produce a few plumes and Qutimately) bubbles of about the same size as those shown in Figures 1 aud 2. which are similar to the X-ray patterus observed.," Calculations repeated at higher spatial resolution (central grid size=50 pc) also produce a few plumes and (ultimately) bubbles of about the same size as those shown in Figures 1 and 2, which are similar to the X-ray patterns observed."
 An exiuniuation of the effective eravity in the eas rest frame indicates that the outer rins of the large bubbles are RT uustable., An examination of the effective gravity in the gas rest frame indicates that the outer rims of the large bubbles are RT unstable.
 At time ts the eravitational τοσα time across the bubble radius is typ©U.Mfa: coherent hot bubbles persist for times 2fyr. even at our highest computational resolution.," At time $t_3$ the gravitational freefall time across the bubble radius is $t_{ff} \approx 0.1 t_3$; coherent hot bubbles persist for times $\gg t_{ff}$, even at our highest computational resolution."
 We performed similar calculations on more extended cooling flows iu galaxy groups aud clusters with very simular results., We performed similar calculations on more extended cooling flows in galaxy groups and clusters with very similar results.
" For example. the lower paucl of Figure 2 illustrates the N-rav surface brightness at time f=108 years produced iu a vich cluster cooling flow (Mj=Lor AZ. aud Ty;%6.9 keV) with heating parameters: Ly=10 ores boy, =2 kpe aud fj=5ν10* τις,"," For example, the lower panel of Figure 2 illustrates the X-ray surface brightness at time $t = 120 \times 10^6$ years produced in a rich cluster cooling flow $M_{vir} = 10^{15}$ $M_{\odot}$ and $T_{vir} \approx 6.9$ keV) with heating parameters: $L_h = 10^{45}$ ergs $^{-1}$, $r_h = 2$ kpc and $t_h = 5 \times 10^7$ yrs."
 Very little gas cools by radiation losses during the brief duratiou of our calculations shown iu Figures 1 aud 2 so the flow is esseutiallv adiabatic., Very little gas cools by radiation losses during the brief duration of our calculations shown in Figures 1 and 2 so the flow is essentially adiabatic.
 This is expected since a defining attribute of cooling flows is that the dynamical time is uch lessthan theraciative cooling time., This is expected since a defining attribute of cooling flows is that the dynamical time is much lessthan theradiative cooling time.
 Cold rims are a generic feature of flows forced to expand into a nediuu of mereasing entropy., Cold rims are a generic feature of flows forced to expand into a medium of increasing entropy.
 This cau be illustrated with sel£siuibu flows described bv the variable yo=rsM (Courant Friedricehs 1918: Rogers. 1957: Parker. 1961: Chevalier Duanmra 1983).," This can be illustrated with self-similar flows described by the variable $\eta = r^{-\lambda}t$ (Courant Friedrichs 1948; Rogers, 1957; Parker, 1961; Chevalier Imamura 1983)."
" The adiabatic equatious for spherical How can be written iu erus of g provided &=(rflty). p=Aye""O(u) aud P=Ayre?oF?IIGp."," The adiabatic equations for spherical flow can be written in terms of $\eta$ provided $u = (r/t)U(\eta)$, $\rho = A_o r^{-\alpha} \Omega(\eta)$ and $P = A_o r^{2-\alpha} t^{-2} \Pi(\eta)$."
 The hot region of coutrallv jeated. σας can be represeuted by an expanding spherical yston at jg=a that drives an advancing shock at y=an<a., The hot region of centrally heated gas can be represented by an expanding spherical piston at $\eta = \eta_2$ that drives an advancing shock at $\eta = \eta_1 < \eta_2$.
 The solutions (iy). O(4) aud. Toy) cau o found by inward iuteeration from post-shock values Cy. Oy and Wy which are functions of A. 5=5/3 and he Mach. number of the shock M (e.g. Rogers 1957).," The solutions $U(\eta)$, $\Omega(\eta)$ and $\Pi(\eta)$ can be found by inward integration from post-shock values $U_1$ , $\Omega_1$ and $\Pi_1$ which are functions of $\lambda$, $\gamma = 5/3$ and the Mach number of the shock ${\cal M}$ (e.g. Rogers 1957)."
" We cousider the spherical flow that results for pistons uuiformly expanding (A= 1) iuto au isothermal. fully ionized imiediunu (T,=10* K) with decreasing electron density p,=0.05267"" ? where a=1.2 approximates the eas density observed m massive elliptical galaxies (6.8. Brigheuti Mathews 1997)."," We consider the spherical flow that results for pistons uniformly expanding $\lambda = 1$ ) into an isothermal, fully ionized medium $T_o = 10^7$ K) with decreasing electron density $n_e = 0.0526 r^{-\alpha}$ $^{-3}$ where $\alpha = 1.2$ approximates the gas density observed in massive elliptical galaxies (e.g. Brighenti Mathews 1997)."
 Figure 3 illustrates the flow eecnerated by three pistons driving shocks at r4=3 kpe with M= 1.5. 2 and 3.," Figure 3 illustrates the flow generated by three pistons driving shocks at $r_1 = 3$ kpc with ${\cal M} = 1.5$ , 2 and 3."
" In cach case the temperature drops sharply uear the piston. T,. creating cold rius where the deusitv rises accordingly."," In each case the temperature drops sharply near the piston, $T_o$ , creating cold rims where the density rises accordingly."
" Post-shock eas that has just cooled back to 7, inFigure 3. located at log(r/iry)= O.A79. -0.108. and -0.075. for"," Post-shock gas that has just cooled back to $T_o$ inFigure 3, located at $\log(r/r_1) = -0.179$ , -0.108, and -0.078 for"
inFTOOLS.,in.
 We are grateful to the referee for identifving shortcomings in our original discussions in Section 3.7., We are grateful to the referee for identifying shortcomings in our original discussions in Section 3.7.
ind for GRB 010222 does not necessarily mean that its lost galaxy is especially metal poor.,find for GRB 010222 does not necessarily mean that its host galaxy is especially metal poor.
 Le could. just. mean hat cold. dense. giant. molecular clouds are. particularly depleted.," It could just mean that cold, dense, giant molecular clouds are particularly depleted."
 Given the fact that the QSOs are unrelated to he absorption-line sites and the GRBs may not be. it is unclear whether the CRB host galaxies are special or onA he GRB environments within hosts.," Given the fact that the QSOs are unrelated to the absorption-line sites and the GRBs may not be, it is unclear whether the GRB host galaxies are special or only the GRB environments within hosts."
 The low resolution « our spectra does not allow us to clisentanele these two cases., The low resolution of our spectra does not allow us to disentangle these two cases.
"ga ""here is one case. GRB 000026. where Castro et al. ("," There is one case, GRB 000926, where Castro et al. ("
2001) ind that the circumburst material is formed by two cülleren clouds with a velocity separation of 168.,2001) find that the circumburst material is formed by two different clouds with a velocity separation of 168.
ον In these wo clouds the relative abundances of Zn and Cr are similar o each other. thus indicating that the metallicity and dus depletion are characteristic of the host galaxy. as a whole. and ruling out the hypothesis of grain. destruction in the vicinity of the burst.," In these two clouds the relative abundances of Zn and Cr are similar to each other, thus indicating that the metallicity and dust depletion are characteristic of the host galaxy as a whole, and ruling out the hypothesis of grain destruction in the vicinity of the burst."
 In the case of GRB 010222 there is also some substructure with a velocity separation of 106 (Castro et al., In the case of GRB 010222 there is also some substructure with a velocity separation of 106 (Castro et al.
 2001)., 2001).
 Lf these subsystems had similar abundances. then the metallicity and dust content that we find would. be representative of the ealaxy as a whole. and again. dust destruction by the burst itself would be ruled oul.," If these subsystems had similar abundances, then the metallicity and dust content that we find would be representative of the galaxy as a whole, and again, dust destruction by the burst itself would be ruled out."
 The observational evidence points out. that CRB 010222 has taken place in a clampecl Lyman-a svstem having the highest column density of neutral hydrogen and the lowest metallicity known up to now. and with low dust. content.," The observational evidence points out that GRB 010222 has taken place in a damped $\alpha$ system having the highest column density of neutral hydrogen and the lowest metallicity known up to now, and with low dust content."
 This result. reinforces the idea that a small cross-section. and not dust. is the cause of not finding DLA systems having very large NOLLD.," This result reinforces the idea that a small cross-section, and not dust, is the cause of not finding DLA systems having very large )."
 I the svstem associated with C1tuD 010222 is not atypical. this means that we may be missing an important fraction of at cosmological distances which in turn may have an impact on Opà (ic. the fraction of neutral gas in DLAs as a function of the closure density). although the exact contribution will depend on their cross section. which is unarguably small.," If the system associated with GRB 010222 is not atypical, this means that we may be missing an important fraction of at cosmological distances which in turn may have an impact on $\Omega_{DLA}$ (i.e. the fraction of neutral gas in DLAs as a function of the closure density), although the exact contribution will depend on their cross section, which is unarguably small."
 1t is not clear whether the dust is destroved by the GRB itself. or if its size clistribution is clilferent. being devoided of small particles.," It is not clear whether the dust is destroyed by the GRB itself, or if its size distribution is different, being devoided of small particles."
 In the latter case the feature due to silicate erains at 9.7 jin should be absent also., In the latter case the feature due to silicate grains at 9.7 $\mu$ m should be absent also.
 Finally. there is strong evidence that other GRB host galaxies (or at least their locations) are as well damped Lyman-a systems.," Finally, there is strong evidence that other GRB host galaxies (or at least their locations) are as well damped $\alpha$ systems."
 Since they seem as well associated with star formation. we conclude that they may occur in a sub-class of DLAs that associatec with active star formation.," Since they seem as well associated with star formation, we conclude that they may occur in a sub-class of DLAs that associated with active star formation."
 LS. and. P.M.V. acknowledge support. from NWO (post-doctoral grant no., I.S. and P.M.V. acknowledge support from NWO (post-doctoral grant no.
 614.051.003. ancl Spninoza grant respectivlev)., 614.051.003 and Spninoza grant respectivley).
 L.Ix. is supported bv the Roval Academy of Arts and Sciences in The Netherlands CINNAN)., L.K. is supported by the Royal Academy of Arts and Sciences in The Netherlands (KNAW).
 We thank the stalf at La Palma who performed the ToO observations., We thank the staff at La Palma who performed the ToO observations.
measurements of the local velocity field. and in Section 7 we provide our conclusions.,"measurements of the local velocity field, and in Section 7 we provide our conclusions."
 In the following we will take c=I., In the following we will take $c=1$.
 The luminosity distance. d;. to a supernova at redshift z is given by the definition where F is the observed flux and L the luminosity. or equivalently bywhere m is the apparent and M the absolute magnitude.," The luminosity distance, $d_L$, to a supernova at redshift $z$ is given by the definition where $F$ is the observed flux and $L$ the luminosity, or equivalently bywhere $m$ is the apparent and $M$ the absolute magnitude."
 The angular-diameter distance. d4. to the same supernova is detined as In a homogeneous and isotropic universe the two distance measures. dy and d4. are given by where we obtain d However. in a perturbed universe. the luminosity distance depends on the detailed trajectory of the individual photons from the supernova to the observer.," The angular-diameter distance, $d_A$, to the same supernova is defined as In a homogeneous and isotropic universe the two distance measures, $d_L$ and $d_A$, are given by where 0.2cm By taking the logarithmic derivative of with respect to $(1+z)$ in the special case of a flat universe $\Omega = 1$ ) we obtain where However, in a perturbed universe, the luminosity distance depends on the detailed trajectory of the individual photons from the supernova to the observer."
 At moderately high redshifts (z= 0.5) the contribution arising from lensing by intervening matter dominates. while at low redshifts (z 0.5) the contribution from the peculiar velocities of the supernova host galaxy relative to the observer dominates.," At moderately high redshifts $z
\gtrsim 0.5$ ) the contribution arising from lensing by intervening matter dominates, while at low redshifts $z \lesssim 0.5$ ) the contribution from the peculiar velocities of the supernova host galaxy relative to the observer dominates."
 In this paper we are only concerned with the local universe., In this paper we are only concerned with the local universe.
 From now on we will therefore only consider the contribution. from the peculiar velocities in determining distances., From now on we will therefore only consider the contribution from the peculiar velocities in determining distances.
" For à given supernova and observer. each with some peculiar velocity. the measured redshift z. angular-diameter distance. d4. and lummosity distance. dj. are modified according to Greene2006:Bonvinetal.2006b) where by definition A bar indicates the quantities. as measured in. à homogeneous and isotropic cosmology. and v,=ντ is the velocity of the supernova projected along the direction from the observer to the supernova."," For a given supernova and observer, each with some peculiar velocity, the measured redshift $z$, angular-diameter distance, $d_A$, and luminosity distance, $d_L$, are modified according to \citep{Hui:2005nm,Bonvin:2006en}
 where by definition A bar $\bar{\ldots}$ indicates the quantities as measured in a homogeneous and isotropic cosmology, and $v_r = v_e\cdot n$ is the velocity of the supernova projected along the direction from the observer to the supernova."
 Having measured the redshift. z. and the flux. F. of the supernova we can calculate the lummosity distance. dj. and the angular-diameter distance. dy. to within a scatter determined by the cosmic variance on the luminosity. £L. of a supernova.," Having measured the redshift, $z$, and the flux, $F$, of the supernova we can calculate the luminosity distance, $d_L$, and the angular-diameter distance, $d_A$, to within a scatter determined by the cosmic variance on the luminosity, $L$, of a supernova."
 As we know the peculiar velocity. vo. of the observer with respect to the CMB with great accuracy it is useful to collect all known quantities on the left hand side of the three equations: The three quantities Z. 4. and d; are the redshift. angular- distance. and luminosity distance as measured and calculated by an observer at rest with respect to the CMB.," As we know the peculiar velocity, $v_0$, of the observer with respect to the CMB with great accuracy it is useful to collect all known quantities on the left hand side of the three equations: The three quantities $\tilde{z}$, $\tilde{d}_A$, and $\tilde{d}_L$ are the redshift, angular-diameter distance, and luminosity distance as measured and calculated by an observer at rest with respect to the CMB."
 In the following we assume that the measured quantities all have been corrected to a frame at rest with respect to the CMB., In the following we assume that the measured quantities all have been corrected to a frame at rest with respect to the CMB.
 The measured redshift. =. is most transparently split into the cosmological redshift. 7. and the radial velocity of the supernova. v. by inverting to find ς. which ts then inserted into this alternative form of(13): It is also possible to begin with to obtain the cosmological redshift. z. in terms of the measured redshift. z. and the radial velocity. v. of the supernova: We now expand the natural logarithm of the angular-diameter distance. Indείς). to first order around the measured redshift. =. and by using we obtain where o4) for a flat universe is given by(6).," The measured redshift, $\tilde{z}$, is most transparently split into the cosmological redshift, $\bar{z}$, and the radial velocity of the supernova, $v_r$ , by inverting to find $\bar{z}$, which is then inserted into this alternative form of: It is also possible to begin with to obtain the cosmological redshift, $\bar{z}$, in terms of the measured redshift, $\tilde{z}$, and the radial velocity, $v_r\,$, of the supernova: We now expand the natural logarithm of the angular-diameter distance, $\ln\bar{d}_A(\bar{z})$, to first order around the measured redshift, $\tilde{z}$, and by using we obtain where $\alpha_A(z)$ for a flat universe is given by."
. This equation applies equally well for the luminosity distance. so we can give a common formula for both distance measures: This expression can easily be transformed into an equation between the measured apparent magnitude. 77. the calculated apparent magnitude. 7(2). at the measured redshift. =. and the peculiar radial velocity. v. of the supernova: Inverting this equation for the special case of a flat universe we obtain For low redshift supernovae the Hubble parameter. the angular-diameter distance. andthe luminosity distance can all be expanded in termsof the deceleration parameter. qo ::," This equation applies equally well for the luminosity distance, so we can give a common formula for both distance measures: This expression can easily be transformed into an equation between the measured apparent magnitude, $\tilde{m}$, the calculated apparent magnitude, $\bar{m}(\tilde{z})$, at the measured redshift, $\tilde{z}$ , and the peculiar radial velocity, $v_r\,$ , of the supernova: Inverting this equation for the special case of a flat universe we obtain For low redshift supernovae the Hubble parameter, the angular-diameter distance, andthe luminosity distance can all be expanded in termsof the deceleration parameter, $q_0$ :"
sindlar (o equations (25)-(28) above. equations (29)-(34) can be written in symbolic form as in which. as belore. the ο coefficients are defined by matching terms from equations (35)-(40) respectively with those of (29)-(34).,"Similar to equations (25)-(28) above, equations (29)-(34) can be written in symbolic form as in which, as before, the $c_{ij}$ coefficients are defined by matching terms from equations (35)-(40) respectively with those of (29)-(34)."
 But here there are many more coefficients that are zero. namely C1).C15.045:Can.C51C35.CatC31.C53.C51:," But here there are many more coefficients that are zero, namely $c_{14},c_{15},c_{16}; 
c_{25},c_{26}; c_{32},c_{36}; c_{51},c_{52},c_{54};c_{62}$."
 ο reduces the number ol svimbolic multiplvs needed {ο apply Cramers rule. rendering it practial to work out all the algebra [or (his 6X6 svstem.," This greatly reduces the number of symbolic multiplys needed to apply Cramers rule, rendering it practial to work out all the algebra for this 6X6 system."
(see.e.g..?)..,"\citep[see, e.g.,][]{Wang:08}."
 In the particular case of SN 2008D. polarization measurements show that the explosion is not spherical. with dramatic asymmetries in the structure of some line-forming regions (22)..," In the particular case of SN 2008D, polarization measurements show that the explosion is not spherical, with dramatic asymmetries in the structure of some line-forming regions \citep{Maund:09, Gorosabel:08}."
 Theoretically. current models for the explosior mechanism of core-collapse SNe produce inherently aspherical shock waves (2222222).," Theoretically, current models for the explosion mechanism of core-collapse SNe produce inherently aspherical shock waves \citep{Wheeler:00, Wheeler:02, Blondin:03, Burrows:06, Obergaulinger:06a, Buras:06, Burrows:07}."
 Here we focus οἱ core-collapse SNe driven by bipolar jets (see??).. as may arise from a magneto-rotational mechanism (???)..," Here we focus on core-collapse SNe driven by bipolar jets \citep[see][]{Khokhlov:99, Couch:09}, as may arise from a magneto-rotational mechanism \citep{Wheeler:00, Wheeler:02, Burrows:07}."
 These models have features that may explain many of the observec features of core-collapse supernovae that indicate asymmetry (22222?) The general features of aspherical shock breakout that we discuss. however. apply to arbitrarily aspherical shocks. not just those produced by bipolar jets.," These models have features that may explain many of the observed features of core-collapse supernovae that indicate asymmetry \citep{Khokhlov:99, Wheeler:00, Wang:01, Hoflich:01, Wang:08, Couch:09} The general features of aspherical shock breakout that we discuss, however, apply to arbitrarily aspherical shocks, not just those produced by bipolar jets."
 The absence of spherical symmetry dramatically modifies the observational characteristics of shock breakout and subsequent stages of emission., The absence of spherical symmetry dramatically modifies the observational characteristics of shock breakout and subsequent stages of emission.
 We assume black-body emission in our models and we apply a detector response function appropriate for the XRT and account for ray absorption due to neutral matter along the line of sight so that we can make a direct comparison to the observations of XRO 080109., We assume black-body emission in our models and we apply a detector response function appropriate for the XRT and account for X-ray absorption due to neutral matter along the line of sight so that we can make a direct comparison to the observations of XRO 080109.
 We show that the timescale of the light curve is not set by the light crossing time of the progenitor star but by the sfock-crossing time., We show that the timescale of the light curve is not set by the light crossing time of the progenitor star but by the -crossing time.
 This can account for the length of the XRO associated with SN 2008D with a Wolf-Rayet star progenitor of reasonable parameters., This can account for the length of the XRO associated with SN 2008D with a Wolf-Rayet star progenitor of reasonable parameters.
 Further. we demonstrate that the spectral shape of aspherical shock breakouts is considerably different from that of a single-temperature. spherically-symmetric black-body even if thermal emission is assumed.," Further, we demonstrate that the spectral shape of aspherical shock breakouts is considerably different from that of a single-temperature, spherically-symmetric black-body even if thermal emission is assumed."
 Recently ? have reported on their study of aspherical supernova shock— breakout from blue supergiant progenitors., Recently \citet{Suzuki:10} have reported on their study of aspherical supernova shock breakout from blue supergiant progenitors.
 ? present a semi-analytic method for calculating. shock breakout light curves based on results of 2D hydrodynamie simulations of aspherical core-collapse supernovae., \citet{Suzuki:10} present a semi-analytic method for calculating shock breakout light curves based on results of 2D hydrodynamic simulations of aspherical core-collapse supernovae.
 They assume. as we do in this work. that the breakout emission is thermal and caleulate bolometric breakout light curves.," They assume, as we do in this work, that the breakout emission is thermal and calculate bolometric breakout light curves."
 find that the asphericity of the explosion and the angle from which the explosion is viewed determine the shapes of the resulting light curves., \citet{Suzuki:10} find that the asphericity of the explosion and the angle from which the explosion is viewed determine the shapes of the resulting light curves.
 We find a similar result in this work., We find a similar result in this work.
 Our study. however. Is targeted to explaining the observations of XRO 080109/SN 2008D and. as such. we caleulate X-ray spectra and light curves that allow a direct comparison to the observations.," Our study, however, is targeted to explaining the observations of XRO 080109/SN 2008D and, as such, we calculate X-ray spectra and light curves that allow a direct comparison to the observations."
 Also. our simulations are carried out in a more compact Wolf-Rayet progenitor star.," Also, our simulations are carried out in a more compact Wolf-Rayet progenitor star."
 This paper is organized as follows., This paper is organized as follows.
 In Section 2 we describe the hydrodynamic simulations of jet-driven SNe., In Section \ref{sec:simulations} we describe the hydrodynamic simulations of jet-driven SNe.
 In Section 3 we describe our method of modeling the shock breakout emission from our simulations., In Section \ref{sec:breakout} we describe our method of modeling the shock breakout emission from our simulations.
 In Section 4.. we present the resulting spectra and light curves and compare our spectral and light curve models with the observations of SN 2008D. We discuss our results and give our conclusions in Section 5..," In Section \ref{sec:radModels}, we present the resulting spectra and light curves and compare our spectral and light curve models with the observations of SN 2008D. We discuss our results and give our conclusions in Section \ref{sec:discussion}."
 We have carried out high-resolution hydrodynamiec simulations of jet-driven core-collapse supernovae using the FLASH code. version 2.5 (?)..," We have carried out high-resolution hydrodynamic simulations of jet-driven core-collapse supernovae using the FLASH code, version 2.5 \citep{Fryxell:00}."
 We use two progenitor model stars in our calculations., We use two progenitor model stars in our calculations.
 The first is model sleSa from ?.., The first is model s1c5a from \citet{Woosley:95}.
 This model is a non-rotating. non-magnetic evolved helium star with a pre-supernova radius of 1.35«10! em (1.93 R.) and a pre-supernova mass of about 2.5M.," This model is a non-rotating, non-magnetic evolved helium star with a pre-supernova radius of $1.35\times10^{11}$ cm (1.93 $R_\odot$ ) and a pre-supernova mass of about 2.5."
.. The second model used is a stretched versionof model sleSa., The second model used is a stretched versionof model s1c5a.
" For this model. we stretch the radial coordinates at each model grid point of sleSa according to ria,=Foie6.17«10777."," For this model, we stretch the radial coordinates at each model grid point of s1c5a according to $r_{\rm new} = r_{\rm orig} + 6.17\times10^{-8} r_{\rm orig}^{1.7}$."
 The physical variables. such as density and temperature. at each grid point are left unchanged.," The physical variables, such as density and temperature, at each grid point are left unchanged."
 This function then leaves the core mass and radius approximately unchanged while extending the envelope of the progenitor., This function then leaves the core mass and radius approximately unchanged while extending the envelope of the progenitor.
 The mass and radius of this model are 6.8 παπά 6.51011 em (9.3 Α.Ο). respectively.," The mass and radius of this model are 6.8 and $6.5\times10^{11}$ cm (9.3 $R_\odot$ ), respectively."
 Both progenitors are surrounded by a wind with à mass loss rate of L5«10?M.ye! and wind velocity of 1000 km s., Both progenitors are surrounded by a wind with a mass loss rate of $1.5\times10^{-5}\ M_\sun\ {\rm yr}^{-1}$ and wind velocity of 1000 km $^{-1}$.
 The wind is assumed to be spherically symmetric. which may not be the case for a rotating progenitor.," The wind is assumed to be spherically symmetric, which may not be the case for a rotating progenitor."
 The transition from the progenitor model profile to the wind profile is made linearly over about eight computational zones., The transition from the progenitor model profile to the wind profile is made linearly over about eight computational zones.
 Figure | shows the density profiles of the two progenitor models., Figure \ref{fig:progenitors} shows the density profiles of the two progenitor models.
 We use an equation of state (EoS) that accounts for contributions to the internal energy and pressure from radiation and gas., We use an equation of state (EoS) that accounts for contributions to the internal energy and pressure from radiation and gas.
 In the wind. the gas and radiation are not in thermal equilibrium and so radiation will not contribute to the pressure or internal energy. however our single-temperature code cannot correctly account. for this using an EoS that calculates contributions from. both. radiation and gas.," In the wind, the gas and radiation are not in thermal equilibrium and so radiation will not contribute to the pressure or internal energy, however our single-temperature code cannot correctly account for this using an EoS that calculates contributions from both radiation and gas."
 Therefore. we initially set the temperature in the wind to a small value.," Therefore, we initially set the temperature in the wind to a small value."
 Thisis justified because for a strong shock the upstream temperature is unimportant to the downstream thermodynamics., Thisis justified because for a strong shock the upstream temperature is unimportant to the downstream thermodynamics.
 We track seven atomic species in our simulations: tHe. C. 'O. Ne. Mg. δι. and ?9Ni.," We track seven atomic species in our simulations: $^4$ He, $^{12}$ C, $^{16}$ O, $^{20}$ Ne, $^{24}$ Mg, $^{28}$ Si, and $^{56}$ Ni."
 We employ the consistent multi-fluid advection scheme implemented in FLASH (?).., We employ the consistent multi-fluid advection scheme implemented in FLASH \citep{Plewa:99}.
 We do not include nuclear burning., We do not include nuclear burning.
" Gravity is calculated using the multipole Poisson solver with #720 and (=1,", Gravity is calculated using the multipole Poisson solver with $m=0$ and $\ell=1$.
 All simulations are carried out in 2D cylindrical geometry., All simulations are carried out in 2D cylindrical geometry.
 In order to cover the enormous dynamic range from the inner regions of the progenitor star (~10° em) to the shock radius several minutes after shock breakout (~10 em). we have implemented a logarithmically-spaced cylindrical mesh.," In order to cover the enormous dynamic range from the inner regions of the progenitor star $\sim 10^8$ cm) to the shock radius several minutes after shock breakout $\sim 10^{15}$ cm), we have implemented a logarithmically-spaced cylindrical mesh."
 This is achieved through a radially-dependent maximum level of refinement limiter., This is achieved through a radially-dependent maximum level of refinement limiter.
 This limiter requires that the grid spacing at radius r. Ax. not fall below some fraction of the radius 1.," This limiter requires that the grid spacing at radius $r$, $\Delta x$, not fall below some fraction of the radius $r$ ."
" The grid spacing then takes the form Ax>uN hr. where N, is the number of zones per block in the x-direction and jj is a small number that sets the resolution scale."," The grid spacing then takes the form $\Delta x > \eta N_{x}^{-1} r$ , where $N_x$ is the number of zones per block in the $x$ -direction and $\eta$ is a small number that sets the resolution scale."
 In effect. 4)N7! is analogous to the minimum angular resolution in spherical," In effect, $\eta N_x^{-1}$ is analogous to the minimum angular resolution in spherical"
&eraciuallv. and that the dense cores may survive for longer.,"gradually, and that the dense cores may survive for longer."
 The extent to which disruption would allect these ‘orphan’ galaxies remains unclear., The extent to which disruption would affect these `orphan' galaxies remains unclear.
 Llowever. survival ofthese ‘orphan’ ealaxies (at least to some extent) has been shown hy ?. to be essential in order to explain the observed correlation signal at small scales.," However, survival of these `orphan' galaxies (at least to some extent) has been shown by \cite{bib:Wang06} to be essential in order to explain the observed correlation signal at small scales."
 This study used the Millennium Simulation to construct a new mocdoel of galaxy. clustering., This study used the Millennium Simulation to construct a new model of galaxy clustering.
 They. found that if orphan’ galaxies were exeluded. (rom the analysis. the correlation signal decreases at small scales in contrast to observations.," They found that if `orphan' galaxies were excluded from the analysis, the correlation signal decreases at small scales in contrast to observations."
 A related question is: if we were to increase the numerical resolution of the simulation. would the fraction of Luminous satellites rise significantly?," A related question is: if we were to increase the numerical resolution of the simulation, would the fraction of luminous satellites rise significantly?"
 Clearly. the number of subbaloes (dark. or luminous) must rise further since the subhalo mass funetion roughly follows a power law with dnfdAlxAL So a=LTLO ο 7: Ts 75 7).," Clearly, the number of subhaloes (dark or luminous) must rise further since the subhalo mass function roughly follows a power law with $dn/dM \propto M^{-\alpha}$ , $\alpha = 1.7 - 1.9$ \citealt{bib:Moore99}; ; \citealt{bib:Ghinga00}; \citealt{bib:DeLucia04}; \citealt{bib:Gao04_dm}; \citealt{bib:Diemand07}) )."
 The lowest mass subhaloes we can resolve have circular velocities at. the virial radius. 6. of the order 50 km +: haloes with ο30 km may be inhibited. from star formation by the UV background radiation (e.g. 2: ?:: Tu ?))," The lowest mass subhaloes we can resolve have circular velocities at the virial radius, $v_c$ , of the order $ \lesssim $ 50 km $^{-1}$; haloes with $v_c \lesssim 30$ km $^{-1}$ may be inhibited from star formation by the UV background radiation (e.g. \citealt{bib:Rees86}; \citealt{bib:Efstathiou92}; \citealt{bib:Thoul96}; \citealt{bib:Gnedin00}) )."
 Thus many subhaloes may remain dark. and the fraction of bright subhaloes will not increase significant," Thus many subhaloes may remain dark, and the fraction of bright subhaloes will not increase significantly."
 Increasing the resolution of the simulation would also meanlv. that some of the ‘orphan’ galaxies would be resolved., Increasing the resolution of the simulation would also mean that some of the `orphan' galaxies would be resolved.
 Since only ‘orphan’ galaxies are allowed to merge with the central galaxy. increasing the resolution may prolong the lifetime of some of our ‘orphan’ galaxies.," Since only `orphan' galaxies are allowed to merge with the central galaxy, increasing the resolution may prolong the lifetime of some of our `orphan' galaxies."
 To fully unclerstancl the impact of this ellect a more quantitative analysis is required., To fully understand the impact of this effect a more quantitative analysis is required.
 Clearly. a firm conclusion can only be reached with higher-resolution simulations with realistic treatment of the gas processes.," Clearly, a firm conclusion can only be reached with higher-resolution simulations with realistic treatment of the gas processes."
 In summary. for the CLASS survey. approximately 5. of he 22 primary lensing galaxies appear to have a faint companion within the projected central 5 f+ kpe.," In summary, for the CLASS survey, approximately 5 of the 22 primary lensing galaxies appear to have a faint companion within the projected central 5 $ h^{-1}$ kpc."
 The companions have luminosities of about 240% of the o»imarv galaxy., The companions have luminosities of about $2 - 40$ of the primary galaxy.
 We have studied host galaxies covering a comparable range of luminosities ancl host-to-satellite separations to the CLASS lenses. and found that the oedieted: fraction. of galaxv-(group-) sized. haloes hosting central luminous satellites (~ (6%)) at z=0: (1) at z= 1) is shiehth lower than (but. possibly consistent with) the observed. value.," We have studied host galaxies covering a comparable range of luminosities and host-to-satellite separations to the CLASS lenses, and found that the predicted fraction of galaxy-(group-) sized haloes hosting central luminous satellites $\sim$ ) at $z = 0$; $\sim$ ) at $z = 1$ ) is slightly lower than (but possibly consistent with) the observed value."
 While this fraction is largely independent of galaxy ype. it is shown to increase with redshift.," While this fraction is largely independent of galaxy type, it is shown to increase with redshift."
 The Poisson oobabilitv of detecting Luminous substructure in 5 out of 22 lenses. given that 3% of haloes host luminous substructure. is ~5«10.0.," The Poisson probability of detecting luminous substructure in 5 out of 22 lenses, given that $3\%$ of haloes host luminous substructure, is $\sim 5 \times 10^{-4}$."
 17& of haloes host luminous substructure. the probability of such a detection is ~ 0.14.," If $17\%$ of haloes host luminous substructure, the probability of such a detection is $\sim 0.14$ ."
 Our prediction of the redshift ancl mass dependence appears to be roughly consistent with the cata: three lenses with Luminous satellites are in groups (see Section 3.2)). and all appear have redshifts close to 1. higher than the median redshift (~ 0.6) of all CLASS lenses (see the right panel of Fie. 3)).," Our prediction of the redshift and mass dependence appears to be roughly consistent with the data: three lenses with luminous satellites are in groups (see Section \ref{sec:observations}) ), and all appear have redshifts close to 1, higher than the median redshift $\sim$ 0.6) of all CLASS lenses (see the right panel of Fig. \ref{obs}) )."
 One possibility that we have not considered. is whether lensing galaxies are biased tracers of substructure: such bias may. arise if substructure enhances the lensing cross-sections significantly.," One possibility that we have not considered, is whether lensing galaxies are biased tracers of substructure; such bias may arise if substructure enhances the lensing cross-sections significantly."
 Previous studies. on cluster scales. for giant ares indicates that the bias is small (7)): it remains to be seen whether this holds true for galaxy- lenses.," Previous studies, on cluster scales, for giant arcs indicates that the bias is small \citealt{bib:Hennawi07}) ); it remains to be seen whether this holds true for galaxy-scale lenses."
 Observationally. the Sloan Lens ACS Survey (SLACS) seems to indicate that the lensing galaxies at z ~ 1.2 are typical earlv-tvpe galaxies (2)).," Observationally, the Sloan Lens ACS Survey (SLACS) seems to indicate that the lensing galaxies at z $\sim$ 0.2 are typical early-type galaxies \citealt{bib:Treu08}) )."
 Another possibility is that some of the luminous ‘satellites’ are not associated with lensing galaxies at all. but just happen to be along the ine of sieht (7)).," Another possibility is that some of the luminous `satellites' are not associated with lensing galaxies at all, but just happen to be along the line of sight \citealt{2005ApJ...629..673M}) )."
 ? recently stuclied the clleet of satellite galaxies on gravitational lensing llux ratios using analytic expressions or the host potential and the satellite galaxies., \cite{bib:Shin08} recently studied the effect of satellite galaxies on gravitational lensing flux ratios using analytic expressions for the host potential and the satellite galaxies.
 They use a spherically svnunetric galaxy distribution. ancl assume hat the three-dimensional number density falls olf like comparable to the Milky Way.," They use a spherically symmetric galaxy distribution, and assume that the three-dimensional number density falls off like $r^{-3.5}$, comparable to the Milky Way."
 Thes. show that. the »obabilitv of a finding a Large dwarf is about LOY wo Einstein radii and about within one Einstein radius., They show that the probability of a finding a large dwarf is about within two Einstein radii and about within one Einstein radius.
 We find that our z=0 results are consistent with this., We find that our $z = 0$ results are consistent with this.
 The Millennium Simulation assumes à power-spectrum normalisation of ex=0.9. slightly higher than the latest five-vear result (2)). where ox=0.8.," The Millennium Simulation assumes a power-spectrum normalisation of $\sigma_8 = 0.9$, slightly higher than the latest five-year result \citealt{bib:wmap5}) ), where $\sigma_8 = 0.8$."
 A lower value oes will mean that haloes are expected to form later and be less concentrated., A lower value of $\sigma_8$ will mean that haloes are expected to form later and be less concentrated.
 Llowever. the impact of this parameter on our results is complicated.," However, the impact of this parameter on our results is complicated."
 The semi-analvtie models allow some fine-tuning of parameters to mateh observations., The semi-analytic models allow some fine-tuning of parameters to match observations.
 For example. ο found. no significant dillerence in the galaxy populations (at the redshift range relevant. here) created rom semi-analvtic models based on the one-vear (?)) and. threc-vear (2)) os values of 0.9 and 0.722. »ovided suitable galaxy formation parameters were chosen (the dillerence becomes significant at high redshift).," For example, \cite{bib:Wang08} found no significant difference in the galaxy populations (at the redshift range relevant here) created from semi-analytic models based on the one-year \citealt{bib:wmap1_03}) ) and three-year \citealt{bib:Spergel07}) ) $\sigma_8$ values of 0.9 and 0.722, provided suitable galaxy formation parameters were chosen (the difference becomes significant at high redshift)."
 We finc hat our results do not change significantly when base on the HWALA/P33 galaxy. catalogue produced. by 2? when he same merger timescale is adopted. (as in their moce C)., We find that our results do not change significantly when based on the 3 galaxy catalogue produced by \cite{bib:Wang08} when the same merger timescale is adopted (as in their model C).
 However. in their model B (which has the same star ormation ellicicney but a shorter merger timescale than the ? catalogue) we find a factor of ~ 2 fewer haloes with centra substructure.," However, in their model B (which has the same star formation efficiency but a shorter merger timescale than the \citealt{bib:DeLucia07} catalogue) we find a factor of $\sim$ 2 fewer haloes with central substructure."
 To summarise. while we find that the fraction of uminous satellites in eroup-sized haloes at 2~| is roughly consistent with the observational data we caution hat a firm conclusion can only be reached. with higher-resolution simulations involving a realistic treatment of he gas processes.," To summarise, while we find that the fraction of luminous satellites in group-sized haloes at $ z \sim 1$ is roughly consistent with the observational data we caution that a firm conclusion can only be reached with higher-resolution simulations involving a realistic treatment of the gas processes."
 At the same time. a larger sample of gravitational lenses will also be beneficial to constrain hese models and. allow more celinitive conclusions on the sroperties of the substructure to be made.," At the same time, a larger sample of gravitational lenses will also be beneficial to constrain these models and allow more definitive conclusions on the properties of the substructure to be made."
 We thank lan Browne. Neal Jackson. Liang Cao. Peter Schneider. Dandan Nu and Simon White for useful discussions and. Gerard Lemson and Jie Wang lor providing the galaxy catalogues.," We thank Ian Browne, Neal Jackson, Liang Gao, Peter Schneider, Dandan Xu and Simon White for useful discussions and Gerard Lemson and Jie Wang for providing the galaxy catalogues."
 Phe Millennium Simulation databases used in this paper and the web application providingonline access to them were constructed. às. part of the activities of the German. Astrophysical Virtual Observatory., The Millennium Simulation databases used in this paper and the web application providingonline access to them were constructed as part of the activities of the German Astrophysical Virtual Observatory.
" SEB acknowledges the support. provided: by the EU.Framework 6 Marie.Curie. Early Stage Training Programme under contract number MIZST-C""E-2005-19660 “ESTRELA™."," SEB acknowledges the support provided by the EUFramework 6 MarieCurie Early Stage Training Programme under contract number MEST-CT-2005-19669 “ESTRELA""."
 SAL acknowledges travel support. from. the European Community's Sixth Framework Marie Curie, SM acknowledges travel support from the European Community's Sixth Framework Marie Curie
 , 
In the scenario outlined above. the properties of the stellar cores which eventually merge to give rise to a LGIUD can differ from one another in a number of respects.,"In the scenario outlined above, the properties of the stellar cores which eventually merge to give rise to a LGRB can differ from one another in a number of respects."
 First. the amount of helium envelope remaining at the time the cores merge can vary from several solar masses to very Little., First the amount of helium envelope remaining at the time the cores merge can vary from several solar masses to very little.
 Secondly the mass of both cores can vary as long as enough of the CO core can be accreted to drive the collapse of the ONe core., Secondly the mass of both cores can vary as long as enough of the CO core can be accreted to drive the collapse of the ONe core.
 There is then a range from zero to about a solar mass of CO that can accrete on to the neutron star., There is then a range from zero to about a solar mass of CO that can accrete on to the neutron star.
 Some of this can burn and be ejected in a dise wind to provide the varving quantities of 7?NI seen in the associated supernovae., Some of this can burn and be ejected in a disc wind to provide the varying quantities of $^{56}$ Ni seen in the associated supernovae.
 ‘These range from the very powerful. with as much as half a solar mass of nickel-56. to very weak. with almost no 56 (Mazzalietal.2003:Watsonct2007).," These range from the very powerful, with as much as half a solar mass of nickel-56, to very weak, with almost no nickel-56 \citep{mazzali2003,watson2007}."
. We can also estimate the time required for the initial binary svstem to evolve to the 5-rav burst., We can also estimate the time required for the initial binary system to evolve to the $\gamma$ -ray burst.
 At a metallicity of Z=107. the age of a typical svstem is around. 100 150Myr.," At a metallicity of $Z=10^{-4}$, the age of a typical system is around $100-150\,$ Myr."
 This gives the upper limit to the redshift of the earliest gamma-ray burst of zzz20 anc we should expect to see bursts back to this redshift., This gives the upper limit to the redshift of the earliest gamma-ray burst of $z\approx 20$ and we should expect to see bursts back to this redshift.
 This is a prediction of our moclel., This is a prediction of our model.
 ultrongly magnetic neutron stars (magnetars) have been identified in our galaxy as anomalous X-ray. pulsars(CANIS) or soft 5-rav repeaters (Sls.Alereghetti2008).," Strongly magnetic neutron stars (magnetars) have been identified in our galaxy as anomalous X-ray pulsars (AXPs) or soft $\gamma$ -ray repeaters \citep[SGRs,][]{mereghetti2008}."
.. Phese are observed. to have magnetic fields of Lob 107€ and spin periods of 2.— 10s Known magnetars have ages of 105—103 vr and it is unclear whether they have spun down or were born slowly rotating.," These are observed to have magnetic fields of $10^{14} - 10^{15}\,$ G and spin periods of $2 - 10\,$ s. Known magnetars have ages of $10^3 - 10^4\,$ yr and it is unclear whether they have spun down or were born slowly rotating."
 With estimated. birthrates of 10?vrft they are unlikely to all be related to the LGRBs.," With estimated birthrates of $10^{-3}\,\rm yr^{-1}$ they are unlikely to all be related to the LGRBs."
 Alagnetars could. therefore be of two types. those that are born from single star evolution and those that. form. [rom the merging of the cores of two stars as proposed. here.," Magnetars could therefore be of two types, those that are born from single star evolution and those that form from the merging of the cores of two stars as proposed here."
 As for the radio pulsars. the fields in the first group are likely to be generated by à dynamo mechanism in the stellar core that subsequentIy collapses to form the neutron star.," As for the radio pulsars, the fields in the first group are likely to be generated by a dynamo mechanism in the stellar core that subsequently collapses to form the neutron star."
 Strong magnetic coupling with the stellar envelope leads to outward transport of angular momentum during stellar. evolution and results in a slowly rotating neutron star following core collapse., Strong magnetic coupling with the stellar envelope leads to outward transport of angular momentum during stellar evolution and results in a slowly rotating neutron star following core collapse.
 Magnetars that are born by the merging of the cores of two stars are. on the other hand. likely to be born rapidly spinning with high fields and so give rise to LGRBs.," Magnetars that are born by the merging of the cores of two stars are, on the other hand, likely to be born rapidly spinning with high fields and so give rise to LGRBs."
 The 5-ay jets seen in the LOGRBs have been attributed AID. extraction of the rotational energy. of a neutron star or black hole. or of a massive. about 0.1M... unstable disc that is aceretecl hy the compact star at the time of its birth.," The $\gamma$ -ray jets seen in the LGRBs have been attributed to MHD extraction of the rotational energy of a neutron star or black hole, or of a massive, about $0.1\,\rm M_\odot$, unstable disc that is accreted by the compact star at the time of its birth."
 The combination of super strong magnetic fields ancl millisecond spin periods that appear to be required to explain the energeties and collimation of the 5-ray jet cannot easily be produced in the stellar remnant through single star evolution., The combination of super strong magnetic fields and millisecond spin periods that appear to be required to explain the energetics and collimation of the $\gamma$ -ray jet cannot easily be produced in the stellar remnant through single star evolution.
 Nor does it appear likely that single star evolution could leac to a rapidly spinning disc around the compact star that could generate the field. required to produce the MILD jet., Nor does it appear likely that single star evolution could lead to a rapidly spinning disc around the compact star that could generate the field required to produce the MHD jet.
 We have argued that the required strong fields and rapid spins may occur more naturally in binary star scenarios where merece stellar cores collapse to form neutron stars or rack holes and have presented a convoluted but still simple jnary star origin for 5-rav bursts in which the progenitors are intermediate-mass stars and the collapsed. star is a stronelv magnetic neutron star., We have argued that the required strong fields and rapid spins may occur more naturally in binary star scenarios where merged stellar cores collapse to form neutron stars or black holes and have presented a convoluted but still simple binary star origin for $\gamma$ -ray bursts in which the progenitors are intermediate-mass stars and the collapsed star is a strongly magnetic neutron star.
 This οους radically from »evious proposals that have envisaged massive stars as the »rogenltors., This differs radically from previous proposals that have envisaged massive stars as the progenitors.
" Our model predicts a birth rate ofa few times 10.""vr.+ ovr galaxy at solar metallicitv."," Our model predicts a birth rate of a few times $10^{-5}\,\rm yr^{-1}$ per galaxy at solar metallicity."
 Phe. fraction of binary stars that leac to 5-rav. bursts does not vary much with metallicity in our model., The fraction of binary stars that lead to $\gamma$ -ray bursts does not vary much with metallicity in our model.
 Thus we expect the rate of +-ray oursts to be proportional to the star formation rate for a wide range of redshifts., Thus we expect the rate of $\gamma$ -ray bursts to be proportional to the star formation rate for a wide range of redshifts.
 We expect that the LORBs should wave characteristies of an intermediate-mass population and need. only. be loosely linked to voung star forming regions., We expect that the LGRBs should have characteristics of an intermediate-mass population and need only be loosely linked to young star forming regions.
 tecent high resolution imaging of nearby LOIDs appears o support this requirement., Recent high resolution imaging of nearby LGRBs appears to support this requirement.
 The estimated. time interval tween the time of formation of the component stars and he final collapse to a neutron star is about )Mr so that LORBs produced by this channel should be seen up to red shifts of zzz20.," The estimated time interval between the time of formation of the component stars and the final collapse to a neutron star is about $100\,$ Myr so that LGRBs produced by this channel should be seen up to red shifts of $z\approx 20$."
 CAT thanks Churchill College for his fellowship and the Australian National University for generously supporting his visit to Australia in 2009., CAT thanks Churchill College for his fellowship and the Australian National University for generously supporting his visit to Australia in 2009.
 LUIBL thanks the ANU for hospitality., HHBL thanks the ANU for hospitality.
radius will be within 9 kpc.,radius will be within 9 kpc.
" For comparison, its H--band effective radius is 4.6 kpc."," For comparison, its -band effective radius is 4.6 kpc."
" While it is unlikely that IC 3418’s disk is being tidally stripped, the tail may be tidally stripped from the galaxy as it passes through Virgo’s core, thereby contributing to the ICL."," While it is unlikely that IC 3418's disk is being tidally stripped, the tail may be tidally stripped from the galaxy as it passes through Virgo's core, thereby contributing to the ICL."
 Tails and arcs of young stars and/or H regions which are consistent with an RPS origin have been observed in several clusters., Tails and arcs of young stars and/or H regions which are consistent with an RPS origin have been observed in several clusters.
 CO7 discovered two tails of young stars in deep imaging., C07 discovered two tails of young stars in deep imaging.
" One, an 0.1L* galaxy well within Abell 1689's X-ray halo, should be completely RP stripped, which is consistent with the declining SFR. in the galaxy's disk observed by C07."," One, an $\sim0.1~L^{\ast} $ galaxy well within Abell 1689's X-ray halo, should be completely RP stripped, which is consistent with the declining SFR in the galaxy's disk observed by C07."
" The other, an L* galaxy in Abell 2667, should be partially stripped (C07)."," The other, an $L^{\ast}$ galaxy in Abell 2667, should be partially stripped (C07)."
 Its asymmetric appearance and nuclear burst of star formation match simulations of partially RP stripped disks (Schulz&Struck2001;Kronbergeretal. 2008).," Its asymmetric appearance and nuclear burst of star formation match simulations of partially RP stripped disks \citep{Schulz01, Kronberger08}."
". Using Equation (3)) and the same choice of parameters as C07, we estimate tidal radii beyond the stellar disk for both galaxies."," Using Equation \ref{tidal2}) ) and the same choice of parameters as C07, we estimate tidal radii beyond the stellar disk for both galaxies."
" Yoshidaetal.(2008, discovered tail of star forming knots and young YOS)stellar filaments abehind RB 199 in deep B, Roc, and Ha observations of the Coma cluster."," \citet[][Y08]{Yoshida08} discovered a tail of star forming knots and young stellar filaments behind RB 199 in deep $B$, $R_C$, and $\alpha$ observations of the Coma cluster."
" Y08 estimate that the stellar disk is not tidally stripped, but that the gas disk is partially RP stripped."," Y08 estimate that the stellar disk is not tidally stripped, but that the gas disk is partially RP stripped."
" The galaxy ESO 137-001 in Abell 3627, a very dynamically young nearby cluster, was discovered to have an extended X-ray tail in which multiple H regions were detected (Sunetal.2006,"," The galaxy ESO 137-001 in Abell 3627, a very dynamically young nearby cluster, was discovered to have an extended X-ray tail in which multiple H regions were detected \citep{Sun06, Sun07}."
 Galaxies 2007)..with arcs of H regions following the bow shock morphology of the H have been observed in Virgo (Kenney&Koopmann and Abell 1367 (Gavazzietal.2001)., Galaxies with arcs of H regions following the bow shock morphology of the H have been observed in Virgo \citep{Kenney99} and Abell 1367 \citep{Gavazzi01}.
". The morphology1999) of these arcs suggests cloud stripping, but cloud formation likely occurred in the shock compressed gas while clouds in the outer disk were dissociated and then stripped."," The morphology of these arcs suggests cloud stripping, but cloud formation likely occurred in the shock compressed gas while clouds in the outer disk were dissociated and then stripped."
" Isolated dense clouds can be ablated by the ICM wind (Nulsen1982;Ragaetal. 1998),, dissociated if a shock is driven directly into them and heated by star formation and then stripped (Aannestad(Kenney1973),,etal.2004)."," Isolated dense clouds can be ablated by the ICM wind \citep{Nulsen82, Raga98}, dissociated if a shock is driven directly into them \citep{Aannestad73}, and heated by star formation and then stripped \citep{Kenney04}."
". We advance a basic mechanism to account for the tails’ morphologies, which are characterized by isolated knots, long filaments, and intermediate “streamers” such as that in IC 3418 (d1/k1)."," We advance a basic mechanism to account for the tails' morphologies, which are characterized by isolated knots, long filaments, and intermediate “streamers"" such as that in IC 3418 (d1/k1)."
 Giant molecular clouds may spawn multiple episodes of star formation., Giant molecular clouds may spawn multiple episodes of star formation.
" The cloud experiences a force from the ICM wind, but the newly formed stars do not."," The cloud experiences a force from the ICM wind, but the newly formed stars do not."
" Due to the differing accelerations, the newly formed stars are essentially dropped from the cloud, which itself is supported by the ICM wind."," Due to the differing accelerations, the newly formed stars are essentially dropped from the cloud, which itself is supported by the ICM wind."
 The cloud trails a filament of stars., The cloud trails a filament of stars.
 'The details of the tails! morphologies provide a probe of their formation., The details of the tails' morphologies provide a probe of their formation.
"To first order, a cloud forms a filament of length Where amc is the acceleration due to the(1/2)aycA2,,,,, ICM wind, ajcw, on the molecular cloud and Atgorm is the time between the cloud's first and last episodes of star formation.","To first order, a cloud forms a filament of length $(1/2)a_{\rm MC}\Delta t_{\rm form}^2$, where $a_{\rm MC}$ is the acceleration due to the ICM wind, $a_{\rm ICM}$, on the molecular cloud and $\Delta t_{\rm form}$ is the time between the cloud's first and last episodes of star formation."
" In the cloud's frame, the stars accelerate toward the galaxy with awc."," In the cloud's frame, the stars accelerate toward the galaxy with $a_{\rm MC}$."
 Large amc create long filaments., Large $a_{\rm MC}$ create long filaments.
" Accordingly, the long stellar filaments behind RB 199 are indicative of large ayc, itself characteristic of dynamical coupling between the gas phases and therefore of cloud stripping."," Accordingly, the long stellar filaments behind RB 199 are indicative of large $a_{\rm MC}$, itself characteristic of dynamical coupling between the gas phases and therefore of cloud stripping."
" The ajcm can be estimated as Pram/@gas, or, for the H disk, which can form a 10 kpc filament in 30 Myr."," The $a_{\rm ICM}$ can be estimated as $P_{\rm ram}/\sigma_{\rm gas}$, or, for the H disk, which can form a 10 kpc filament in 30 Myr."
" Assuming molecular clouds that are not dynamically coupled to the HL, ag,/ag=ou,/oH7107% (see Section 4)), corresponding to much shorter “filaments”."," Assuming molecular clouds that are not dynamically coupled to the H, $a_{\rm H_2}/a_{\rm H}=\sigma_{\rm H_2}/\sigma_{\rm H} \approx 10^{-3}$ (see Section \ref{env}) ), corresponding to much shorter “filaments""."
 The treatment above assumes that the cloud and its filament orbit together., The treatment above assumes that the cloud and its filament orbit together.
" In contrast, both tides, which increase the distance between two episodes as they fall, and the ICM wind, which continuously increases the orbital energy of a cloud while it exists, work to lengthen filaments."," In contrast, both tides, which increase the distance between two episodes as they fall, and the ICM wind, which continuously increases the orbital energy of a cloud while it exists, work to lengthen filaments."
" Therefore, if auc=ay, then long filaments are formed, which conflicts with the observations."," Therefore, if $a_{\rm MC}=a_{\rm H}$, then long filaments are formed, which conflicts with the observations."
" We show below that a plethora of morphologies can be formed with ayc=ag,.", We show below that a plethora of morphologies can be formed with $a_{\rm MC}=a_{\rm H_2}$.
" The influences of both gas physics and gravity are observed, requiring us to consider both forces."," The influences of both gas physics and gravity are observed, requiring us to consider both forces."
" The streamer in IC 3418’s tail and the filaments observed by Y08 and C07 are located close to the galaxy, where tidal forces are strongest."," The streamer in IC 3418's tail and the filaments observed by Y08 and C07 are located close to the galaxy, where tidal forces are strongest."
" The extended objects are also redder, indicating that they are older and therefore potentially more dynamically evolved."," The extended objects are also redder, indicating that they are older and therefore potentially more dynamically evolved."
" On the other hand, the streamer in IC 3418 shows a strong UV color gradient and, in RB 199, when a star forming knot is associated with a filament, it is found downwind."," On the other hand, the streamer in IC 3418 shows a strong UV color gradient and, in RB 199, when a star forming knot is associated with a filament, it is found downwind."
 Both highlight the role of the ICM wind., Both highlight the role of the ICM wind.
 Figure 4_ illustrates the next step in modeling complexity., Figure \ref{model} illustrates the next step in modeling complexity.
 We postulate clouds that are generated with a moderate initial velocity heading away from the galaxy., We postulate clouds that are generated with a moderate initial velocity heading away from the galaxy.
 These clouds fall back toward the galaxy., These clouds fall back toward the galaxy.
" Each track with its related line segments represents two episodes of star formation separated by tform=100 Myr, which is motivated by the UV gradient in the streamer, as they fall in an NFW potential."," Each track with its related line segments represents two episodes of star formation separated by $t_{\rm form}=100$ Myr, which is motivated by the UV gradient in the streamer, as they fall in an NFW potential."
" Both episodes have the same initial orbit, but the higher, later forming, episode has an additional acceleration for the first trorm. The line segments represent potential knots and filaments."," Both episodes have the same initial orbit, but the higher, later forming, episode has an additional acceleration for the first $t_{\rm form}$ The line segments represent potential knots and filaments."
" In simulations, RP stripped gas is accelerated to the rest velocity of the ICM, but travels away from the galaxy"," In simulations, RP stripped gas is accelerated to the rest velocity of the ICM, but travels away from the galaxy"
Groups found by the linkage of points generally require a separate procedure to assess statistical significance.,Groups found by the linkage of points generally require a separate procedure to assess statistical significance.
 Grahametal.(1995) considered this problem previously. and ereated the m.σ method. which requires that groups show significant sub-clustering: Pilipenko(2007) addressed it. differently by counting the frequency. of comparable groups. in. random sets.," \citet{Graham1995} considered this problem previously, and created the $m, \sigma$ method, which requires that groups show significant sub-clustering; \citet{Pilipenko2007} addressed it differently by counting the frequency of comparable groups in random sets."
 A procedure is also needed to estimate overdensities., A procedure is also needed to estimate overdensities.
 Pilipenko(2007) estimated. overdensity [rom properties of the AIST. which method is convenient but was not justified in that paper.," \citet{Pilipenko2007}
 estimated overdensity from properties of the MST, which method is convenient but was not justified in that paper."
 In this section we present à new procedure based on the convex hull that we can use for both statistical significance and overdensity., In this section we present a new procedure based on the convex hull that we can use for both statistical significance and overdensity.
 We use a measure of the volume. occupied. by a LOG to assess the statistical significance ancl estimate the overdensity: a LQG must occupy a smaller volume than the expectation for the same number of random points., We use a measure of the volume occupied by a LQG to assess the statistical significance and estimate the overdensity: a LQG must occupy a smaller volume than the expectation for the same number of random points.
 A simple way to estimate the volume is to define a “RA-Dec-z box” with corners determined by the LA. Dec. z limits of the LOG.," A simple way to estimate the volume is to define a “RA-Dec-z box” with corners determined by the RA, Dec, z limits of the LQG."
 Llowever. since LOCs are typically not cuboicds aligned with the coordinate axes. the box typically. overestimates the volume quite severely.," However, since LQGs are typically not cuboids aligned with the coordinate axes, the box typically overestimates the volume quite severely."
 Phe RA-Dec-z boxes will generally. also include non-members and so their volumes are not then solely those of the LOGs., The RA-Dec-z boxes will generally also include non-members and so their volumes are not then solely those of the LQGs.
 Another way to estimate. the volume ds. to use he volume of the convex hull7., Another way to estimate the volume is to use the volume of the convex hull.
 Llowever. the convex hull tvpically underestimates the volume (although overestimates are also possible. depending on morphology) oecause it can wrap tightly. around the surface points and does not then consider any surrounding region as belonging o them.," However, the convex hull typically underestimates the volume (although overestimates are also possible, depending on morphology) because it can wrap tightly around the surface points and does not then consider any surrounding region as belonging to them."
 We can construct a measure of volume that should οι(ος reflect the volume occupied by the LQG than either he RA-Dece-box or the convex hull of the points., We can construct a measure of volume that should better reflect the volume occupied by the LQG than either the RA-Dec-box or the convex hull of the points.
 We do this w expanding each member point of a unit to be a sphere. with radius set to be half of the mean linkage (AIST edge ength) of the unit.," We do this by expanding each member point of a unit to be a sphere, with radius set to be half of the mean linkage (MST edge length) of the unit."
 In this way. each point is associated with a spherical volume.," In this way, each point is associated with a spherical volume."
 We then take the volume of the LOG o be the volume of the convex hull of these spheres., We then take the volume of the LQG to be the volume of the convex hull of these spheres.
 We shall refer to this method as the CLIALS metho: COLVCX. ull of member spheres., We shall refer to this method as the CHMS method — convex hull of member spheres.
 Note that this measure refers to he LQG members only. unlike the ItX-Dec-z box. which vpically incorporates non-members too.," Note that this measure refers to the LQG members only, unlike the RA-Dec-z box, which typically incorporates non-members too."
 The distribution. of CLIAIS volumes resulting. [roni random points that have been distributed. with known density allows the statistical significance of a LOG to be assessed., The distribution of CHMS volumes resulting from random points that have been distributed with known density allows the statistical significance of a LQG to be assessed.
 The CLALS volumes for sets of random points can, The CHMS volumes for sets of random points can
have amplitude 1056 and to oceur ou timescales 9/<3.3hours=0.01/ to explain the March 1:61 data. aud 9/~7ininutes=0.003/ to explain tle March 3.21 data.,"have amplitude $\sim 10\%$ and to occur on timescales $\delta t \la 3.3
\,\hbox{hours} = 0.04 t$ to explain the March 4.64 data, and $\delta t
\sim 7 \, \hbox{minutes} = 0.003 t$ to explain the March 3.21 data."
 The latter in particular is physically implatsible for an alterglow expaucline 1to a realistic external imedium., The latter in particular is physically implausible for an afterglow expanding into a realistic external medium.
 Finally. if the error bars on the photometry are systematically 1(erestimated. the data may be recouciled with Lic» color evolutio1," Finally, if the error bars on the photometry are systematically underestimated, the data may be reconciled with no color evolution."
" In order to reduce477 to 5- ( Ιvat X7Dy/dofJ=1). we need to add a systematic"" error contributio1 of 0.015 magnitudes in quadratL‘eto all of the A’ and interpolated 2 fluxes."," In order to reduce$\chi^2$ to $5$ (so that $\chi^2 / \hbox{dof} = 1$ ), we need to add a systematic error contribution of $0.078$ magnitudes in quadrature to all of the $K'$ and interpolated $R$ fluxes."
 Given the discussion of systematic errors above al( our effort to adjust all Ro band fluxes to the Heucden et al caliation. this large an additional eror is unlikely.," Given the discussion of systematic errors above and our effort to adjust all R band fluxes to the Henden et al calibration, this large an additional error is unlikely."
 We therefore believe that at least some of the measured color variatiou is real., We therefore believe that at least some of the measured color variation is real.
 The loug spect‘al baseline between the I baud (2.1pini) αιd the optical/UV bands. allows oue to obtain accurate spectral slope nieasuremeuts., The long spectral baseline between the K' band $2.1 \micron$ ) and the optical/UV bands allows one to obtain accurate spectral slope measurements.
 We lave calculated the burst spectral energy (istributiou at selected times based on the availability of multiwavelength data., We have calculated the burst spectral energy distribution at selected times based on the availability of multiwavelength data.
 Where necessary. {X lueasuremenuts were interpolated between adjacent data poius at one waveleneth in order to determine a «'ontemporaueous flux with another waveleneth.," Where necessary, flux measurements were interpolated between adjacent data points at one wavelength in order to determine a contemporaneous flux with another wavelength."
 For this operation. we always interpolated the light ¢rves with good sampling (BR. Iv) to matcl the time of a sparsely sauupled Wiweleneth (UV 28254. B).," For this operation, we always interpolated the light curves with good sampling (R, K') to match the time of a sparsely sampled wavelength (UV $2825
\mbox{\AA}$, B)."
 Photometric zero points for the conversion of magnitudes to [lux deusity ]its were taken from Fukugita. Shimasaku. Ichikawa (1995) for the optical. aud from Caumpius. Rieke. Lebofsky (1985) for the near-IR.," Photometric zero points for the conversion of magnitudes to flux density units were taken from Fukugita, Shimasaku, Ichikawa (1995) for the optical, and from Campins, Rieke, Lebofsky (1985) for the near-IR."
 The resulting specαἱ energy. distributious are shown in figure 3.., The resulting spectral energy distributions are shown in figure \ref{sedfig}.
 The best-fit spectral slopes assunmiug uubrokeu power law spectra are given in table : .. first uncorrected for extiuction: then with a correction for foreground Galactic extinction (Eg;\=0.053. Schlegel et al 1998) assuming an Ry=AyΕυy3.1 extinetion law: and finally with au additional correction for possible extinction in the CRB lost galaxy (see below).," The best-fit spectral slopes assuming unbroken power law spectra are given in table \ref{slopetable}, , first uncorrected for extinction; then with a correction for foreground Galactic extinction $E_{B-V} = 0.053$, Schlegel et al 1998) assuming an $R_V \equiv A_V / E_{B-V} = 3.1$ extinction law; and finally with an additional correction for possible extinction in the GRB host galaxy (see below)."
 Also included are the slopes derived [rom Ro aud IN data alone. correcec only for Galactic extinction.," Also included are the slopes derived from R and K' data alone, corrected only for Galactic extinction."
 The broad baud spectral energy distributions plotted in ligure 3. show uo compelliug evideuce for extinction al the redshift of the GRB. (, The broad band spectral energy distributions plotted in figure \ref{sedfig} show no compelling evidence for extinction at the redshift of the GRB. (
This is a change from the first preprint of this paper. largely due to revised calibration of the March 6.375 UV data point at 30504.),"This is a change from the first preprint of this paper, largely due to revised calibration of the March 6.375 UV data point at $3050 \mbox{\AA}$ .)"
 However. a modest amount of host gaaxy extinction retains consisteut with the data.," However, a modest amount of host galaxy extinction remains consistent with the data."
 The 21754 dlust. absorption eature falls into the observed Ro baud at the redshift 2=2.03) of this burst., The $2175$ dust absorption feature falls into the observed R band at the redshift $z=2.03$ of this burst.
" Assuming tha there is uo intrinsic eiulssion feature at this wavelength in the GRB alterelow spectrum. we can place an upper limit of .A,-<0.1 magnitude for Milky Way type cust at z=2.03."," Assuming that there is no intrinsic emission feature at this wavelength in the GRB afterglow spectrum, we can place an upper limit of $A_V \la 0.1$ magnitude for Milky Way type dust at $z=2.03$."
 Dust with a different recdeniug law is less strongly constrained., Dust with a different reddening law is less strongly constrained.
 A Simall Magellanic Cloud reddening law is plausible for Ay<0.12. while a Large Magellanic Cloud (LMLC) law is plausible for Ay 0.2.," A Small Magellanic Cloud reddening law is plausible for $A_V \la 0.12$, while a Large Magellanic Cloud (LMC) law is plausible for $A_V \la 0.2$ ."
 The measure of plausibility here isthat a single power law should approximately fit the spectrum at, The measure of plausibility here isthat a single power law should approximately fit the spectrum at
optimization of GAPS.,optimization of GAPS.
 We discuss the relevaut issues iu the next section., We discuss the relevant issues in the next section.
 Ouce an antiparticle is slowed down to au eunerev of order the ionization enerev of the atom. the autiparticle is captured iuto au exotic ato replacing a bound electron.," Once an antiparticle is slowed down to an energy of order the ionization energy of the atom, the antiparticle is captured into an exotic atom replacing a bound electron."
 The cross section for trapping of Nis of the order of a molecular cross-section (Becketal.1993)., The cross section for trapping of is of the order of a molecular cross-section \citep{beck93}.
".. The captured antiparticle is in a highlv excited state. my=GM/m,yh? (asstnune that replaces a I&-hell electron) due to energy. conservation (Wavanoctal.199 1).."," The captured antiparticle is in a highly excited state, $n_0=(M^*/m_e)^{1/2}$ (assuming that replaces a K-shell electron) due to energy conservation \citep{hayano94}. ."
 For instance. vy10 for autiproton.," For instance, $n_0\sim40$ for antiproton."
 AL” is the reduced mass of the N-atom svstem., $M^*$ is the reduced mass of the $\bar{X}$ -atom system.
 decays via a radiative transition or an Auger process., decays via a radiative transition or an Auger process.
 A radiative ransition cents a photon from the exotic atom. while he Auger process ionizes a bound electron.," A radiative transition emits a photon from the exotic atom, while the Auger process ionizes a bound electron."
 The Auger Xocess sets in when the transition energy of exceeds the ionization energv of bound clectrous., The Auger process sets in when the transition energy of exceeds the ionization energy of bound electrons.
 Usually. the Auger Xocess is fast compared to the radiative transitions.," Usually, the Auger process is fast compared to the radiative transitions."
" Siuce arecr changes in An. and A?=£1 are preferred for he deexcitation (dipole selection rule). will drop to he so-called circular state iu which the orbital angular uonientuni fakes its largest value bv the transition QI)o>teτμLü.,-n2)."," Since larger changes in $\Delta n$, and $\Delta l=\pm 1$ are preferred for the deexcitation (dipole selection rule), will drop to the so-called circular state in which the orbital angular momentum takes its largest value by the transition $(n,l)\rightarrow(n'=n-1,l'=n-2)$."
 Althonel the Auger xocess may lead to other than the circular states with siidl branching ratio (ILutniaun1990).. will be in a circular state at an carly stage in its decavius process.," Although the Auger process may lead to other than the circular states with small branching ratio \citep{hartmann90}, will be in a circular state at an early stage in its decaying process."
 Decay of circular. states proceeds by An=LAJ1l due to the selection rules for ladder transitions.," Decay of circular states proceeds by $\Delta n=-1, \Delta
l=-1$ due to the selection rules for ladder transitions."
" A simplified. version of the decay process is: The photon cnergy in ladder transition »)ΣΕ ls eiven by. where A7 is the reduced mass of in units of proton lüass my ud jg=i, fie."," A simplified version of the decay process is: The photon energy in ladder transition $n \rightarrow n-1$ is given by, where $M^*$ is the reduced mass of in units of proton mass $m_p$ and $\eta=m_p/m_e$ ."
 + aud Z are the charge of the incident particle aud the absorbing atom aud we defined Z=iZ for convenience., $z$ and $Z$ are the charge of the incident particle and the absorbing atom and we defined $\tilde{Z}=zZ$ for convenience.
 Ladder transition energies are listed in table 1- 3.., Ladder transition energies are listed in table \ref{e_pbar}- \ref{e_hebar}.
" The transition rate. ων) 1|. for ladder transition (9.»9» l)is given by. Therefore. the transition time is: When replaces an electron in shell Ve, GN,=1.2.3 corresponds to IK. L. M-xhiell respectively). its initial bouud state is ny~Ne,GM)7,"," The transition rate, $\Gamma_{ladder}(n)$ $^{-1}$ ], for ladder transition $n\rightarrow n-1$ ) is given by, Therefore, the transition time is: When replaces an electron in shell $N_{sh}$ $N_{sh}=1, 2, 3$ corresponds to K, L, M-shell respectively), its initial bound state is $n_0
\sim N_{sh}(\eta M^*)^{1/2}$."
 The lifetime Ty Is given by. This is an upper lait to the lifetime which assumes that all the transitions are via radiative deexcitation.," The lifetime $\tau_{\bar{X}}$ is given by, This is an upper limit to the lifetime which assumes that all the transitions are via radiative deexcitation."
 Delay of the annuihilation of antiprotous iu helium was observed by Νακάνταctal.(1990.. indicating the formation of autiprotonic helimu atoms with lifetime of order 10. sec.," Delay of the annihilation of antiprotons in helium was observed by \citet{nakamura94}, indicating the formation of antiprotonic helium atoms with lifetime of order $10^{-6}$ sec."
 When the transition energv of becomes larger thau the Is-shell ionization cuerey of a bound electron. the Απο process jonizes a K-shell electron.," When the transition energy of becomes larger than the K-shell ionization energy of a bound electron, the Auger process ionizes a K-shell electron."
 The quantum munber vy where electrons are completely depleted is eiven by. Iu this case. ng is independent of Z.," The quantum number $n_K$ where electrons are completely depleted is given by, In this case, $n_K$ is independent of $Z$ ."
 yg is 15.19.38 forp. aud vespectively.," $n_K$ is $\sim 15, 19, 38$ for, and respectively."
 Bacherctal.(1988) observed strong suppression of 7=15.16 ladder N-ravs in several noble eases iluumated by autiproton beams.," \citet{bacher88} observed strong suppression of $n=15,16$ ladder X-rays in several noble gases illuminated by antiproton beams."
 All the transitions from i<ng are radiative aud ladder photons from lower 0 (e.g. 90=Lo7 for p) are observed by X-ray detectors., All the transitions from $n<n_K$ are radiative and ladder photons from lower $n$ (e.g. $n=4-7$ for $\bar{p}$ ) are observed by X-ray detectors.
 Tn order to increase cfiiciency for detecting VY. high density targets are desirable for the detection medi.," In order to increase efficiency for detecting $\bar{X}$, high density targets are desirable for the detection medium."
 As the density increases. several atomic processes may interfere with the identification of characteristic N-ravs from the ladder trausitious.," As the density increases, several atomic processes may interfere with the identification of characteristic X-rays from the ladder transitions."
 Figure 6 schematically slows the deexcitation path of an exotic atom and figure presents the rates of differcut atomic transitions discussed in the following section., Figure \ref{decay_path} schematically shows the deexcitation path of an exotic atom and figure \ref{rate} presents the rates of different atomic transitions discussed in the following section.
 A quautumn state of becomes degenerate after the full ionization of bound , A quantum state of becomes degenerate after the full ionization of bound electrons.
Then the electric field from an adjacent atom Eelectrons.=(e/R). distorts differcut | quantum states. inducing An=0 transitions.," Then the electric field from an adjacent atom $\vec{E}=(e/R^3)\vec{R}$, distorts different $l$ quantum states, inducing $\Delta{n}=0$ transitions."
 is the intermolecular distance or iapact parameter of apassing atom iu the eas., $R$ is the intermolecular distance or impact parameter of apassing atom in the gas.
 This transition. leads to anu S-xtatefollowed by a unelear anuihilation or to an nS»1S radiative transition., This transition leads to an $S$ -statefollowed by a nuclear annihilation or to an $nS\rightarrow 1S$ radiative transition.
 The photon cucrey from the 0S.>15 radiativetransition in high Z materials is too high to be measured bv a thin N-rav detector., The photon energy from the $nS\rightarrow 1S$ radiativetransition in high $Z$ materials is too high to be measured by a thin X-ray detector.
 A measure of the, A measure of the
The leavy-clement abundance iu galaxies ix a kev physical property for understandiug galaxy evolution.,The heavy-element abundance in galaxies is a key physical property for understanding galaxy evolution.
 Metals axe formed in massive stars and are dispersed mto the iuterstellar medium (ISAT) via miass-loss processes., Metals are formed in massive stars and are dispersed into the interstellar medium (ISM) via mass-loss processes.
 Hence. metallicity provides an important record of the star formation history of a galaxy.," Hence, metallicity provides an important record of the star formation history of a galaxy."
 Galaxies aud their chemical abundances. however. do not evolve as closed systems where iaetals are dispersed iuto the ISM as eas is converted into stars;," Galaxies and their chemical abundances, however, do not evolve as closed systems where metals are dispersed into the ISM as gas is converted into stars."
 Rather. they are modulated by inflow of uneuriched eas (c.e..7) and a complex “feedback” anechanisia where supernovae and stellar winds influence the ISAL by inducing outfows of gas iuto the intergalactic medium (ICAL) and regulatiug star ormation through reheating (72???7)..," Rather, they are modulated by inflow of unenriched gas \citep[e.g.,][]{Kewley2010} and a complex ""feedback"" mechanism where supernovae and stellar winds influence the ISM by inducing outflows of gas into the intergalactic medium (IGM) and regulating star formation through reheating \citep{Larson1974, Larson1975, Kauffmann1998, Somerville1999, Springel2003}."
" A simple “closed-box” imodel in which feedback is nof considered predicts a correlation between two Mudamental mass and metallicity (22??)., "," A simple ""closed-box"" model in which feedback is not considered predicts a correlation between two fundamental $-$ mass and metallicity \citep{vandenbergh1962, Schmidt1963, Searle1972, Erb2006}."
The pioneering work of ? fist showed this correlation in local mregulu galaxies., The pioneering work of \citet{Lequeux1979} first showed this correlation in local irregular galaxies.
" May subsequent cfforts ive focused ou the relation between Imuuimositv aud uctallicity (227277). where Iuninosity is taken as à proxy Or nass because of the difücultv in iuferriug ealaxy uasses fron observables,"," Many subsequent efforts have focused on the relation between luminosity and metallicity \citep{Skillman1989, Zaritsky1994, Garnett1997, Lamareille2004, Salzer2005} where luminosity is taken as a proxy for mass because of the difficulty in inferring galaxy masses from observables."
 Early attempts το extend lis relation to iutermiediate redshifts found that the uninositvauctalliaty (LZ) relation at earlier epochs was consistent with the local universe (27)...," Early attempts to extend this relation to intermediate redshifts found that the luminosity-metallicity (LZ) relation at earlier epochs was consistent with the local universe \citep{Kobulnicky1999, Carollo2001}."
 However. many recent investieations have provided compclling evideuce or evolution of the LZ relation over cosmological times C2222277) ," However, many recent investigations have provided compelling evidence for evolution of the LZ relation over cosmological times \citep{Kobulnicky2000, Pettini2001, Kobulnicky2003a, Kobulnicky2004, Maier2004, Shapley2004, Shapley2005}."
Lunadnosifies of galaxies evolve along with heir chemical abundances on cosmological timescales and diseutausling the contribution for each has posed difficultics iu these studies., Luminosities of galaxies evolve along with their chemical abundances on cosmological timescales and disentangling the contribution for each has posed difficulties in these studies.
 Evolution iu the LZ relatiou nuplies au evolution iu the amore fuucdinueutal miassanuetalliitv. (MZ). relation., Evolution in the LZ relation implies an evolution in the more fundamental mass-metallicity (MZ) relation.
 Development of more sophisticated models for stellar population svuthesis (7)— and gaseous uebula (77) allowed ονhereafterTOL to establish aud quantity the local MZ relation for ~523.000 star-forming galaxies from the Sloan. Digital Skv Survey (SDSS).," Development of more sophisticated models for stellar population synthesis \citep{Bruzual2003} and gaseous nebula \citep{Ferland1996, Charlot2001} allowed \citet[hereafter T04]{Tremonti2004} to establish and quantify the local MZ relation for $\sim\!\!53,000$ star-forming galaxies from the Sloan Digital Sky Survey (SDSS)."
" They fouud that metallicity increases linearly with stellar uass for galaxies Laving masses between LOSTAL. and 1019?ay, and flattens out at higher masses.", They found that metallicity increases linearly with stellar mass for galaxies having masses between $10^{8.5} M_{\odot}$ and $10^{10.5} M_{\odot}$ and flattens out at higher masses.
 They attribute this depletion of iietals iu less massive galaxies to ubiquitous ealactic winds that strip metals more effectivele frou ealaxies with shallow potential wells. dispersing them iuto the IGAL," They attribute this depletion of metals in less massive galaxies to ubiquitous galactic winds that strip metals more effectively from galaxies with shallow potential wells, dispersing them into the IGM."
 Alternatively. it has been suggested that low mass svstenis have ongoing star formation and have vot to couvert much of their gas into stars aud therefore are less evolved. less metalaiclhi systems as compared to more luassive galaxies which have uudergone rapid star formation (777).," Alternatively, it has been suggested that low mass systems have ongoing star formation and have yet to convert much of their gas into stars and therefore are less evolved, less metal-rich systems as compared to more massive galaxies which have undergone rapid star formation \citep{Brooks2007, DeRossi2007, Mouchine2008}."
 Finallh. a varving galaxy-iutegrated initial mass function has also emerged as a possible explanation for the observed MZ relation (?)..," Finally, a varying galaxy-integrated initial mass function has also emerged as a possible explanation for the observed MZ relation \citep{Koppen2007}."
 Cosmological νάνοςναλος simulations (77) lave predicted. an evolution iu the MZ relation aud receut observationsat intermediate redshifts (2~ 0.7) have provided strong evidence for an evolution.," Cosmological hydrodynamic simulations \citep{Brooks2007, Oppenheimer2008} have predicted an evolution in the MZ relation and recent observationsat intermediate redshifts $z\sim0.7$ ) have provided strong evidence for an evolution."
 Using the Gemini Deep Deep Survey. ?/— first showed the MZ relation bevoud the local universe for galaxies at τ=OL. LO. ," Using the Gemini Deep Deep Survey, \citet{Savaglio2005} first showed the MZ relation beyond the local universe for galaxies at $z = 0.4-1.0$ ."
"Their relation iutersects the local relation at AL~LOMAT, but is ~0.3 dex lower at AL~LOOAT... imiplàug a much steeper slope."," Their relation intersects the local relation at $M \sim 10^{10.4}M_{\odot}$ but is $\sim0.3$ dex lower at $M \sim 10^{9.3}M_{\odot}$, implying a much steeper slope."
 They conclude that the slope of the MZ relation must become flatter over time., They conclude that the slope of the MZ relation must become flatter over time.
 ? found a difference 70.2 dex for a redshift range of 0.175«i0.9 with a simular slope., \citet{Cowie2008} found a difference $>0.2$ dex for a redshift range of $0.475<z<0.9$ with a similar slope.
 More receutly. find a flatter slope iu the MZ relation out to :—0.9 and an average difference in imoetalliitv of ~0.2 dex.," More recently, \citet{Lamareille2009} find a flatter slope in the MZ relation out to $z \! \sim \! 0.9$ and an average difference in metallicity of $\sim\!0.2$ dex."
 To date. the MZ relation has been extended bevoud οσον. ," To date, the MZ relation has been extended beyond $z\!\sim\!3$ \citep{Shapley2005, Erb2006, Maiolino2008,Perez-Montero2009,Mannucci2009}. ."
ΑΙ these studies have shown evideuce, All these studies have shown evidence
We show the ΕΟΝ and parameter errors predieted. for different surveys individually and combined in Table &..,We show the FoM and parameter errors predicted for different surveys individually and combined in Table \ref{surveyFoM:table}.
 We performed an analysis of the optimal dark energy survey that could be carried. out in a given time period by an experiment similar in design to WEALOS., We performed an analysis of the optimal dark energy survey that could be carried out in a given time period by an experiment similar in design to WFMOS.
 We estimate the accuracy of the distance measurements that utilize only the oscillatory part of the power spectrum (the wigeles using the fitting formula of Seo Eisenstein (2007).," We estimate the accuracy of the distance measurements that utilize only the oscillatory part of the power spectrum (the 'wiggles'), using the fitting formula of Seo Eisenstein (2007)."
 Our measure). of utility. or FoM. was defined. to be proportional to the area of the error ellipse for the CPL parameters of the dark energy equation of state.," Our measure of utility, or FoM, was defined to be proportional to the area of the error ellipse for the CPL parameters of the dark energy equation of state."
 Our results can be summarized as follows:, Our results can be summarized as follows:
where the eas density strictly increases with lookback time.,where the gas density strictly increases with lookback time.
 The trianele points with dotted errors in Figure 20 show the results of this reconstruction., The triangle points with dotted errors in Figure \ref{schmidt} show the results of this reconstruction.
 These points. including the one at 0.9 Car. are all consistent with the known correlation between eas deusitv and star formation rate.," These points, including the one at 0.9 Gyr, are all consistent with the known correlation between gas density and star formation rate."
 This consistency allows the possibility that eas was recently accreted. as was suggested by based on the warped nature of the outer eas disk.," This consistency allows the possibility that gas was recently accreted, as was suggested by based on the warped nature of the outer gas disk."
" ""Together. these results favor a scenario where a gas disk was in place for many Cyr."," Together, these results favor a scenario where a gas disk was in place for many Gyr."
" This disk was of low surface deusity (<5 M. 2) for most of the historv of the disk and formed stars at a low rate (510 τν t 3], ", This disk was of low surface density $<$ 5 $_{\odot}$ $^{-2}$ ) for most of the history of the disk and formed stars at a low rate $<$ $^{-3}$ $_{\odot}$ $^{-1}~$ $^{-2}$ ).
This gas may have been completely depleted by ~1 Gar ago. when new eas eutered the svstem. either from sumroundius eas filaments. or frou a stall. eas-vich. dwarf galaxy that merged to produce the current gas disk.," This gas may have been completely depleted by $\sim$ 1 Gyr ago, when new gas entered the system, either from surrounding gas filaments, or from a small, gas-rich, dwarf galaxy that merged to produce the current gas disk."
 This eveut triggered a small amount of new star formation over the past ~500 Myr., This event triggered a small amount of new star formation over the past $\sim$ 500 Myr.
 Tt is still possible that the preseut eas disk is a τος of the eas that originally formed the disk aud that there was uo recent mereer., It is still possible that the present gas disk is a relic of the gas that originally formed the disk and that there was no recent merger.
 Tf so. the low-density gas disk must have formed stars in short. weal. episodes that lasted ~200LOO Myr.," If so, the low-density gas disk must have formed stars in short, weak, episodes that lasted $\sim$ 200–400 Myr."
" Such episodes. like the oue we have resolved in the most recent Cyr. form ouly Z 107 ML, of stellar mass in the immer disk (r«5 kpc. assuniünug axisviunietry). which could reconcile the amount of stars formed in the past Car with the amount of gas available in our closed-box assumption."," Such episodes, like the one we have resolved in the most recent Gyr, form only $\lap$ $^7$ $_{\odot}$ of stellar mass in the inner disk $r\,<\,5$ kpc, assuming axisymmetry), which could reconcile the amount of stars formed in the past Gyr with the amount of gas available in our closed-box assumption."
 Such star formation would be insienificaut compared to the 2 «10? NE. of stellar mass formed in the mner disk prior to 10 Cir ago and the 1.5105 AL. of1 present2001)., Such star formation would be insignificant compared to the $\gap$ $\times$ $^9$ $_{\odot}$ of stellar mass formed in the inner disk prior to 10 Gyr ago and the $\times$ $^8$ $_{\odot}$ of present.
. Such a scenario is favored by the isolated cuviromment of NGC lot but is not consistent with the star formation rates we mcasure from 0.61.0 Cyr., Such a scenario is favored by the isolated environment of NGC 404 but is not consistent with the star formation rates we measure from 0.6–1.0 Gyr.
 Furthermore. sugeest the warped nature of the gas disk as additional evidence for the recent merecr.," Furthermore, suggest the warped nature of the gas disk as additional evidence for the recent merger."
 While the warp is consistent with the merecr scenario. warping can also occur froin iuisaligunieut of angular moment duriug late eas accretion from filaments and therefore does not require a catastrophic merger.," While the warp is consistent with the merger scenario, warping can also occur from misalignment of angular momentum during late gas accretion from filaments and therefore does not require a catastrophic merger."
 We have measured resolved stellar photometry for ~15 ⋅⋅≻⋟ ⋅ ⋅ ⋜⋯⊳∐∐∐−∪↕↑∐↸∖∐∐∐∖↥⋅≼∐↴∖↴↘↽⋜⋯≼↧∿↓∩⋜⋯⊳∐∐∐−∪↕↑∐↸∖∪∏↑↸∖↥⋅↽ ⋅⋅≻⋟ disk of the S0 ealaxy NGC. 104., We have measured resolved stellar photometry for $\sim$ 15 $^2$ of the inner disk and $\sim$ 10 $^2$ of the outer disk of the S0 galaxy NGC 404.
 Detailed fittiug of the resulting CMDs shows that the disk is dominated bv stars older than το Cir. with all of our inner disk data coustraiine the aee of ~90% of the stellar mass to 28 Cor and our best data constrainiug the age of of the stellar mass to 710 Gyr.," Detailed fitting of the resulting CMDs shows that the disk is dominated by stars older than $\sim$ 10 Gyr, with all of our inner disk data constraining the age of $\sim$ of the stellar mass to $>$ 8 Gyr and our best data constraining the age of of the stellar mass to $>$ 10 Gyr."
 These results show that the difference between this disk aud most later-type disks that have been studied in detail is seen in the οἱ stellar population as well as the young., These results show that the difference between this disk and most later-type disks that have been studied in detail is seen in the old stellar population as well as the young.
 We fouud no treuds between the age distributions « the vouug or old populations and distance from the ealaxy center within the observed regions., We found no trends between the age distributions of the young or old populations and distance from the galaxy center within the observed regions.
 However. as seen by the mean ROB color anc ACB/RGB ratio of the innermost regions are cousisteut with a bulee population in NGC 101 that is vounger than he disk.," However, as seen by the mean RGB color and AGB/RGB ratio of the innermost regions are consistent with a bulge population in NGC 404 that is younger than the disk."
" Despite its low gas density. NGC LOL appears to have con. formüug stars throughout its historv on timescales of Gr. even in the inner disk (~1 1.5"")."," Despite its low gas density, NGC 404 appears to have been forming stars throughout its history on timescales of Gyr, even in the inner disk $\sim$ $'$ )."
 The SEIT aud he current density of the eas disk are cousisteut with NGC 101 having evolved roughly iu a closed box over uuch of its historv. starting as a laree disk of high surface density aud approxinatelv following the Schmidt aw as its overall star formation rate and eas density decreased.," The SFH and the current density of the gas disk are consistent with NGC 404 having evolved roughly in a closed box over much of its history, starting as a large disk of high surface density and approximately following the Schmidt law as its overall star formation rate and gas density decreased."
 Such loue-lived disks have been observed iu isolated galaxies before. as in the compact dwarf ADDS 113815|20082009).. whose gas disk is also of low density.," Such long-lived disks have been observed in isolated galaxies before, as in the compact dwarf ADBS 113845+2008, whose gas disk is also of low density."
 Perhaps these low deusity disks have very oue lifetimes due to very low star formation efficiencies at these densities., Perhaps these low density disks have very long lifetimes due to very low star formation efficiencies at these densities.
 Our ΕΠΣ indicate that star formation in the low deusitv relic disk may be episodic. geome hrough minor episodes of star formation that last a few mudred Myr every Cyr or so.," Our SFHs indicate that star formation in the low density relic disk may be episodic, going through minor episodes of star formation that last a few hundred Myr every Gyr or so."
 Ou the other haud. in NGC LOL there is some evidence hat this passive evolution was slightly disrupted bv an event 0.G-0.9 Car ago. possibly a small merecr.," On the other hand, in NGC 404 there is some evidence that this passive evolution was slightly disrupted by an event 0.6-0.9 Gyr ago, possibly a small merger."
" Such a nerecr event would alleviate the necessity for very low star formation efficiency 71 Cyr ago by allowing that the curreutlv-observed eas disk was not fully iu place at that Hue,", Such a merger event would alleviate the necessity for very low star formation efficiency $\sim$ 1 Gyr ago by allowing that the currently-observed gas disk was not fully in place at that time.
 Support for this work was provided by NASA hrough erauts GO-LO9L5 aud GO-11719 from the Space Telescope Science. Tustitute. which is operated by the Association of Universities for Researcli in Astrouoniv. Iucorporated. under NASA contract NÀS5-26555.," Support for this work was provided by NASA through grants GO-10915 and GO-11719 from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Incorporated, under NASA contract NAS5-26555."
explanation for why twin QPOs are seen in both neutron-star and black-hole svsteuis.,explanation for why twin QPOs are seen in both neutron-star and black-hole systems.
 llowever. the fact that the Irequency. difference between (win kIHlz QPOs in neutron star svstems is often of order the neutron star frequency has no simple explanation.," However, the fact that the frequency difference between twin kHz QPOs in neutron star systems is often of order the neutron star frequency has no simple explanation."
 Also. since the model explicitly invokes general relativistic elfects. it cannot explain the continuitv in QPO properties between neutron stars and black holes on the one hand and while clwarls on the other (Warner οἱ al.," Also, since the model explicitly invokes general relativistic effects, it cannot explain the continuity in QPO properties between neutron stars and black holes on the one hand and white dwarfs on the other (Warner et al."
 2003)., 2003).
 A munber of scientists have investigated modes of oscillation of an accretion disk as a model of high frequency. QPOs (forreviewsseeKato.Fukue.&Mineshige1908:Wagoner1999:Ἱναίο 2001a).," A number of scientists have investigated modes of oscillation of an accretion disk as a model of high frequency QPOs \citep[for reviews
see][]{kat98,wag99,kat01a}."
. Ia an important. paper. Okazaki.Kato.&Fukue(1987). showed that axisvmmnmeirie g-mnodes are (rapped in the inner regions of a relativistic disk. where the epievclic frequency. & reaches a maxinum.," In an important paper, \citet{oka87}
 showed that axisymmetric $g$ -modes are trapped in the inner regions of a relativistic disk, where the epicyclic frequency $\kappa$ reaches a maximum."
 This idea was exploited by and a munber of other workers (Perezetal.1997:2001:Wagoner.Silbergleit.&Ortega-RodriguezAbramowiczIxIuzniak2001) who worked out the physical properties of the trapped modes.," This idea was exploited by \citet{now91,now92} and a number of other workers \citep{per97,now97,sil01,wag01,abr01} who worked out the physical properties of the trapped modes."
 Much of the work has focused on neutral (i.e.. non-growing) modes.," Much of the work has focused on neutral (i.e., non-growing) modes."
 The idea in such models is that. although the mocles are stable. they might nevertheless be excited by some driving mechanism such as disk turbulence to produce the observed QPOs.," The idea in such models is that, although the modes are stable, they might nevertheless be excited by some driving mechanism such as disk turbulence to produce the observed QPOs."
 Recently. IXato(2001b.2002) claimed to find that nonaxisvimetric ganodes are trapped between (wo forbidden zones that lie on either side of the corotation radius and that the modes are highly unstable.," Recently, \citet{kat01b,kat02} claimed to find that nonaxisymmetric $g$ -modes are trapped between two forbidden zones that lie on either side of the corotation radius and that the modes are highly unstable."
 Sueh dvnamically unstable modes are very interesting since they would be sell-excited. and would spontaneously grow to non-linear amplitudes without the need for an external driving mechanism., Such dynamically unstable modes are very interesting since they would be self-excited and would spontaneously grow to non-linear amplitudes without the need for an external driving mechanism.
 Motivated by Ixato's work. Li.(2003) studied nonaxisvimmnietric ganodeand panode instabilities in an unmagnetized isothermal accretion disk.," Motivated by Kato's work, \citet{li03} studied nonaxisymmetric $g$ -modeand $p$ -mode instabilities in an unmagnetized isothermal accretion disk."
 Thev found that g- and p-modes with a nonzero number of vertical nodes are stronglv absorbed at corotation and thus not amplified., They found that $g$ - and $p$ -modes with a nonzero number of vertical nodes are strongly absorbed at corotation and thus not amplified.
 Waves without. vertical nodes are amplified (as known earlier. see Papaloizou Pringle 1985: Goldreich. Goodman. Naravan 1986: Naravan. Goldreich. Goodman 1987). but the amplification is verv weak.," Waves without vertical nodes are amplified (as known earlier, see Papaloizou Pringle 1985; Goldreich, Goodman, Narayan 1986; Narayan, Goldreich, Goodman 1987), but the amplification is very weak."
 Therefore. anv energv loss. either during propagation of the wave or during rellection at the boundaries. would kill the instability.," Therefore, any energy loss, either during propagation of the wave or during reflection at the boundaries, would kill the instability."
 Thus. Li et al. (," Thus, Li et al. ("
2003) concluded. in agreement with Kato (2003). (hat nonasxisvinmetric disk modes are not sufficiently unstable to be of interest for the QPO problem.,"2003) concluded, in agreement with Kato (2003), that nonaxisymmetric disk modes are not sufficiently unstable to be of interest for the QPO problem."
 Ortega-Bodriguez Wagoner (2000) found (hat viscosity causes the fundamental g- and panodes in a disk to grow. though the relevance of their result [or QPOs is presently unclear.," Ortega-Rodriguez Wagoner (2000) found that viscosity causes the fundamental $g$ - and $p$ -modes in a disk to grow, though the relevance of their result for QPOs is presently unclear."
 since the amplitude of intensity fluctuations observed in QPOs is often quite large. it is the opinion of the present authors that the oscillations are likely to be the result of a stronglv-growing znstabililaj of some sort in the accretion flow.," Since the amplitude of intensity fluctuations observed in QPOs is often quite large, it is the opinion of the present authors that the oscillations are likely to be the result of a strongly-growing of some sort in the accretion flow."
 Having argued above that a, Having argued above that a
The mass conservation equation becomes The equation for hydrostatic balance is (Thorne 1977) where P is the locally measured pressure. and we neglect contributions to the gravitational mass from the fluid energy density.,"The mass conservation equation becomes The equation for hydrostatic balance is (Thorne 1977) where $P$ is the locally measured pressure, and we neglect contributions to the gravitational mass from the fluid energy density."
 Using the definition of column depth. we rewrite equation (22)) for the thin layer as where Rewriting the entropy and flux equations (Thorne 1977. eqs. [," Using the definition of column depth, we rewrite equation \ref{eq:HB}) ) for the thin layer as where Rewriting the entropy and flux equations (Thorne 1977, eqs. ["
11d] and [11h in a similar way. we find and where F is the heat flux. e is the nuclear energy release rateper unit rest mass. & is the opacity. and 7 is the locally measured temperature.,"11d] and [11h]) in a similar way, we find and where $F$ is the heat flux, $\epsilon$ is the nuclear energy release rateper unit rest mass, $\kappa$ is the opacity, and $T$ is the locally measured temperature."
 Equations (21)). (23)). (25)). and (26)). which describe the structure of the atmosphere. are identical to the Newtonian equations (compare CB 2)," Equations \ref{eq:1}) ), \ref{eq:2}) ), \ref{eq:3}) ), and \ref{eq:4}) ), which describe the structure of the atmosphere, are identical to the Newtonian equations (compare CB 2)."
" These results also hold in the slow rotation approximation. for which the star is spherical to first order in (Q/OQ,y."," These results also hold in the slow rotation approximation, for which the star is spherical to first order in $(\Omega/\Omega_K)^2$."
 The inclusion of rapid rotation has three effects., The inclusion of rapid rotation has three effects.
 The first is to change the volume correction factor Y., The first is to change the volume correction factor ${\cal{V}}$.
 In addition. the mass conservation equation must be covariant so that all observers agree on the rest mass of the star.," In addition, the mass conservation equation must be covariant so that all observers agree on the rest mass of the star."
 This requirement gives an additional Lorentz factor in equation (20)) (e.g.. Friedman et al.," This requirement gives an additional Lorentz factor in equation \ref{eq:mass}) ) (e.g., Friedman et al."
 1986)., 1986).
 However. we find that for the models we consider in this paper. both of these effects are <1%.. and so we neglect them.," However, we find that for the models we consider in this paper, both of these effects are $\lesssim 1$, and so we neglect them."
 Much more important is the correction to gravity due to the rapid rotation., Much more important is the correction to gravity due to the rapid rotation.
 This correction factor is the relativistic generalisation of the Newtonian reduction in gravitational acceleration due to the centrifugal force., This correction factor is the relativistic generalisation of the Newtonian reduction in gravitational acceleration due to the centrifugal force.
" To calculate this. we write the four-velocity as 4=RU+00""). and use the normalization «d,2—1 to solve for the redshift factor where v is given by equation (17))."," To calculate this, we write the four-velocity as $u^\alpha = {\cal{R}}(t^\alpha + \Omega \phi^\alpha)$, and use the normalization $u^\alpha u_\alpha = -1$ to solve for the redshift factor where $v$ is given by equation \ref{eq:v}) )."
 The equation of hydrostatic balance for a rigidly rotating star is then (Friedman et al 1986) where we assume the pressure is much smaller than the rest mass density., The equation of hydrostatic balance for a rigidly rotating star is then (Friedman et al 1986) where we assume the pressure is much smaller than the rest mass density.
 To reduce this equation to the form of equation (23)). we write where the correction factor A accounts for the reduction in g at the equator due to rapid rotation. tending to unity in the limit of zero rotation.," To reduce this equation to the form of equation \ref{eq:2}) ), we write where the correction factor $\cal{A}$ accounts for the reduction in $g$ at the equator due to rapid rotation, tending to unity in the limit of zero rotation."
 This is the most important effect coming from rotation. giving a correction of up to for a rotation frequency of 600Hz and a stiff equation of state.," This is the most important effect coming from rotation, giving a correction of up to for a rotation frequency of $600\ {\rm Hz}$ and a stiff equation of state."
 In summary. the general relativistic equations look the same as the Newtonian equations if we use the correct value of surface gravity e (eq. [29] .," In summary, the general relativistic equations look the same as the Newtonian equations if we use the correct value of surface gravity $g$ (eq. \ref{eq:g2}] ]),"
 and if we identify the column depth y with the rest mass column depth. and the thickness z with the proper thickness of the layer.," and if we identify the column depth $y$ with the rest mass column depth, and the thickness $z$ with the proper thickness of the layer."
 Values of g. V. and A are given in Table | for different models.," Values of $g$, ${\mathcal V}$, and ${\cal A}$ are given in Table \ref{tab:models} for different models."
 To incorporate the general relativistic angular momentum conservation law into our caleulations of spin down. we write the angular momentum per unit rest mass at the equator in the form where f(r) is a function of radius. e.g. f(r)=r in the Newtonian limit.," To incorporate the general relativistic angular momentum conservation law into our calculations of spin down, we write the angular momentum per unit rest mass at the equator in the form where $f(r)$ is a function of radius, e.g. $f(r)=r^2$ in the Newtonian limit."
 If the angular momentum is constant. we find that the change of O with radius is determined by the gradient off with radius:./Z2(-1/2)dInInr)2CL/2)(dInr).," If the angular momentum is constant, we find that the change of $\Omega$ with radius is determined by the gradient of $f$ with radius:$\beta=(-1/2)(d\ln\Omega/d\ln r)=(1/2)(d\ln f/d\ln r)$."
 Now consider a particle which movesO/d from radius RIn tof /dR+Ar. where Ar«R.," Now consider a particle which moves from radius $R$ to $R+\Delta r$, where $\Delta
r\ll R$."
" Conservation of angular momentum gives u,(R+Ar)=u,,(R)."," Conservation of angular momentum gives $u_\phi(R+\Delta
r)=u_\phi(R)$."
 Expanding in a Taylor series around r=R. we find In the Newtonian limit. for which .)=|. we recover the familiar result AOο=—2Ar/R.," Expanding in a Taylor series around $r=R$ , we find In the Newtonian limit, for which $\beta=1$, we recover the familiar result $\Delta\Omega/\Omega=-2\Delta r/R$."
 To calculate the moment of inertia. we must modify equation (4)).," To calculate the moment of inertia, we must modify equation \ref{eq:right}) )."
 Again writing the angular momentum in the form of equation (31)). and expanding f(r) around r=R. we find which reduces to equation (4)) in the Newtonian limit.," Again writing the angular momentum in the form of equation \ref{eq:f}) ), and expanding $f(r)$ around $r=R$, we find which reduces to equation \ref{eq:right}) ) in the Newtonian limit."
 The calculations presented in $2 were for a surface gravity ο-1.9«10!ems? (the Newtonian gravity for a 1.4M... R=10km neutron star).," The calculations presented in 2 were for a surface gravity $g=1.9\times 10^{14}\ {\rm cm\ s^{-2}}$ (the Newtonian gravity for a $1.4 M_\odot$, $R=10\ {\rm km}$ neutron star)."
 We now rescale these fiducial results to different masses and equations of state. including the effects of rapid rotation and general relativity.," We now rescale these fiducial results to different masses and equations of state, including the effects of rapid rotation and general relativity."
 To calculate the thickness of the atmosphere in $2. we integrated hydrostatic balance in the form dP/dz=—pGM/R-. giving the thickness Az/R«IH/gRR/M (see also 82 of CB).," To calculate the thickness of the atmosphere in 2, we integrated hydrostatic balance in the form $dP/dz=-\rho\,GM/R^2$ , giving the thickness $\Delta z/R\propto 1/gR\propto R/M$ (see also 2 of CB)."
 In addition. we must include a factor Y to obtain the correct general relativistic value for e. and afactor A to include the effects of rotation on gravity.," In addition, we must include a factor ${\cal V}$ to obtain the correct general relativistic value for $g$ , and afactor ${\cal A}$ to include the effects of rotation on gravity."
 We thus find the fractional change in thickness of the atmosphere is, We thus find the fractional change in thickness of the atmosphere is
Iluetuations.,fluctuations.
 Cen presents 32 different cosmological modcls to determine the fitting parameters in his Gaussian. peas method. using them to obtain the cluster mass function.," Cen presents 32 different cosmological models to determine the fitting parameters in his Gaussian peak method, using them to obtain the cluster mass function."
 In particular. we take four of these models that. present dillerent characteristics related to the present radii and present density contrasts of the voids.," In particular, we take four of these models that present different characteristics related to the present radii and present density contrasts of the voids."
 Ehe models chosen have their parameters described in Table 1 (see also Cen 19982.h).," The models chosen have their parameters described in Table 1 (see also Cen 1998a,b)."
 All models considered here have ‘excess power EP=1.30+ 0.15. (, All models considered here have `excess power' $EP=1.30\pm 0.15$ . (
Excess power. is defined as £2?=3.405fox. where σος is the linear rms density [uctuation in a 25h!XMpe top-hat. sphere at 2=0. and σ has a similar definition in an Sf!Mpe top-bat sphere.),"`Excess power' is defined as $EP\equiv 3.4 \sigma_{25}/\sigma_8$, where $\sigma_{25}$ is the linear rms density fluctuation in a $25h^{-1}{\rm Mpc}$ top-hat sphere at $z=0$, and $\sigma_8$ has a similar definition in an $8h^{-1}{\rm Mpc}$ top-hat sphere.)"
" ""This parameter describes the shape of the power spectrum on scales ~S03005+Alpe and its values range from 1.15 to 145 to it the data.", This parameter describes the shape of the power spectrum on scales $\sim 8-300h^{-1}{\rm Mpc}$ and its values range from $1.15$ to $1.45$ to fit the data.
 In the first column of Table 1: we present the number of 1e model we use here the same numbers presented by Cen (1998a.b)]: in the second. and third columns we present the lensity. parameter for the matter and that associated with 16 cosmological constant. respectively: the fourth column esents the rms density DDuctuation in an Sf1Alpe top-hat sphere at >=0: in the fifth column we present the power --ndex n: in the sixth column we present the value of the Llubble constant: anc finally the seventh column indicates Ye type of power spectrum usec.," In the first column of Table 1 we present the number of the model [we use here the same numbers presented by Cen (1998a,b)]; in the second and third columns we present the density parameter for the matter and that associated with the cosmological constant, respectively; the fourth column presents the rms density fluctuation in an $8h^{-1}{\rm Mpc}$ top-hat sphere at $z=0$; in the fifth column we present the power index $n$; in the sixth column we present the value of the Hubble constant; and finally the seventh column indicates the type of power spectrum used."
 For the normalization of the power spectra studied here we followed the usual methodology. found in the literature. where the value of as permits one to find the constant zd of he power spectra.," For the normalization of the power spectra studied here we followed the usual methodology found in the literature, where the value of $\sigma_{8}$ permits one to find the constant $A$ of the power spectra."
 Our normalized spectra follow the sanic shape seen in fig., Our normalized spectra follow the same shape seen in fig.
 2 of Cen (199Sa)., 2 of Cen (1998a).
 ‘Thus. for a ὃνgiven perturbation with mass Ave. and for he mocels described in Table 1. we can obtain the present racius and the present density contrast.," Thus, for a given perturbation with mass $M_{\rm C}$, and for the models described in Table 1, we can obtain the present radius and the present density contrast."
 Phen. to determine he initial conditions at the beginning of the recombination era. we used the general formula. given by Opher. Pires de Araujo (1997) for the amplification of the density »erturbations (their studs incorporates both growing ancl decaving mocles).," Then, to determine the initial conditions at the beginning of the recombination era, we used the general formula given by Opher, Pires de Araujo (1997) for the amplification of the density perturbations (their study incorporates both growing and decaying modes)."
 We start the calculations. at the. beginning of the recombination era. at redshift spo~1000. where the ionization degree begins to be significantly lower than unity.," We start the calculations at the beginning of the recombination era, at redshift $z_{\rm rec} \sim 1000$, where the ionization degree begins to be significantly lower than unity."
 For zz1000 the matter is completely ionized and density perturbations (positive or negative) cannot evolve due to the photon drag that inhibits the growth of any peculiar velocity.," For $z_{\rm rec}
> 1000$ the matter is completely ionized and density perturbations (positive or negative) cannot evolve due to the photon drag that inhibits the growth of any peculiar velocity."
 We start our calculations assuming that the underdense regions evolve initially with a velocity field. given by the Llubble Dow., We start our calculations assuming that the underdense regions evolve initially with a velocity field given by the Hubble flow.
" In the present study we assume that initially where pi is the total density. (dark matter and barvons). (n. is the barvon density. pq is the dark matter density. and 9p. pn, and dp. are the respective density. perturbations."," In the present study we assume that initially where $\bar{\rho}_{\rm t}$ is the total density (dark matter and baryons), $\bar{\rho}_{\rm b}$ is the baryon density, $\bar{\rho}_{\rm d}$ is the dark matter density, and $\delta
\rho_{\rm t}$, $\delta \rho_{\rm b}$ and $\delta \rho_{\rm d}$ are the respective density perturbations."
 The barvonic mass contained within the underdense region is obtained from and. analogously. the collisionless clark matter within the unclerdense region is The initial radius of the underdense region is given hy With such an initial profile. the density at the void boundary does not fall smoothly to the density of the Universe.," The baryonic mass contained within the underdense region is obtained from and, analogously, the collisionless dark matter within the underdense region is The initial radius of the underdense region is given by With such an initial profile, the density at the void boundary does not fall smoothly to the density of the Universe."
 Thus. the interface (between the external medium and. the last shell of the computational grid) is a membrane-like one.," Thus, the interface (between the external medium and the last shell of the computational grid) is a membrane-like one."
 This membrane is maintained throughout the void time evolution. because the mass inside the void region is taken to be constant throughout the void evolution.," This membrane is maintained throughout the void time evolution, because the mass inside the void region is taken to be constant throughout the void evolution."
 In Tables 2-5 we present the main results for the models described. in Table 1., In Tables 2-5 we present the main results for the models described in Table 1.
 We see that the spectra. of perturbations used here produce cdillerent results concerning the final diameter of the voids (and [or their density contrasts) for the same mass perturbation Mc., We see that the spectra of perturbations used here produce different results concerning the final diameter of the voids (and for their density contrasts) for the same mass perturbation $M_{\rm C}$.
 In particular. if we consider that voids are regions with density contrasts 0—0.6. then large voids with ciameters of D>20.30h!Mpe the sizes found. by El-Ad et al. (," In particular, if we consider that voids are regions with density contrasts $\bar \delta \leq -0.6$, then large voids with diameters of $\bar D > 20-30 h^{-1} {\rm Mpc}$ [the sizes found by El-Ad et al. ("
1996) and Mülller et al. (,1996) and Mülller et al. (
2000)] could be produced only for the spectrum 26 or similar one (see Tables 1 and 5).,2000)] could be produced only for the spectrum 26 or similar one (see Tables 1 and 5).
 1n fact. it is worth mentioning that even if (2003)& negative density perturbations are considered for the models 14. 2 and 23 or similar ones. large. void. regions cannot be accounted. for.," In fact, it is worth mentioning that even if $(2-3)\sigma$ negative density perturbations are considered for the models 14, 20 and 23 or similar ones, large void regions cannot be accounted for."
 For example. for model 23. (see Tables ] and 4) a 2e negative density perturbation of 1034M. would produce a void with barvonic (dark matter) diameter ol 14.7*\Ipe (16.7h ‘\Ipe) with dpe=0.639 (Opp= 0.639).," For example, for model 23 (see Tables 1 and 4) a $2\sigma$ negative density perturbation of $10^{14}{\rm M}_{\odot}$ would produce a void with baryonic (dark matter) diameter of $14.7 h^{-1}{\rm
Mpc}$ $16.7 h^{-1}{\rm Mpc}$ ) with $\delta_{\rm BF}=-0.639$ $\delta_{\rm DF}= -0.639$ )."
 If instead a 30 negative perturbation is considered. avoid with barvonie (dark matter) diameter of 17.2hINpe (18.95. *Mpe) with opr=O.776 (ου= 0.776) would be produced.," If instead a $3\sigma$ negative perturbation is considered, a void with baryonic (dark matter) diameter of $17.2 h^{-1}{\rm Mpc}$ $18.9 h^{-1}{\rm Mpc}$ ) with $\delta_{\rm BF}=-0.776$ $\delta_{\rm DF}= -0.776$ ) would be produced."
 Although large. these voids are not enough to explain those reported in the literature.," Although large, these voids are not enough to explain those reported in the literature."
 Lt is worth noting that some authors define void regions as having a density of 0.2 times the mean galaxy. density (Schmidt. Iden Melott 2001).," It is worth noting that some authors define void regions as having a density of 0.2 times the mean galaxy density (Schmidt, Ryden Melott 2001)."
 Phe density contrast in clusters of galaxies is around 1.5-3.0. which means that the density contrast in void regions would amount to 0=(0.52: 0.2). which is less conservative than our definition.," The density contrast in clusters of galaxies is around 1.5-3.0, which means that the density contrast in void regions would amount to $ \delta = -(0.5 \pm 0.2)$ , which is less conservative than our definition."
 Even considering that regions with 0~0.3 are volds. it is casily seen from our results that the largest. voids in the Universe cannot be accounted for by models 14. 20 and 23 or similar ones.," Even considering that regions with $\delta \sim -0.3$ are voids, it is easily seen from our results that the largest voids in the Universe cannot be accounted for by models 14, 20 and 23 or similar ones."
 A possible explanation for the largest. voids could. be related to the interaction and eventual merger of two or more small voids., A possible explanation for the largest voids could be related to the interaction and eventual merger of two or more small voids.
 We refer the reader to à paper by Dubinski οἱ al. (, We refer the reader to a paper by Dubinski et al. (
1993) who investigated. the evolution. interaction and merger of voids formed from negative density perturbations.,"1993), who investigated the evolution, interaction and merger of voids formed from negative density perturbations."
 Then. the merger of small voids could be à way to build up laree voids [rom the spectra of density perturbations related to mocels 14. 20 and 23.," Then, the merger of small voids could be a way to build up large voids from the spectra of density perturbations related to models 14, 20 and 23."
"a correlation will not qualitatively alter our results, but will need to be taken into account once precision data becomes available.","a correlation will not qualitatively alter our results, but will need to be taken into account once precision data becomes available."
" The distribution function for dp from existing data has a positive mean value (dp= 0.049), implying that the post-calibration luminosity of the delayed population is dimmer than the prompt population by ~5%.."," The distribution function for $\delta_D$ from existing data has a positive mean value $\delta_D=0.049$ ), implying that the post-calibration luminosity of the delayed population is dimmer than the prompt population by $\sim$."
" The standard deviation of this dp is c=0.097, which suggests that current SN data is consistent with the absence of a two-population bias."," The standard deviation of this $\delta_D$ is $\sigma=0.097$, which suggests that current SN data is consistent with the absence of a two-population bias."
 The Howelletal.(2007) result of ~ (12+4)% difference in the intrinsic luminosity is a pre-calibration difference., The \citet{howell:07} result of $\sim$ $(12 \pm 4)$ difference in the intrinsic luminosity is a -calibration difference.
 We find a 7»(53:9)96 difference in the post-calibration luminosity between the two types of SNe., We find a $\sim(5 \pm 9)$ difference in the -calibration luminosity between the two types of SNe.
" While this uncertainty is larger than the value, it is estimated from the Hubble diagram, pre-calibrationand is independent of the directlyempirical stretch-luminosity relation."," While this uncertainty is larger than the pre-calibration value, it is estimated directly from the Hubble diagram, and is independent of the empirical stretch-luminosity relation."
 This important consistency check will improve as the SN sample sizes increase., This important consistency check will improve as the SN sample sizes increase.
" To study the impact of a potential luminosity difference on the measurement of the dark energy EOS, we also fit the same data (SN+BAO+CMB) assuming a wCDM model, and take the corresponding priors Ho=72.13:7.5 and Q,,/?20.132930.0066 (Komatsuetal.2008;Freedman 2001).."," To study the impact of a potential luminosity difference on the measurement of the dark energy EOS, we also fit the same data (SN+BAO+CMB) assuming a $w$ CDM model, and take the corresponding priors $H_0=72.1\pm7.5$ and $\Omega_m
h^2=0.1329\pm0.0066$ \citep{komatsu:08,freedman:01}. ."
" We first consider two extreme cases: ignoring the two-population systematic (δρ= 0), and allowing for a completely unconstrained (óp with no "," We first consider two extreme cases: ignoring the two-population systematic $\delta_D=0$ ), and allowing for a completely unconstrained systematic $\delta_D$ with no priors)."
"With óp free, our analysis yields a systematicbest-fit time-independent priors).EOS parameter w=—0.986+0.180, with a best-fit dp=0.040+0.279."," With $\delta_D$ free, our analysis yields a best-fit time-independent EOS parameter $w=-0.986 \pm 0.180$, with a best-fit $\delta_D = 0.040 \pm 0.279$."
" If we set dp=0, the same data provides a constraint of w=—0.960+ 0.066."," If we set $\delta_D=0$, the same data provides a constraint of $w=-0.960 \pm
0.066$ ."
" By incorporating a two-population effect in the fit, the errors on the best-fit EOS degrade by factor of ~3."," By incorporating a two-population effect in the fit, the errors on the best-fit EOS degrade by a factor of $\sim3$."
" In addition, we note a shift in the best-fit avalue of w between the two cases."," In addition, we note a shift in the best-fit value of $w$ between the two cases."
" If there is a residual difference in the calibrated luminosity, ignoring dp might bias the dark energy estimate."," If there is a residual difference in the calibrated luminosity, ignoring $\delta_D$ might bias the dark energy estimate."
" As we show in the left panel of Figure 1, the increase in the error of the best-fit EOS is due to the degeneracy between dark energy (w) and the two-population systematic (dp)."," As we show in the left panel of Figure 1, the increase in the error of the best-fit EOS is due to the degeneracy between dark energy $w$ ) and the two-population systematic $\delta_D$ )."
" Following Howelletal.(2007),, it may be possible to calibrate out the luminosity difference between the two components with large samples of SNe (e.g., by looking for characteristic properties in the SN spectra, or by correlating SN luminosities with host galaxy types)."," Following \citet{howell:07}, it may be possible to calibrate out the luminosity difference between the two components with large samples of SNe (e.g., by looking for characteristic properties in the SN spectra, or by correlating SN luminosities with host galaxy types)."
" This will lead to a prior constraint on dp, although uncertainties will still remain on the redshift evolution (as estimated by fp(z))."," This will lead to a prior constraint on $\delta_D$, although uncertainties will still remain on the redshift evolution (as estimated by $f_D(z)$ )."
" Assuming the two-population fraction, and its redshift evolution, is perfectly known, we analyze the same SN data a Gaussian prior on 6p."," Assuming the two-population fraction, and its redshift evolution, is perfectly known, we analyze the same SN data assuming a Gaussian prior on $\delta_D$."
" Taking this prior to have zeroassuming mean and o=(0.25,0.1,0.05), the errors on w are (0.130,0.086, 0.072), with the best-fitvalue of w being approximately the same as was found for the dp=0 case."," Taking this prior to have zero mean and $\sigma=
(0.25, 0.1, 0.05)$, the errors on $w$ are $0.130, 0.086, 0.072$ ), with the best-fitvalue of $w$ being approximately the same as was found for the $\delta_D=0$ case."
" Thus, by including a two-population systematic in the fit, with a prior on dp centered at 0 with a 0.05 mag dispersion, we recover the same best-fit w, but with ~10% degradation in the error bars."," Thus, by including a two-population systematic in the fit, with a prior on $\delta_D$ centered at 0 with a 0.05 mag dispersion, we recover the same best-fit $w$ , but with $\sim10\%$ degradation in the error bars."
" We now turn to proposed SN surveys, and investigate how uncertainties in a possible two-population systematic impact measurements of w. We mock SN catalogs with 300 SNe uniformly distributed at generatez<0.1, and 2000 SNe in the range 0.1<z€1.7, similar to a JDEM--like survey (Kimetal. 2004)."," We now turn to proposed SN surveys, and investigate how uncertainties in a possible two-population systematic impact measurements of $w$ We generate mock SN catalogs with 300 SNe uniformly distributed at $z<0.1$, and 2000 SNe in the range $0.1 \leq
z \leq 1.7$, similar to a -like survey \citep{kim:04}."
". We incorporate an intrinsic Gaussian scatter of 0.1 mag for each SN, and take the relative fractionof delayed and prompt SNe, fp(z), given by Eq."," We incorporate an intrinsic Gaussian scatter of 0.1 mag for each SN, and take the relative fractionof delayed and prompt SNe, $f_D(z)$, given by Eq."
 1 with SB values of A24.4x10? and B=2.6., \ref{eqn:twopop} with SB values of $A=4.4\times 10^{-2}$ and $B=2.6$.
" We assume different values (0.025, 0.05, 0.1 mag) for the underlying two-population bias (dp)."," We assume different values (0.025, 0.05, 0.1 mag) for the underlying two-population bias $\delta_D$ )."
" We fit each mock data set of 2300 SNe, along with the 2 existing BAO measurements, to a wWCDMmodel, with the corresponding wCDM priors as before."," We fit each mock data set of 2300 SNe, along with the 2 existing BAO measurements, to a $w$ CDMmodel, with the corresponding $w$ CDM priors as before."
" When the data we consider two extreme cases: dp= 0, and óp fittingcompletelyunconstrained."," When fitting the data we consider two extreme cases: $\delta_D=0$ , and $\delta_D$ completelyunconstrained."
"Our results for dp= 0.025, from 200 separate mocks, are summarized in the right panel of Figure 1.","Our results for $\delta_D=0.025$ , from 200 separate mocks, are summarized in the right panel of Figure 1."
" The hatched histogram, which peaks at -0.974,"," The hatched histogram, which peaks at -0.974,"
is nonetheless informative.,is nonetheless informative.
 We shall compute this quantity for a quantum Arnold cat that encounters multiparticle scattering., We shall compute this quantity for a quantum Arnol'd cat that encounters multi–particle scattering.
 Indeed. that scattering by simall particles can provide a fully dynamical model of decoherence las been already shown by Joos aud Ζο [1]: we shall only reuder this model explicit aud computable in a form that well ties with the classical 11odel for the Arnold cat map. that can be seen as the Floquet evolution of a kicked particle ou a torus.," Indeed, that scattering by small particles can provide a fully dynamical model of decoherence has been already shown by Joos and Zeh \cite{joos}: we shall only render this model explicit and computable in a form that well ties with the classical model for the Arnol'd cat map, that can be seen as the Floquet evolution of a kicked particle on a torus."
" As noted above. many works have alrewdy appeared liakine decolerence and chaoticity. so that any list of references. dnchiding ours |21.20.22,29| is forced to be utterly partial."," As noted above, many works have already appeared linking decoherence and chaoticity, so that any list of references, including ours \cite{ian1,monpaz,alik0,marcos,zurek0,todd} is forced to be utterly partial."
 Nonetheless. we believe that the two theoretical poiuts raised above have never been studied together so fu.," Nonetheless, we believe that the two theoretical points raised above have never been studied together so far."
 We therefore present iu this paper the first results of our analysis. via iuforiiatiou dynamical eutropies. of a fully dynamical model of a decohering system.," We therefore present in this paper the first results of our analysis, via information--theoretical dynamical entropies, of a fully dynamical model of a decohering system."
 Further study. both from the theoretical aud the munerical point of view. will follow.," Further study, both from the theoretical and the numerical point of view, will follow."
 Iu the next section we review the classical Arnold cat aud its quantization., In the next section we review the classical Arnol'd cat and its quantization.
 Then. in Sect.," Then, in Sect."
 3 we introduce the multiparticle Arnold cat. in its classical aud «quautuni versions.," \ref{sec-mpac} we introduce the multi–particle Arnol'd cat, in its classical and quantum versions."
 The problems involved in deriving a scattering matrix for quanti particles on a torus are solved iu Sect. [., The problems involved in deriving a scattering matrix for quantum particles on a torus are solved in Sect. \ref{sec-scat}.
 Next. in Sect.," Next, in Sect."
 5. we review the esseutial ingredieuts of A-F entropy. particularly the finite time Shannon Alicki Fauues eutropy. that is then computed ummerically in Sect.," \ref{sec-af} we review the essential ingredients of A-F entropy, particularly the finite time Shannon Alicki Fannes entropy, that is then computed numerically in Sect."
 6 for the nmltiparticle Arnold cat., \ref{sec-resu} for the multi–particle Arnol'd cat.
 Its scaling with the system parameters is examined and it is discussed in the Conclusions., Its scaling with the system parameters is examined and it is discussed in the Conclusions.
 Consider a point particle of mass AL subject to move on a ring., Consider a point particle of mass $M$ subject to move on a ring.
 Mathematically. this ring is a torus. of length £. labelled by the variable Qc[0.£L).," Mathematically, this ring is a torus, of length $L$, labelled by the variable $Q \in [0,L)$."
 Tu the absence of auv interaction. this particle rotates with coustant velocity around the rine.," In the absence of any interaction, this particle rotates with constant velocity around the ring."
 To the contrary. it is also subject to the action of a periodic impulsive force. of period T. thathas the effect of changing instautancously its momentum: the relative IEuniltouian is therefore P being the conjugate momentum to Q aud κ a coupling coustaut.," To the contrary, it is also subject to the action of a periodic impulsive force, of period $T$, thathas the effect of changing instantaneously its momentum: the relative Hamiltonian is therefore $P$ being the conjugate momentum to $Q$ and $\kappa$ a coupling constant."
 This is the Hamiltonian for the Arnold cat mapping., This is the Hamiltonian for the Arnol'd cat mapping.
 In fact. it cutails a classical period evolution operator that acts as follows on the dvuamical variables observed at the time inunediately following the action of the mupulsive force: It is now convenicut to rescale momentum 2 by multiplviue it by T/AJ.," In fact, it entails a classical period evolution operator that acts as follows on the dynamical variables observed at the time immediately following the action of the impulsive force: It is now convenient to rescale momentum $P$ by multiplying it by $T/M$ ."
 The new momentum. 2=Lp has the dimensions of a leugth.," The new momentum, $\tilde{P}:= \frac{T}{M} P$ has the dimensions of a length."
 In the new variables.," In the new variables,"
There have been several analyses of the clustering of the SDSS LRG spectroscopic sample.,There have been several analyses of the clustering of the SDSS LRG spectroscopic sample.
 Percivaletal.(2010) use power spectrum analysis to extract the BAO signals., \cite{Percival:2009xn} use power spectrum analysis to extract the BAO signals.
 The sample they use include not only LRGs but also main galaxies sample., The sample they use include not only LRGs but also main galaxies sample.
 They obtain the constraint on. rs/L)4-(0.35)0.1007+ 0.0036., They obtain the constraint on $r_s/D_V(0.35) = 0.1097\pm0.0036$ .
 Reidetal.(2010) apply power spectrum analysis on the reconstructed halo density field derived from SDSS DR7full (flux-limited LRG) sample., \cite{Reid:2009xm} apply power spectrum analysis on the reconstructed halo density field derived from SDSS DR7full (flux-limited LRG) sample.
 Their measurement of refDy(0.385)Μ.Μ.”0.2) which is similar to Percivaletal.0103..," Their measurement of $r_s/D_V(0.35) = 0.1097^{+0.0039}_{-0.0042}\, (k_{max}=0.2)$ which is similar to \cite{Percival:2009xn}."
. There is about Lo difference between their results and ours., There is about $1\sigma$ difference between their results and ours.
 However. in Reidetal.(2010). they show a dependence on the minimum scales of the range. he. r./Do(0.35)=ΟΠΠfron(uus0.15). and refPy(0.35)=0.1136onze(ku0.1).," However, in \cite{Reid:2009xm}, they show a dependence on the minimum scales of the range, i.e. $r_s/D_V(0.35) = 0.1118^{+0.0043}_{-0.0046}\, 
(k_{max}=0.15)$ and $r_s/D_V(0.35) = 0.1136^{+0.0070}_{-0.0072}\, (k_{max}=0.1)$."
 Their results are more consistent with ours when a smaller Ανω Ge. a larger minimum scale) is used. which represents a more conservative choice for the scale range.," Their results are more consistent with ours when a smaller $k_{max}$ (i.e., a larger minimum scale) is used, which represents a more conservative choice for the scale range."
 reftable:test shows the systematic tests that we have done varying key assumptions made in our analysis., \\ref{table:test} shows the systematic tests that we have done varying key assumptions made in our analysis.
 These include the range of scales used to calculate the correlation function. the nonlinear damping scale. an overall shift in the measured correlation function due to a systematic error.," These include the range of scales used to calculate the correlation function, the nonlinear damping scale, an overall shift in the measured correlation function due to a systematic error."
 We vary the effective redshift (from ορ=0.33 to zy 0.35) used to calculate the theoretical model., We vary the effective redshift (from $z_{eff}=0.33$ to $z_{eff}=0.35$ ) used to calculate the theoretical model.
 We rescale the results to 2=0.35 for comparison and find that the results are insensitive to the effective redshift., We rescale the results to $z=0.35$ for comparison and find that the results are insensitive to the effective redshift.
" We also test the sensitivity of our results to the nonlinear damping scale. &,."," We also test the sensitivity of our results to the nonlinear damping scale, $k_{\star}$."
 Although ἂν can be predicted accurately in the real space (Crocce&Scoccimarro2006:Matsubara2007).. in the redshift space. it would also depend on the redshift distortions which cannot be well determined from the spherically-averaged correlation function.," Although $k_{\star}$ can be predicted accurately in the real space \citep{Crocce:2005xz,Matsubara:2007wj}, in the redshift space, it would also depend on the redshift distortions which cannot be well determined from the spherically-averaged correlation function."
 In reftable:test.. one can tell that the results are not sensitive to Ay.," In \\ref{table:test}, one can tell that the results are not sensitive to $k_{\star}$."
 In principle. the range of scales chosen for the analysis should be as large as »ossibe. in order to derive the tightest constraints.," In principle, the range of scales chosen for the analysis should be as large as possible, in order to derive the tightest constraints."
 However. we do not use the small scales (s.<404+Mpci. where the scale dependence of redshift distortion and galaxy bias are not negligiple and cannot be accurately determined at present.," However, we do not use the small scales $s<40\ h^{-1}$ Mpc), where the scale dependence of redshift distortion and galaxy bias are not negligible and cannot be accurately determined at present."
 According to Fig., According to Fig.
 5 in Eisensteinetal.(2005)... these effects are negligible at 540h!Mpe.," 5 in \cite{Eisenstein:2005su}, these effects are negligible at $s>40\ h^{-1}$ Mpc."
" Wevary the minimum scales used and find that the εςσον(0.35) is insensitive to it but OQ,h? is not.", Wevary the minimum scales used and find that the $r_s(z_d)/D_V(0.35)$ is insensitive to it but $\Omega_m h^2$ is not.
" This indicates the robustness of the measurement of relay)fD(0:35) ebut not Q,,5?) from this paper.", This indicates the robustness of the measurement of $r_s(z_d)/D_V(0.35)$ (but not $\Omega_m h^2$ ) from this paper.
 On larger scales (s>1305. Mp. the observed correlation is significantly higher than expected in conventional models of galaxy clustering.," On larger scales $s>130\ h^{-1}$ Mpc), the observed correlation is significantly higher than expected in conventional models of galaxy clustering."
 This high tail problem was reported in previous work. see. e.g. Eisensteinetal.(2005)... Hutsi(2005). and Sanchezetal.(2009)..," This high tail problem was reported in previous work, see, e.g., \cite{Eisenstein:2005su}, \cite{Hutsi:2005qv}, and \cite{Sanchez:2009jq}."
 They found that the observed correlation function could be fitted better by lowering all the data points by a constant., They found that the observed correlation function could be fitted better by lowering all the data points by a constant.
 In other words. they assumed a constant shift from some systematic error.," In other words, they assumed a constant shift from some systematic error."
 Although this systematic error is unknown. we could minimize its effect by using smaller scale.," Although this systematic error is unknown, we could minimize its effect by using smaller scale."
 The reason is that the correlation function has larger value at smaller scale so that the results are less sensitive to the shift., The reason is that the correlation function has larger value at smaller scale so that the results are less sensitive to the shift.
 We choose s=1205.*Alpe as our boundary for the large scale and show that the results are insensitive to the constant shift by lowering down the data points of the observed correlation function by 0.002., We choose $s=120 h^{-1}Mpc$ as our boundary for the large scale and show that the results are insensitive to the constant shift by lowering down the data points of the observed correlation function by 0.002.
" We find that Q,,.A° varies by Lo and ry/4:(0.35) only varies by 0.26.", We find that $\Omega_m h^2$ varies by $\sigma$ and $r_s/D_V(0.35)$ only varies by $\sigma$.
 Therefore. our measurementry/4:(0.35) is robust to the systematic shift.," Therefore, our measurement of $r_s/D_V(0.35)$ is robust to the systematic shift."
" This is another indicationof that our measurementof Ενσον(0.35) (but not that of ©,,,47) is robust.", This is another indication that our measurementof $r_s(z_d)/D_V(0.35)$ (but not that of $\Omega_m h^2$ ) is robust.
 We have presented our first results for the model independent constraints on dark energy from. the spherically-averaged correlation function of SDSS DR7 data. using an MCMC likelihood analysis.," We have presented our first results for the model independent constraints on dark energy from the spherically-averaged correlation function of SDSS DR7 data, using an MCMC likelihood analysis."
 Our. constraints on [DU(0.35). Oh. Γι)Do(0.35). (0.35)1 are summarized by reftable:mean and 2..," Our constraints on $\{D_V(0.35)$, $\Omega_mh^2$, $r_s(z_d)/D_V(0.35)$, $A(0.35)\}$ are summarized by \\ref{table:mean} and \ref{table:covar_matrix}."
" Applying these results to constraining a constant dark energy equation of state without assuming a flat Universe (the owCDM model). and combining with WMAP7 and Union2 SN data sets. we find that Q,=0.0082(00d and wn= (10. consistent with a flat universe with a cosmological constant (£9,— Qi= 1)."," Applying these results to constraining a constant dark energy equation of state without assuming a flat Universe (the owCDM model), and combining with WMAP7 and Union2 SN data sets, we find that $\Omega_k=-0.0032^{+0.0074}_{-0.0072}$ and $w=-1.010^{+0.046}_{-0.045}$ , consistent with a flat universe with a cosmological constant $\Omega_k=0$ , $w=-1$ )."
 We have also measured the model independent constraints of 1(0.55) and D41(0.35) from the spherically-averaged correlation function from SDSS DR7 LRGs. as a baseline for comparison with," We have also measured the model independent constraints of $H(0.35)$ and $D_A(0.35)$ from the spherically-averaged correlation function from SDSS DR7 LRGs, as a baseline for comparison with"
transitions Irom adjusted ab initio potential energv surfaces (IIuangetal.2008.201Lab:Yurchenkoetal.2009. 2011).,"transitions from adjusted ab initio potential energy surfaces \citep{huang2008, huang11a, huang11b, yurchenko09, yurchenko11}."
. Caleulated line lists suitable for spectra of ammonia at 300 Ix (Yurchenkoetal.2009). and 1500 IX. (Yurchenkoetal.2011). are now available., Calculated line lists suitable for spectra of ammonia at 300 K \citep{yurchenko09} and 1500 K \citep{yurchenko11} are now available.
 Cool objects ave less Iuminous than hot stars and emit most of (heir radiation in the inlrared., Cool objects are less luminous than hot stars and emit most of their radiation in the infrared.
 High resolution infrared spectra of these objects are best recorded. with large 10 m class telescopes., High resolution infrared spectra of these objects are best recorded with large 10 m class telescopes.
 In 1995 the first brown εναν]. Gliese 229D was identified bv the presence of methane in ils spectrum (Oppenheimerοἱal.1995).," In 1995 the first brown dwarf, Gliese 229B was identified by the presence of methane in its spectrum \citep{oppen95}."
. Objects wilh masses less than 0.0075 solar masses are unable to fuse hydrogen in their cores and are called brown dwarls., Objects with masses less than 0.0075 solar masses are unable to fuse hydrogen in their cores \citep{burrows01} and are called brown dwarfs.
 Brown dwarf atmospheres are cool enough (iwpically 700 2.000 IX) to sustain a rich molecular environment (Ixirkpatrick2005).," Brown dwarf atmospheres are cool enough (typically 700 – 2,000 K) to sustain a rich molecular environment \citep{kirkpatrick05}."
. Since the discovery ol Gliese 229D. the L and T classes have been introduced to describe objects cooler than A-dwarls (Leggettetal.2001).," Since the discovery of Gliese 229B, the L and T classes have been introduced to describe objects cooler than M-dwarfs \citep{leggett01}."
.. L-dwarls display near infrared. electronic transitions of metal hvdrides such as Fell (Hargreavesetal.2010) ancl Crll (INirkpatricketal.1999) and have weak TiO and VO transitions which are characteristic of M-cdwarls 2007).., L-dwarfs display near infrared electronic transitions of metal hydrides such as FeH \citep{hargreaves10} and CrH \citep{kirkpatrick99} and have weak TiO and VO transitions which are characteristic of M-dwarfs \citep{boch2007}. .
" In contrast. T-dwarls like Gliese 229D have prominent overtone Uransitions of hot 1,0 and CIL; (Durgasserοἱal.2006)."," In contrast, T-dwarfs like Gliese 229B have prominent overtone transitions of hot $_{2}$ O and $_{4}$ \citep{burg2006}."
. Dased on observations taken [rom theTelescope. Cushingetal.(2006) identified prominent HO. CIL; and NIL; absorption bands between ο 14 jan. including the NIL; ο band near Ll jan. IH is predicted that a new cooler Classification. a Y-cdwarl (< 100 IX) is required below the T-dwarf class (INirkpatrick2005) with (he prediction that strong NIL; absorptions will help to differentiate (he Y-dwarls from T-dwarls (Leggettοἱal.2007:Delorme2003).," Based on observations taken from the, \citet{cushing06} identified prominent $_{2}$ O, $_{4}$ and $_{3}$ absorption bands between 6 – 14 $\mu$ m, including the $_{3}$ $\nu_{2}$ band near 11 $\mu$ m. It is predicted that a new cooler classification, a Y-dwarf $<$ 700 K) is required below the T-dwarf class \citep{kirkpatrick05} with the prediction that strong $_{3}$ absorptions will help to differentiate the Y-dwarfs from T-dwarfs \citep{leg07, delorme08}."
. Numerous unassigned features are present in (he observed SEDs of brown clwarls., Numerous unassigned features are present in the observed SEDs of brown dwarfs.
 Due to a lack of experimental data recorded in thetemperature ranges appropriate for these objects (or reliable theoretical predictions). the best option is (ο extrapolate room temperature ‘cold’ line parameters such as the IIITRAN 2008 database," Due to a lack of experimental data recorded in thetemperature ranges appropriate for these objects (or reliable theoretical predictions), the best option is to extrapolate room temperature `cold' line parameters such as the HITRAN 2008 database"
(ο equation (19)) ean be given in terms of the parameter [ and the Alfvénn speed.,to equation \ref{dv2}) ) can be given in terms of the parameter $f$ and the Alfvénn speed.
 The results are shown in Figures 11. and 12.. where the rms velocity is given in kilometers per second.," The results are shown in Figures \ref{fig11} and \ref{fig12}, where the rms velocity is given in kilometers per second."
 The left column of Figure LL shows (007)? for the nonlinear case. for three different optical depths.," The left column of Figure \ref{fig11} shows $\langle \delta v^2
\rangle^{1/2}$ for the nonlinear case, for three different optical depths."
 The value of (007)? for the turbulent damping is shown in the right column of Figure 11.., The value of $\langle \delta v^2 \rangle^{1/2}$ for the turbulent damping is shown in the right column of Figure \ref{fig11}.
 In Figure 12. we plot the rms velocity for the collisional ancl viscous-resistive mechanisms in the left and in (he right columns. respectively.," In Figure \ref{fig12} we plot the rms velocity for the collisional and viscous-resistive mechanisms in the left and in the right columns, respectively."
 We find that the ereater the optical depth. the greater the rms velocity. which increases with decreasing wave frequency.," We find that the greater the optical depth, the greater the rms velocity, which increases with decreasing wave frequency."
 For the nonlinear and collisional mechanisms. the rms velocity can be as great as GO km | for an optically thick medium and a very low wave frequency (F=107).," For the nonlinear and collisional mechanisms, the rms velocity can be as great as 60 km $^{-1}$ for an optically thick medium and a very low wave frequency $F = 10^{-5}$ )."
 The (002)? is smallest for the case of the viscous-resistive mechanism., The $\langle \delta v^2 \rangle^{1/2}$ is smallest for the case of the viscous-resistive mechanism.
 In this work. we evaluated the role of damped Alfvénn waves in the heating of the tubes in Magnetospheric accretion models.," In this work, we evaluated the role of damped Alfvénn waves in the heating of the tubes in magnetospheric accretion models."
 We analyzed four damping mechanisms: 1) nonlinear: 2) turbulent: 3) viscous-resistive: ancl 4) collisional., We analyzed four damping mechanisms: 1) nonlinear; 2) turbulent; 3) viscous-resistive; and 4) collisional.
 Using the values f = 0.004 and F = 0.1. all the clamping lengths are found to be small compared with the length of the tube.," Using the values f = 0.004 and F = 0.1, all the damping lengths are found to be small compared with the length of the tube."
 Since our calculations show that the thermal conductivity of the gas is very small. (hese hieh-f[requency. waves. if generated only al the surface of the star. cannot heat the whole tube.," Since our calculations show that the thermal conductivity of the gas is very small, these high-frequency waves, if generated only at the surface of the star, cannot heat the whole tube."
 Extending the lrequency range of the Alfvénn waves from F = 0.1 to 10.7. we found that Che damping lengths varied strongly with frequency. spanning eight orders of magnitude for collisional damping.," Extending the frequency range of the Alfvénn waves from F = 0.1 to $10^{-5}$, we found that the damping lengths varied strongly with frequency, spanning eight orders of magnitude for collisional damping."
 The parameter / was eliminated. imposing the constraint that the Allvenic heating was converted into radiation in different optical media.," The parameter $f$ was eliminated, imposing the constraint that the Alfvenic heating was converted into radiation in different optical media."
 The damping lengths were [ound to be much smaller than the tube length. except when the lowest frequency waves had viscous-resistive damping.," The damping lengths were found to be much smaller than the tube length, except when the lowest frequency waves had viscous-resistive damping."
 Ii (his case. the damping length was greater (ian (the tube.," In this case, the damping length was greater than the tube."
 If this were the dominant mechanism. Alfvénn waves generated at the surface of the star could heat the whole tube.," If this were the dominant mechanism, Alfvénn waves generated at the surface of the star could heat the whole tube."
 For the other cases. waves produced locally are required to explain the observations.," For the other cases, waves produced locally are required to explain the observations."
 The degree of turbulence (002)/? necessary to heat the tube was constrained. Lalàng into account the optical depth of the gas.," The degree of turbulence $\langle \delta v^2 \rangle^{1/2}$ necessary to heat the tube was constrained, taking into account the optical depth of the gas."
 The larger the optical depth. the greater the degree of turbulence recquired.," The larger the optical depth, the greater the degree of turbulence required."
 We found that the turbulent velocity of the Alfvénn waves can reach LOO kins J|. for verv low frequency waves (F=5 10/7) undergoing collisional damping.," We found that the turbulent velocity of the Alfvénn waves can reach 100 km $^{-1}$, for very low frequency waves $F = 10^{-5}$ ) undergoing collisional damping."
 The authors thank an anonymous referee for useful suggestions., The authors thank an anonymous referee for useful suggestions.
 MJV would like to thank Adriano II. Cerqueira. Silvia IL. P. Alencar and Bruno V. Castilho lor very useful discussions," MJV would like to thank Adriano H. Cerqueira, Silvia H. P. Alencar and Bruno V. Castilho for very useful discussions"
that evaporation times. for given mass. are approximately independent of concentration.,"that evaporation times, for given mass, are approximately independent of concentration."
 The weak dependence. of evaporation time on density is similar re[secidensitv))., The weak dependence of evaporation time on density is similar \\ref{sec:density}) ).
 The results of our calculations. therefore. do not depend significantly on the distribution of initial concentration.," The results of our calculations, therefore, do not depend significantly on the distribution of initial concentration."
 Estimates of initial population size. are. derived assuming that an initially spherical. cluster. distribution remains spherical., 	Estimates of initial population size are derived assuming that an initially spherical cluster distribution remains spherical.
 Lowever. clusters on orbits confined near he disk evolve more rapidly than those on high inclination orbits. leaving an axisvmmetrie distribution with minimum: density about the midplane of the disk.," However, clusters on orbits confined near the disk evolve more rapidly than those on high inclination orbits, leaving an axisymmetric distribution with minimum density about the midplane of the disk."
 Estimates of initial 1alo cluster population size should not change significantly when accounting for additional cluster loss at low inclination »eause only 104 more depletion. of halo clusters. is required to account for the remaining disk clusters., Estimates of initial halo cluster population size should not change significantly when accounting for additional cluster loss at low inclination because only $\sim 10\%$ more depletion of halo clusters is required to account for the remaining disk clusters.
 Llowever. he observed disk. cluster population is about larger han our estimate of 16.," However, the observed disk cluster population is about larger than our estimate of 16."
 Assuming that our estimate [ου the initial population size scales accordingly. we estimate that the disk population may have had roughly SO members initially.," Assuming that our estimate for the initial population size scales accordingly, we estimate that the disk population may have had roughly 80 members initially."
 The choice of the mass spectrum of clusters also inlluences estimates of initial population size., 	The choice of the mass spectrum of clusters also influences estimates of initial population size.
 While use of a single power law is necessitated by sample size. the two-component power law clearly provides the best fit.," While use of a single power law is necessitated by sample size, the two-component power law clearly provides the best fit."
 To gauge the importance of this choice. we use the two-component fit and the results of Paper Ll as a guide.," To gauge the importance of this choice, we use the two-component fit and the results of Paper II as a guide."
" We estimate that the spectral index may. decrease by 0.2 for ALzM; and by 0.3 for ALSM, since the high mass range is very. similar to that considered in Paper HE while more rapid evolution occurs in the low mass range."," We estimate that the spectral index may decrease by $0.2$ for $M\gta
M_{cut}$ and by $0.3$ for $M\lta M_{cut}$ since the high mass range is very similar to that considered in Paper II while more rapid evolution occurs in the low mass range."
" This implies initial spectral indices of a,—0.3 and e»=LS and of the population Wil LAL< Adv.", This implies initial spectral indices of $\alpha_1=0.3$ and $\alpha_2=1.8$ and of the population with $M\leq M_{cut}$ .
 ln the single index fit with a=0.95. of the population has Al<Aljar.," In the single index fit with $\alpha=0.95$, of the population has $M\leq M_{cut}$."
 Since most of these clusters evaporate within a Hubble time and since they dominate the depleted. population. estimates of initial size will not change substantially whenusing a two-index moclel.," Since most of these clusters evaporate within a Hubble time and since they dominate the depleted population, estimates of initial size will not change substantially whenusing a two-index model."
accretion than is the case in one dimension.,accretion than is the case in one dimension.
 Another way of saying this is that mechanical feedback is less effective in two dimensions than one dimension: it takes more energy for a given galaxy model to reach equilibrium between cooling-driven inflow and feedback-driven /heating., Another way of saying this is that mechanical feedback is less effective in two dimensions than one dimension: it takes more energy for a given galaxy model to reach equilibrium between cooling-driven inflow and feedback-driven outflow/heating.
" At a fixed mechanical efficiency, the SMBH outflowself-regulates to larger mean accretion rates in order to provide the required additional energy."," At a fixed mechanical efficiency, the SMBH self-regulates to larger mean accretion rates in order to provide the required additional energy."
 Two-dimensional outflows are less effective than spherical ones in two critical aspects., Two-dimensional outflows are less effective than spherical ones in two critical aspects.
 less effective in protecting the SMBH from infalling gas and less effective in ejecting gas from the galaxy., less effective in protecting the SMBH from infalling gas and less effective in ejecting gas from the galaxy.
" For our models with wind efficiencies of 107? and 10-4, the distribution of Eddington ratios is in reasonable agreement with the observed fraction of galaxies above/below a given Eddington ratio over four orders of magnitude in luminosity."," For our models with wind efficiencies of $10^{-3}$ and $10^{-4}$, the distribution of Eddington ratios is in reasonable agreement with the observed fraction of galaxies above/below a given Eddington ratio over four orders of magnitude in luminosity."
 These values of the efficiency are both physically plausible and supported by recent observations (Aravetal.2011).., These values of the efficiency are both physically plausible and supported by recent observations \citep{arav:11-preprint}.
 Our plans for future work include implementing optically thick radiative transfer on dust to bring the physics of our treatment of AGN and star-formation, Our plans for future work include implementing optically thick radiative transfer on dust to bring the physics of our treatment of AGN and star-formation
Systems produciug absorption iu the spectra of distant quasars offer an excellent probe of the early Universe.,Systems producing absorption in the spectra of distant quasars offer an excellent probe of the early Universe.
 At high redshifts. they easily outuumber other observed tracers of cosmic structure. includiug both normal aud active galaxies.," At high redshifts, they easily outnumber other observed tracers of cosmic structure, including both normal and active galaxies."
 The interpretation of low column density quasar absorption systems has uudergoue somewhat of a revolution curing the past several years. with the recognitiou that they may cousist of eas aggregatingMOD into mildly uoulinear structures analogous in their dynamical structure to today’s galaxy superclusters sen1997:Hui.Cinediu.&Zhang 1997).," The interpretation of low column density quasar absorption systems has undergone somewhat of a revolution during the past several years, with the recognition that they may consist of gas aggregating into mildly nonlinear structures analogous in their dynamical structure to today's galaxy superclusters \citep{CM94,PM95,ZA95,ZA97,HK96,MC96,bi97,hui97}."
. However. damped ((DLA)) absorbers. with neutral hydrogen column densities Vy>10797em.2. are usually thought to be associated with tlie dense interstellar gas of hieh-recdshilt galaxies. based ou several lines of cireumstantial evidence: similarity between the column deusities of damped systems aud the columu densities through. typical spiral disks today. rough agreement between the total mass of atomic hydrogen in damped absorbers at 23 aud the total mass of stars today 1998).. rneasureiments of radial extents =10/1kpc in two ssvstems (Brigesetal.1089:Wolfe1993).. aud direct imagine of a number of hhosts from ground based aud HST observatious (Rao&Turushek1998:etal.2000: 1997).," However, damped ) absorbers, with neutral hydrogen column densities $\NHI \geq 10^{20.3}\;\cm^{-2}$, are usually thought to be associated with the dense interstellar gas of high-redshift galaxies, based on several lines of circumstantial evidence: similarity between the column densities of damped systems and the column densities through typical spiral disks today, rough agreement between the total mass of atomic hydrogen in damped absorbers at $z \sim 3$ and the total mass of stars today \citep{WP98}, measurements of radial extents $\ga 10\hkpc$ in two systems \citep{briggs89,wolfe93}, and direct imaging of a number of hosts from ground based and HST observations \citep{RT98,TR00,djorgovski96,fontana96,moller98,lebrun97}."
. The nature of Lyman limit (LL)) absorbers. with Nyy>10572cm2. is less well understood. though most moclels associate them with the outer regions of galaxies," The nature of Lyman limit ) absorbers, with $\NHI \geq 10^{17.2}\;\cm^{-2}$, is less well understood, though most models associate them with the outer regions of galaxies"
constant along the projected major axis of the polar ring.,constant along the projected major axis of the polar ring.
 It is only slightly declining. differently from what it is observed for spiral galaxies. where there is a very steep gradient in metallicity.," It is only slightly declining, differently from what it is observed for spiral galaxies, where there is a very steep gradient in metallicity."
 The absence of metallicity gradient is a typical behavior found in LSB galaxies (deBlok&vander 1998).. and also in the polar disk galaxy studied by Brooks and in NGC4650A (Spavoneetal..2010).," The absence of metallicity gradient is a typical behavior found in LSB galaxies \citep{deB98}, and also in the polar disk galaxy studied by \cite{Bro09} and in NGC4650A \citep{Spav10}."
. This suggests that. as already pointed out for NGC4650A (Spavoneetal.. 2010).. the star formation and the metal enrichment in the polar structure is not influenced by the stellar evolution of the central spheroid and that the polar ring was formed later.," This suggests that, as already pointed out for NGC4650A \citep{Spav10}, the star formation and the metal enrichment in the polar structure is not influenced by the stellar evolution of the central spheroid and that the polar ring was formed later."
 As also suggested by Rupkeetal.(2010).. numerical simulations predict that the absence of metallicity gradient is accounted for by inflow of low metallicity σας from the outskirts.," As also suggested by \cite{Rup10}, numerical simulations predict that the absence of metallicity gradient is accounted for by inflow of low metallicity gas from the outskirts."
 On the contrary. ordinary and oxygen-rich spiral galaxies. where no primordial gas is accreted from outside. show adecreasing abundance with increasing radii (see Fig.," On the contrary, ordinary and oxygen-rich spiral galaxies, where no primordial gas is accreted from outside, show adecreasing abundance with increasing radii (see Fig."
 7. and Pilyuginetal. 2006))., \ref{fit} and \citealt{Pil06}) ).
 These observed features in spiral disks are well explained by the infall models of galaxy formation: they predict that these systems build up through the accretion of on-site gas. which become more metal rich while it flows towards the galaxy center (Matteucci&Francois 1989:: Boissier&Prantzos 1999)).," These observed features in spiral disks are well explained by the infall models of galaxy formation: they predict that these systems build up through the accretion of on-site gas, which become more metal rich while it flows towards the galaxy center \citealt{Mat89}; \citealt{Boi99}) )."
 Such process generates the observed gradients., Such process generates the observed gradients.
 Galaxies with a polar ring/disk have a special role in the studies of physical processes at work during galaxy interactions and merging., Galaxies with a polar ring/disk have a special role in the studies of physical processes at work during galaxy interactions and merging.
 In order to account both for the featureless norphology of the central spheroidal galaxy and for the more complex structure of the polar ring/disk. so far three main formation processes have been proposed: a major dissipative nerger: tidal aceretion of material (gas and/or stars) by outside: cold accretion of pristine gas along a filament.," In order to account both for the featureless morphology of the central spheroidal galaxy and for the more complex structure of the polar ring/disk, so far three main formation processes have been proposed: a major dissipative merger; tidal accretion of material (gas and/or stars) by outside; cold accretion of pristine gas along a filament."
" In the merging scenario. the PRG results from a ""polar nerger of two disk galaxies with unequal mass. (Bekki1997:: Bekki 1998:; Bournaudetal. 2005))."," In the merging scenario, the PRG results from a “polar” merger of two disk galaxies with unequal mass, \citealt{Bek97}; \citealt{Bek98}; \citealt{Bou05}) )."
 In the accretion scenario. the polar ring/disk may form by a) the disruption of a dwarf companion galaxy orbitating around an early-type system. or by b) the tidal accretion of gas stripping from a disk galaxy outskirts. captured by an early-type galaxy on a parabolic encounter (Reshetnikov&Sotnikova1907: Bournaud&Combes 2003:: Hancocketal. 2009)).," In the accretion scenario, the polar ring/disk may form by a) the disruption of a dwarf companion galaxy orbitating around an early-type system, or by b) the tidal accretion of gas stripping from a disk galaxy outskirts, captured by an early-type galaxy on a parabolic encounter \citealt{Res97}; \citealt{Bou03}; \citealt{Han09}) )."
 Both major merger and accretion scenarios are able to account for many observed PRGs norphologies and kinematics. such as the existence of both wide and narrow rings. helical rings and double rings (Whitnoreetal.. 1990)..," Both major merger and accretion scenarios are able to account for many observed PRGs morphologies and kinematics, such as the existence of both wide and narrow rings, helical rings and double rings \citep{Whi90}. ."
 The cold accretion scenario has been proposed. very recently for the formation of a wide disk-like polar rings (see Sec. 1)):, The cold accretion scenario has been proposed very recently for the formation of a wide disk-like polar rings (see Sec. \ref{intro}) ):
 a long-lived polar structure may form through cold gas accretion along a filament. extended for ~1 Mpe. into the virialized dark matter halo (Maceidetal..2006).," a long-lived polar structure may form through cold gas accretion along a filament, extended for $\sim 1$ Mpc, into the virialized dark matter halo \citep{Mac06}."
. In this formation scenario. there ts no limits to the mass of the accreted material. thus a very massive polar disk may develop either around a stellar disk or a spheroid.," In this formation scenario, there is no limits to the mass of the accreted material, thus a very massive polar disk may develop either around a stellar disk or a spheroid."
 As suggested by the previous studies on PRGs (Spavoneet2010.. Iodiceetal.2006)). the critical physical parameters that allow to diseriminate among the three formation scenarios are 1) the total baryonie mass (stars plus gas) observed in the polar structure with respect to that in the central spheroid: 2) the kinematies along both the equatorial and meridian planes: 3) the metallicity and SER in the polar In the tidal accretion scenario the total amount of accreted gas by the early-type object is about 10% of the gas in the disk donor galaxy. ie. up to 10?M...," As suggested by the previous studies on PRGs \citealt{Spav10}, \citealt{Iod06}) ), the critical physical parameters that allow to discriminate among the three formation scenarios are 1) the total baryonic mass (stars plus gas) observed in the polar structure with respect to that in the central spheroid; 2) the kinematics along both the equatorial and meridian planes; 3) the metallicity and SFR in the polar In the tidal accretion scenario the total amount of accreted gas by the early-type object is about $10\%$ of the gas in the disk donor galaxy, i.e. up to $10^{9}M_{\odot}$."
 In the case of UGC7576 and UGC9796. the polar structure is characterized by an high baryonic mass (see Tab. 4)).," In the case of UGC7576 and UGC9796, the polar structure is characterized by an high baryonic mass (see Tab. \ref{scenarios}) ),"
 comparable with the total luminous mass in the central spheroid., comparable with the total luminous mass in the central spheroid.
 Similarly to the case of NGC4650A (Spavoneetal..2010).. the large HI mass exclude the formation of the polar structures through the tidal accretion of gas and stars by an external donor galaxy (Bournaud&Combes. 2003).," Similarly to the case of NGC4650A \citep{Spav10}, the large HI mass exclude the formation of the polar structures through the tidal accretion of gas and stars by an external donor galaxy \citep{Bou03}."
". In the merging scenario, the morphology and kinematics of the merger remnants depends on the merging initial orbital parameters and the initial mass ratio of the two galaxies (Bournaudetal.2005)."," In the merging scenario, the morphology and kinematics of the merger remnants depends on the merging initial orbital parameters and the initial mass ratio of the two galaxies \citep{Bou05}."
 In the case of UGC7576 and UGC9796. this scenario is ruled out because. according to simulations (e.g. Bournaudetal. 2005)). à high mass ratio of the two merging galaxies is required to form a massive and extended polar disk as observed in both the PRGs: this would convert the intruder into an elliptical-like. not rotationally supported. stellar system.," In the case of UGC7576 and UGC9796, this scenario is ruled out because, according to simulations (e.g. \citealt{Bou05}) ), a high mass ratio of the two merging galaxies is required to form a massive and extended polar disk as observed in both the PRGs: this would convert the intruder into an elliptical-like, not rotationally supported, stellar system."
 Since this is in contrast with the high maximum rotation velocities observed (see Reshetnikov&Combes1994 and Tab. 4) , Since this is in contrast with the high maximum rotation velocities observed (see \citealt{Res94} and Tab. \ref{scenarios}) )
in UGC7576 (~212éan/ 5) and UGC9796 (~ 157kmj/s). the merging scenario is ruled out for both galaxies.," in UGC7576 $\sim\ 212 km/s$ ) and UGC9796 $\sim\ 157 km/s$ ), the merging scenario is ruled out for both galaxies."
 As already mentioned. if the polar structure. both around an elliptical and disk galaxy. forms by the cold accretion of gas from filaments there is no limit to the acereted mass.," As already mentioned, if the polar structure, both around an elliptical and disk galaxy, forms by the cold accretion of gas from filaments there is no limit to the accreted mass."
 Moreover. due to the inflow of pristine gas. the metallicity is lower with respect to that observed in galaxies of comparable total lummosity. and its value derived by the present SFR is higher than those directly measured by the chemical abundances.," Moreover, due to the inflow of pristine gas, the metallicity is lower with respect to that observed in galaxies of comparable total luminosity, and its value derived by the present SFR is higher than those directly measured by the chemical abundances."
 Spavoneetal.(2010) have found that such predictions were consistent with observations in the polar disk galaxy NGC46504. yielding to the conclusion that the cold accretion of gas by cosmic web filaments is the most realistic formation scenario for this object.," \cite{Spav10} have found that such predictions were consistent with observations in the polar disk galaxy NGC4650A, yielding to the conclusion that the cold accretion of gas by cosmic web filaments is the most realistic formation scenario for this object."
 By studying the chemical abundances in the polar structure of UGC7576 and UGC9796. in the present work we aim at testing the cold acecretion scenario for these objects.," By studying the chemical abundances in the polar structure of UGC7576 and UGC9796, in the present work we aim at testing the cold accretion scenario for these objects."
 The main results are as follows (see also Sec. ??)):, The main results are as follows (see also Sec. \ref{disc}) ):
 1) both PRGs have on average lower metallicity with. respect to that of same- spiral disks. in particular. for UGC7576Z=0.4Z.. and for UGC9796 Z= 0.1Z4: 2) the metallicity expected for the present SER at three different epochs are higher than those measured from the element abundances. and they varies 1n the range 0.5Z.<Z€IZ» for UGC7576 and," 1) both PRGs have on average lower metallicity with respect to that of same-luminosity spiral disks, in particular, for UGC7576$Z=0.4 Z_{\odot}$ and for UGC9796 $Z=0.1 Z_{\odot}$ ; 2) the metallicity expected for the present SFR at three different epochs are higher than those measured from the element abundances, and they varies in the range $0.5 Z_{\odot} \leq Z \le 1 Z_{\odot}$ for UGC7576 and"
The sightlines for this study were chosen to maximize the possibility of finding enhanced © depletions as compared to the diffuse cloud value.,The sightlines for this study were chosen to maximize the possibility of finding enhanced C depletions as compared to the diffuse cloud value.
 In the sample of interstellar oxvgen measurements made by Cartledgeοἱal.(2001).. the sightlines toward IID 27778. IID 37021. Η 37061. WD 147588. WD 152590 and LD 207198 appeared to show iilel to moderate enhancements of oxveen depletions as compared to (he average diffuse cloud. abundance (Meveretal.1993).," In the sample of interstellar oxygen measurements made by \citet{car01}, the sightlines toward HD 27778, HD 37021, HD 37061, HD 147888, HD 152590 and HD 207198 appeared to show mild to moderate enhancements of oxygen depletions as compared to the average diffuse cloud abundance \citep{mey98}."
. The O depletion toward HD 152590 in Cartledge et al., The O depletion toward HD 152590 in Cartledge et al.
 was determined wilh respect to Ixr since no Il measurement was available: a recent FUSE II determination (Cartledgeοἱal.2003) shows that Ο/Η toward: (his star is actually consistent with the dilfuse cloud value.," was determined with respect to Kr since no H measurement was available; a recent $FUSE$ H determination \citep{cml03}
 shows that O/H toward this star is actually consistent with the diffuse cloud value."
 A summary of the sightline eharacteristies for our sample is given in Table 1., A summary of the sightline characteristics for our sample is given in Table 1.
 The intrinsie weakness of the 2325 CC uf feature that is the focus of our program mandated Chat the observations be high resolution and have high sienal-(o-uoise ratios (S/N 2 50 - 140 per pixel)., The intrinsic weakness of the 2325 C ] feature that is the focus of our program mandated that the observations be high resolution and have high signal-to-noise ratios (S/N $\approx$ 50 - 140 per pixel).
 To satisfy. these requirements. the data were all obtained with the echelle mode of the STIS instrument aboard the (~ 2.5 kin/s resolution).," To satisfy these requirements, the data were all obtained with the echelle mode of the STIS instrument aboard the $\sim$ 2.5 km/s resolution)."
 Four spacecraft orbits were allocated to each (target., Four spacecraft orbits were allocated to each target.
 Data lor the first observed target. IID 207193. were all obtained through the 0 72 x 0709 aperature at a central wavelength of 2263A.," Data for the first observed target, HD 207198, were all obtained through the 0 2 x 0 09 aperature at a central wavelength of 2263."
. The S/N of the data was below what would be expected from photon statistics. so the observing strategy. was changed [or the other targets.," The S/N of the data was below what would be expected from photon statistics, so the observing strategy was changed for the other targets."
 The data for WD 27778. WD 37061. WD 1437888 and WD 152590 were obtained through the four 0 72 x 0 72 FP-SPLIT apertures and were all centered at 2263As: the FPA aperture observations were shorter Chan the others because of target acquisition in the first orbit.," The data for HD 27778, HD 37061, HD 147888 and HD 152590 were obtained through the four 0 2 x 0 2 FP-SPLIT apertures and were all centered at 2263; the FPA aperture observations were shorter than the others because of target acquisition in the first orbit."
 The possibility of a nearby contaminating star made the use of the EP-SPLIT apertures impractical for the IID 37021 observations., The possibility of a nearby contaminating star made the use of the FP-SPLIT apertures impractical for the HD 37021 observations.
 Instead. those data were all obtained through STIS's 072 x 0 709 aperture but al several central wavelength settings: the first two orbits centered αἱ 2263 aand one orbit each centered at 2313 and 2363Ἀ.," Instead, those data were all obtained through STIS's 0 2 x 0 09 aperture but at several central wavelength settings; the first two orbits centered at 2263 and one orbit each centered at 2313 and 2363."
. In order to obtain the most reliable measurements of the weak carbon line. the data were iveated independenily at the three institutions represented by the authors: LSU. Northwestern and Whitman.," In order to obtain the most reliable measurements of the weak carbon line, the data were treated independently at the three institutions represented by the authors: LSU, Northwestern and Whitman."
 For the LSU analysis. the data used were (hose calibrated and delivered by the SIScl pipeline just alter each observation.," For the LSU analysis, the data used were those calibrated and delivered by the STScI pipeline just after each observation."
 The Northwestern and Whitman analvses were based on spectra delivered by the STScl archive at the end of October 2003 with on-the-fly calibrations., The Northwestern and Whitman analyses were based on spectra delivered by the STScI archive at the end of October 2003 with on-the-fly calibrations.
 We recalibrated some spectra at Nortliwestern using the latest STSDAS (version 3.1) varying the extraction slit width to test whether this produced any significant change to the data: it did not., We recalibrated some spectra at Northwestern using the latest STSDAS (version 3.1) varying the extraction slit width to test whether this produced any significant change to the data; it did not.
 Each institution independently performed a weighted coaddition of the individual observations for each target with either IRAF or IDL software., Each institution independently performed a weighted coaddition of the individual observations for each target with either IRAF or IDL software.
 The final spectra, The final spectra
ln the absence of any heating mechanism. a sclf-eravitating aceretion dise will cool down until it becomes gravitationally unstable. which happens when the stability parameter Q llere. ὃς is the sound speed. s denotes the epievelic [requeney. and. XM is the surface density of the disc.,"In the absence of any heating mechanism, a self-gravitating accretion disc will cool down until it becomes gravitationally unstable, which happens when the stability parameter $Q$ Here, $\cs$ is the sound speed, $\kappa$ denotes the epicyclic frequency, and $\Sigma$ is the surface density of the disc."
 When Q approaches unity ancl gravitational instabilities kick in. spiral shocks are generated that heat up the disc(??).," When $Q$ approaches unity and gravitational instabilities kick in, spiral shocks are generated that heat up the disc."
. An equilibrium can then be set up in which shock heating exactly balances energy losses through cooling(?)., An equilibrium can then be set up in which shock heating exactly balances energy losses through cooling.
. This is called the eravito-turbulent state of the disc., This is called the gravito-turbulent state of the disc.
 Lo the cooling time scale is parametrised as with 2j constant and © the local angular velocity of the disc. a local balance of heating and cooling dictates that the spiral shocks generate an elfective a parameter where 5 is the acliabatic exponent.," If the cooling time scale is parametrised as with $\beta$ constant and $\Omega$ the local angular velocity of the disc, a local balance of heating and cooling dictates that the spiral shocks generate an effective $\alpha$ parameter where $\gamma$ is the adiabatic exponent."
 | lot. of work has been put into determining whether a local equilibrium really is established and whether an a-description is valid at all for eravitoturbulent disces(222). which is not necessarily the case(?).," A lot of work has been put into determining whether a local equilibrium really is established and whether an $\alpha$ -description is valid at all for gravitoturbulent discs, which is not necessarily the case."
. For strong enough cooling. the disc is unable to provide enough shock heating to balance the cooling. and the disc fragments (2).. possibly forming gas giant planets.," For strong enough cooling, the disc is unable to provide enough shock heating to balance the cooling, and the disc fragments , possibly forming gas giant planets."
 The critical value of 3 was found to be 3.z3 in for 2D simulations with ?=2., The critical value of $\beta$ was found to be $\betac\approx 3$ in for 2D simulations with $\gamma=2$.
 interpreted the fragmentation criterion as a maximum stress the disc can sustain., interpreted the fragmentation criterion as a maximum stress the disc can sustain.
 They. found: 0.06. which Leads to a value of 3 that depends on  through equation (3)).," They found $\alpha_\mathrm{max}\approx 0.06$ , which leads to a value of $\betac$ that depends on $\gamma$ through equation \ref{eqabc}) )."
 Lt was found by that the value of Guus can increase by a factor of z2 when 2 is decreased slowly., It was found by that the value of $\alpha_\mathrm{max}$ can increase by a factor of $\approx 2$ when $\beta$ is decreased slowly.
 The determination of the critical value of 3 is important in models of giant planet formation that rely on. disc instability22)., The determination of the critical value of $\beta$ is important in models of giant planet formation that rely on disc instability.
. Since according to the results above. the cooling time scale should be of the order of the orbital time scale. planet. formation through disc instability is more likely to occur in the outer parts of real protoplanetary disces(?7).," Since according to the results above, the cooling time scale should be of the order of the orbital time scale, planet formation through disc instability is more likely to occur in the outer parts of real protoplanetary discs."
. Alore recently. the numerical convergence of the results on te was questioned. byο," More recently, the numerical convergence of the results on $\betac$ was questioned by."
"ν, Phere. it was shown through SPLH simulations that the critical value of 2 had not converged with respect to resolution."," There, it was shown through SPH simulations that the critical value of $\beta$ had not converged with respect to resolution."
 At higher ancl higher resolution. discs with higher and higher :2would fragment. with no sien of convergence.," At higher and higher resolution, discs with higher and higher $\beta$would fragment, with no sign of convergence."
 Ht was suggested by that this behaviour, It was suggested by that this behaviour
"(SSAP, ?)) and the Theoretical Spectra Access Protocol (TSAP).","(SSAP, \citealt{tody08}) ) and the Theoretical Spectra Access Protocol (TSAP)."
 The SSAP defines a uniform interface to track and access 1-D spectra., The SSAP defines a uniform interface to track and access 1-D spectra.
" However, being originally designed to access observational data, it cannot be used directly to query the POLLUX database."," However, being originally designed to access observational data, it cannot be used directly to query the POLLUX database."
 A way to work around the impossibility to locate (with coordinates on the sky) the synthetic spectra is to use the FORMAT=METADATA mechanism to query theoretical data in the context of the SSAP., A way to work around the impossibility to locate (with coordinates on the sky) the synthetic spectra is to use the FORMAT=METADATA mechanism to query theoretical data in the context of the SSAP.
" This corresponds to the TSAP, that has been included in the SSAP as a usecase!!,, and that allows to make a query and visualize data from the POLLUX database using the VOSpec tool as shown in Fig."," This corresponds to the TSAP, that has been included in the SSAP as a use, and that allows to make a query and visualize data from the POLLUX database using the VOSpec tool as shown in Fig."
 5 for the POLLUX MARCS solar spectrum in units of absolute flux versus wavelength., \ref{VOspec-sun} for the POLLUX MARCS solar spectrum in units of absolute flux versus wavelength.
" All the relevant characteristics of the data, appearing in the POLLUX query forms, as well as those listed in the header files coming with any POLLUX spectra, have been associated to an UnifiedContent Descriptor (UCD!”,, that ensures the data interoperability within the VO context."," All the relevant characteristics of the data, appearing in the POLLUX query forms, as well as those listed in the header files coming with any POLLUX spectra, have been associated to an UnifiedContent Descriptor \citep[UCD\footnote{http://www.ivoa.net/Documents/latest/UCD.html}, that ensures the data interoperability within the VO context."
" Finally, the Pollux database has been registered as a VO service providing theoretical spectra that can be handled, through VO tools, within the framework of Science Cases."," Finally, the Pollux database has been registered as a VO service providing theoretical spectra that can be handled, through VO tools, within the framework of Science Cases."
 In particular two scientific applications are already in progress :, In particular two scientific applications are already in progress :
features of the accretion flow around the black hole iu «üffeveut states.,features of the accretion flows around the black hole in different states.
 The observed N-ray spectra in differeut spectral stedes can be explained by the Comptonization ( ¢.e.. soft photons from the cold thin accretion disk are Coniptonized iu a hot corona) with different model parameters (Nowaketal.2002:Corolicmct 9005). N," The observed X-ray spectra in different spectral states can be explained by the Comptonization ( e.g., soft photons from the cold thin accretion disk are Comptonized in a hot corona) with different model parameters \citep{n02,c03}. ."
aravan(1996) proposed hat these different spectral states can be explained bv outer thin disk plus iuner advection «onünated accretion How (ADAF) models with varviug accretion rate aud transition radius (Exineal.1997.1998).," \citet{n96} proposed that these different spectral states can be explained by outer thin disk plus inner advection dominated accretion flow (ADAF) models with varying accretion rate and transition radius \citep{e97,e98}."
. Iu his scenario. the transition radius increases with decreasingmass accretion rate ip (i=M/ η Milleretal.," In this scenario, the transition radius increases with decreasingmass accretion rate $\dot{m}$ $\dot{m}=\dot{M}/\dot{M}_{\rm Edd}$ )."
(2001.206) sugeested that a standard thin disk may extend at or near to the innermost stable circular orbi (ISCOJ. at least iu xieht phases of the lowliuL state. after i6 spectra of CUN. 339-1 observed by Chandra luring its decline of 2002-203 outburst aud by NADLNewton during its 2001 outburst with a cool in disk component and a relativistic Fe I line nodel.," \citet{m04,m06} suggested that a standard thin disk may extend at or near to the innermost stable circular orbit (ISCO), at least in bright phases of the low/hard state, after the spectra of GX 339-4 observed by Chandra during its decline of 2002-2003 outburst and by XMM-Newton during its 2004 outburst with a cool thin disk component and a relativistic Fe K line model."
 that the «isk reflection inodel can also fit the spectra in fjio hard state when ιο total1-100. keV luminosity even lower than πια (Tomsicletal.2XWOs}., that the disk reflection model can also fit the spectra in the hard state when the total1-100 keV luminosity even lower than $0.023L_{\rm Edd}$ \citep{t08}.
. Moreover. the orescence of jet/outflow in the accretion flows to explain the radio clussion aud he radio/X-rav correlations in the low/hard state (Feuderetal.1999:Corbel2000:Markotffctal.2003:C'orbelet 2003).," Moreover, the presence of jet/outflow in the accretion flows to explain the radio emission and the radio/X-ray correlations in the low/hard state \citep{f99,c00,m03,corbel03}."
. Fluorescent Fe Ke πιο enisson ds a vorv powerful trace of the iuner οσο CR) of the thin accretion disk around the black holes., Fluorescent Fe $\alpha$ line emission is a very powerful trace of the inner edge $R_{\rm in}$ ) of the thin accretion disk around the black holes.
 Tonisicketal.(2009) used. Suzaku acl RATE observations of durimg 2008 to study the iron line profile at low lhunünositv state when the total 1-100 keV Iuuimositv is 0.001πια. aud showed that the inner edee of the thin disk recedes aid the truncated accretion disk is present around a black hole at low huuinositv.," \citet{t09} used Suzaku and RXTE observations of during 2008 to study the iron line profile at low luminosity state when the total 1-100 keV luminosity is $0.0014 L_{\rm Edd}$, and showed that the inner edge of the thin disk recedes and the truncated accretion disk is present around a black hole at low luminosity."
 If the ADAE exists in the inner region of: the aceretion flow as stgeested by Naravan(1996).. thermal N-rav iron lines may also enütted from the hot plasma in the ADAF surouudiug the black hole at very low huniuositv (Naravan&Ravinond1999:Pernaetal.2000:Xu2006:&Cao 2009)...," If the ADAF exists in the inner region of the accretion flow as suggested by \citet{n96}, thermal X-ray iron lines may also emitted from the hot plasma in the ADAF surrounding the black hole at very low luminosity \citep{n99,p00,xu06,xu09}. ."
 We describe the calculation ofdwuamucal structures of the ADAF in 82., We describe the calculation ofdynamical structures of the ADAF in 2.
 The calculation of the equivalent, The calculation of the equivalent
"where A4,(0) is the initial stellar mass. 5h(/)=fatodi. ancl 3 is basically the accretion clock which depends on the eas content and on JUS but not on time.","where $\msi $ is the initial stellar mass, $h(t) = \int g(t) dt$ , and $\beta$ is basically the accretion clock which depends on the gas content and on $\Rclus^{3\eta-\zeta}$ but not on time."
 Phe primary dillerence between this case and the preceding ticallohe accretion is the higher dependency on the stellar mass such that the initial mass is always important., The primary difference between this case and the preceding tidal-lobe accretion is the higher dependency on the stellar mass such that the initial mass is always important.
 This is the case for any accretion rate which depends on the stellar mass to a power greater than one., This is the case for any accretion rate which depends on the stellar mass to a power greater than one.
 The resulting mass spectrum from such an accretion rate was studied. by Zinnecker (1982). and. we follow his derivation here., The resulting mass spectrum from such an accretion rate was studied by Zinnecker (1982) and we follow his derivation here.
" As there is a direct. correlation between initial ancl final mass. then there is a direct mapping from the initial mass spectrum to the final one. lnverting equation(28)) and inserting it into equation(20)). as long as 3 is independent. of the initial masses. we have In the limit 2M,(D)>>1 (CAL,- AL.(O)) this gives and thus as long as there is an initial (small) range of stellar masses. the asvmptotic limit of the mass spectrum is This is somewhat complicated if there is a correlation between the initial masses ancl radii in the cluster."," As there is a direct correlation between initial and final mass, then there is a direct mapping from the initial mass spectrum to the final one, Inverting \ref{maccbhsol}) ) and inserting it into \ref{imfmap}) ), as long as $\beta$ is independent of the initial masses, we have In the limit $\beta \ms h(t) >> 1$ $\ms >> \msi$ ) this gives and thus as long as there is an initial (small) range of stellar masses, the asymptotic limit of the mass spectrum is This is somewhat complicated if there is a correlation between the initial masses and radii in the cluster."
" Ln this case S ds dependent on AZ,(0) and we cannot derive an analytical expression for the resulting mass spectrum.", In this case $\beta$ is dependent on $\msi$ and we cannot derive an analytical expression for the resulting mass spectrum.
 In the case where the incoming mass distribution is from the tidal-lobe accretion. then From 83.1 we have an expression for the initial mass-raclius correlation (equation 17)).," In the case where the incoming mass distribution is from the tidal-lobe accretion, then from 3.1 we have an expression for the initial mass-radius correlation (equation \ref{massvsrad1}) )."
 Using this. we can generate a mass spectrum from accretion for a cluster in anx27 configuration.," Using this, we can generate a mass spectrum from accretion for a cluster in a $n\propto R^{{-2}}$ configuration."
 Numerical simulations of accretion in stellar dominated. clusters have values of 5j=1/3 and &=3/2 (Bonnell 2000)., Numerical simulations of accretion in stellar dominated clusters have values of $\eta = 1/3$ and $\zeta = 3/2$ (Bonnell 2000).
" Equation (27)) then becomes Figure 2 shows the resultant IME for clusters accreting from an initial delta function mass distribution ancl those with an initially M,2 mass spectrum that accretes according to equation 2 (33)).", Equation \ref{maccbh}) ) then becomes Figure \ref{bhimf3o2} shows the resultant IMF for clusters accreting from an initial delta function mass distribution and those with an initially $\ms^{-3/2}$ mass spectrum that accretes according to equation \ref{maccbhseg}) ).
 In the former case. the stars. in anoxr7 distribution. accrete with a spatiallv-indepencdoent accretion rate (M.xMZ).," In the former case, the stars, in a $n\propto r^{-2}$ distribution, accrete with a spatially-independent accretion rate $\macc \propto \ms^2$ )."
" These clusters develop. an Al,= guass spectrum.", These clusters develop an $\ms^{-2}$ mass spectrum.
 In contrast. with the initial mass segregation. the ΙΔΗΣ develops into a steeper (52— 2.5) IME as the stars acerete in the stellar dominated: regime.," In contrast, with the initial mass segregation, the IMF develops into a steeper $\gamma \approx -2.5$ ) IMF as the stars accrete in the stellar dominated regime."
" This is somewhat steeper than the analytical AZ,7 mass spectrum due to the initial correlation between stellar mass ancl position in the cluster.", This is somewhat steeper than the analytical $\ms^{{-2}}$ mass spectrum due to the initial correlation between stellar mass and position in the cluster.
 We expect stellar dominated clusters to accrete towards a mass spectrum with asvmptotic limits of 2.5τσx2., We expect stellar dominated clusters to accrete towards a mass spectrum with asymptotic limits of $-2.5 \simless \gamma \le -2$.
 ‘This mass spectrum is expected for the higher mass stars as it is they that are located in thecentre of the cluster and have attained a higher mass due to their previous acecretion in the eas dominated regime., This mass spectrum is expected for the higher mass stars as it is they that are located in thecentre of the cluster and have attained a higher mass due to their previous accretion in the gas dominated regime.
 Clusters that have mass, Clusters that have mass
formation take place eutirelv at lower redshifts. since this would violate observational limits from umber counts and redshift survers.,"formation take place entirely at lower redshifts, since this would violate observational limits from number counts and redshift surveys."
 Hf the galaxies were indecel μια] at hieh-redshüft. they have had to erow by mergers (ellipticals). or star formation in their disks has had to proceed from the iusicde out (spirals).," If the galaxies were indeed small at high-redshift, they have had to grow by mergers (ellipticals), or star formation in their disks has had to proceed from the inside out (spirals)."
 Whatever the mechanisin bv which galaxies “hide” at high redshift. it is Iucreasinely more difficult to conceal them iu laree-volune universes. such as open or A universes with a low deceleration.1.," Whatever the mechanism by which galaxies “hide” at high redshift, it is increasingly more difficult to conceal them in large-volume universes, such as open or $\lambda$ universes with a low deceleration,."
. For those cosmological models the total uuuber ofobserved galaxies 1s less than predicted. respective of size.," For those cosmological models the total number of observed galaxies is less than predicted, irrespective of size."
 Star formation would then have to be episodal to account for the observed uumuber of objects; aud more of the star formation would have to take place behind a shroud of dust.," Star formation would then have to be episodal to account for the observed number of objects, and more of the star formation would have to take place behind a shroud of dust."
 Αν realistic model of galaxy formation will have to face this issue. particularly if other observational evidence contiuues to point toward open or A models. as it now inereasinely docs.," Any realistic model of galaxy formation will have to face this issue, particularly if other observational evidence continues to point toward open or $\lambda$ models, as it now increasingly does."
 Further comments are Ooeiven in the concludingC» remarks to this coufereuce., Further comments are given in the concluding remarks to this conference.
 Based on the experience of the last few vears we expect most progress to come from observations., Based on the experience of the last few years we expect most progress to come from observations.
 Deep near-infrared ΠαρΊο is verv useful to improve the photometric redshifts of all fainter galaxies. and is esseutial for galaxies at redshifts 2z5.," Deep near-infrared imaging is very useful to improve the photometric redshifts of all fainter galaxies, and is essential for galaxies at redshifts $z \ga 5$."
 Spectroscopy of Daliner Hues can determine the dust content of the galaxies. telling us how wich star formation is hidden from UV view.," Spectroscopy of Balmer lines can determine the dust content of the galaxies, telling us how much star formation is hidden from UV view."
 It will also confirma spectroscopic redshifts in the range 1«:<2. which have not con. adequately observed from the erouud.," It will also confirm spectroscopic redshifts in the range $1 < z < 2$, which have not been adequately observed from the ground."
 Deep UST imaging in intermediate xuads. especially atTOOOA.. can improve the photometric redshifts aud increase heir reliability at fainter magnitudes.," Deep HST imaging in intermediate bands, especially at, can improve the photometric redshifts and increase their reliability at fainter magnitudes."
 Finallv. better nuage reconstruction should inprove both resolution aud photometry. enabling more accurate flux ucasureineut aud the detection of faimter objects.," Finally, better image reconstruction should improve both resolution and photometry, enabling more accurate flux measurement and the detection of fainter objects."
 Ou the theoretical side. modeling of galaxy formation is ouly beginning seriously to tackle eas dvuamics and star formation.," On the theoretical side, modeling of galaxy formation is only beginning seriously to tackle gas dynamics and star formation."
" Expect important breakthroughs 100,", Expect important breakthroughs here.
 Iu anv event. we can no lounger coutent ourselves with the traditional ΠΠΡΟ counts; Va). [IN(t). or even /NGn.τὸν," In any event, we can no longer content ourselves with the traditional number counts, $N(m)$, $N(z)$, or even $N(m,z)$."
 The sizes of galaxies at different wavelengths. or better vet their profiles. are observable aud provide crucial constraints both to galaxy formation aud to the underlying cosmological model.," The sizes of galaxies at different wavelengths, or better yet their profiles, are observable and provide crucial constraints both to galaxy formation and to the underlying cosmological model."
 We need to obtain these data and to confront our models with them., We need to obtain these data and to confront our models with them.
 This work was supported in part by NASA eraut. AR-O7551.01-96A., This work was supported in part by NASA grant AR-07551.01-96A.
dvuaimics. stellar population svuthesis (SPS) stellar mass cstimates APPS edven an IMF. and WL shear lueasurements to constrain a model for the relationship between stellar and virial mass;,"dynamics, stellar population synthesis (SPS) stellar mass estimates $_{*}^{\rm SPS}$ given an IMF, and WL shear measurements to constrain a model for the relationship between stellar and virial mass."
 We allow for several different DM halo profiles and we fit for an IME nüsnmateh pavaneter relating to MEPS (Trenet.al. 2010)., We allow for several different DM halo profiles and we fit for an IMF mismatch parameter relating $_*$ to $_*^{\rm SPS}$ \citep{t10}.
. The two main improvements over the analysis by Treuctal.(2010) ave the inclusion of WL data aud the adoption of more flexible models., The two main improvements over the analysis by \citet{t10} are the inclusion of WL data and the adoption of more flexible models.
 We use a staudard fat ACDAL cosmology with O4;=0.3 and h—0.7., We use a standard flat $\Lambda$ CDM cosmology with $\Omega_{\rm M} = 0.3$ and $h=0.7$.
" Our gsuuple is derived from the Sloan Lens ACS Survey (SLACS:Boltonetal.2006:Cavazzi2007:al.2008:Augeret 20001, which consists of nearly LOO strong gravitational leusiug galaxies."," Our sample is derived from the Sloan Lens ACS Survey \citep[SLACS;][]{slacsi,slacsiv,slacsv,slacsix}, which consists of nearly 100 strong gravitational lensing galaxies."
 We exclude 6 ealaxies which appear to be structurally different than the rest of the SLACS carly-type lenses (Augeretal. 2010aj., We exclude 6 galaxies which appear to be structurally different than the rest of the SLACS early-type lenses \citep{slacsx}.
. These svstenis are significant outliers from the otherwise very fieht relation between effective radius. velocity dispersion. total central mass. aud stellar mass: this is likely due to systematic errors in the stellar velocity dispersious of these svstenis (alsoseeJiang&ISochanck 2007).," These systems are significant outliers from the otherwise very tight relation between effective radius, velocity dispersion, total central mass, and stellar mass; this is likely due to systematic errors in the stellar velocity dispersions of these systems \citep[also see][]{jiang}."
. We therefore use a stibsample of 53 ETCs that have well-determuned central stella velocity dispersions. SPS stellar masses assume Chabrier(2003) and Salpeter(1955). IMES. and ceutral projected mass estimates from SL.," We therefore use a subsample of 53 ETGs that have well-determined central stellar velocity dispersions, SPS stellar masses assuming \citet{chabrier} and \citet{salpeter} IMFs, and central projected mass estimates from SL."
 We also use WL shear mieasuremoeuts for a subset of Ll lenses with deep ACS imaging., We also use WL shear measurements for a subset of 44 lenses with deep ACS imaging.
 The details of how these data are derived cau be fouud elsewhere (Gavazzietal.2007:Boltou2008:Augeretal.2009. 2010a).," The details of how these data are derived can be found elsewhere \citep{slacsiv,slacsv,slacsix,slacsx}."
. We eiipliasize that the SLACS Ίος ealaxies are indistinemshable from twin ETCs selected from SDSS to havethe same velocity dispersion aud redshift aud can therefore be considered as represcutative of the overall population of massive ETCs (Trenetal.2006.2009:Augerctal.2009. 2010a).," We emphasize that the SLACS lens galaxies are indistinguishable from twin ETGs selected from SDSS to havethe same velocity dispersion and redshift and can therefore be considered as representative of the overall population of massive ETGs \citep{slacsii,slacsviii,slacsix,slacsx}."
. We ainodel the leuses as two-coniponent lass distributions consisting of stars and a CDAL halo., We model the lenses as two-component mass distributions consisting of stars and a CDM halo.
 Both conrponeuts are modeled as spherical since our data do iof constrain the halo fatteniug and our dyvuamical analvsis ds restricted. to spherical Jeans modeling., Both components are modeled as spherical since our data do not constrain the halo flattening and our dynamical analysis is restricted to spherical Jeans modeling.
" Comparisons with nou-spherical models shows that the nost relevant quantities such as mass-density slope aud otalinass are relatively inscusitive to this approximation ίοιο,,Barnabeetal.2009).."," Comparisons with non-spherical models shows that the most relevant quantities such as mass-density slope and total mass are relatively insensitive to this approximation \citep[e.g.,][]{barnabe}."
" The stars are modeled with the Ieruquist(1990) motile aud the scale radius is set bv the rest-frame V-baud effective radius. às,=0.5510: our conclusions are unchauged if we adopt a Jaffe(1983) profile."," The stars are modeled with the \citet{hernquist} profile and the scale radius is set by the rest-frame $V$ -band effective radius, $r_a =
0.551 r_{\rm e}$; our conclusions are unchanged if we adopt a \citet{jaffe} profile."
 We assunie isotropic orbits. consistent with typical findings or massive ETCs (Dirnabóetal.2009).," We assume isotropic orbits, consistent with typical findings for massive ETGs \citep{barnabe}."
. MBld radial anisotropy (dl&oopiuansetal.2009) onlv altes the interred velocity dispersous bv a few percent. ie. au ainount simaller than the errors. aud therefore does not change our conclusious.," Mild radial anisotropy \citep{k09} only alters the inferred velocity dispersions by a few percent, i.e. an amount smaller than the errors, and therefore does not change our conclusions."
 We investigate three different density profiles for the CDAL halos., We investigate three different density profiles for the CDM halos.
 First. we consider the NEW profile with the concentration fixed by e{Ala.)x followine Maccioetal.(2008).. neglecting intrinsic ALscatter for conrputational simplicity.," First, we consider the NFW profile with the concentration fixed by $c({\rm M_{\rm vir}}) \propto {\rm M}_{\rm vir}^{-0.094}$ following \citet{maccio}, neglecting intrinsic scatter for computational simplicity."
 We also modify the NEW halo to include AC’ using the prescriptions of BSG aud GOL. with the barvon fraction set by Equation 1 and the scale radius given by rq: these Έπος profiles are the most widely adopted descriptions for CDM. halos (c.¢..Jiang2010).. aud the NEW (Ba6) halo has the lowest (highest) central DM deusity.," We also modify the NFW halo to include AC using the prescriptions of B86 and G04, with the baryon fraction set by Equation 1 and the scale radius given by $r_a$; these three profiles are the most widely adopted descriptions for CDM halos \citep[e.g.,][]{jiang,schulz,napolitano}, and the NFW (B86) halo has the lowest (highest) central DM density."
 We fix the relationship between stellar aud virial 11ass to have the form suggested by Mosteretal.(2010). (who use a IE&roupa IME) based upon abundance matching techniques.We assuue where ME is the stellar mass inferred from the SPS models assuiiug a Salpeter IME. 54=7.17. 5»=0.201. and 30.557 and we fit for y aud Mago.," We fix the relationship between stellar and virial mass to have the form suggested by \citet{moster} (who use a Kroupa IMF) based upon abundance matching We assume _* = where $_*^{\rm SPS}$ is the stellar mass inferred from the SPS models assuming a Salpeter IMF, $\gamma_1 = 7.17$, $\gamma_2 = 0.201$, and $\beta = 0.557$ and we fit for $_{*,0}$ and $_{\rm vir,0}$."
 We note that our results are qualitatively the same if we use different values for 541.5». aud (e.e.. fourDuttonctal.2010).. although the eoocduess-offit changes: we will explicitly fit for these paraimcters in a future paper to ascertain the extent to which one model prefers a giveu set of parameters over another.," We note that our results are qualitatively the same if we use different values for $\gamma_1,
\gamma_2,$ and $\beta$ \citep[e.g., from][]{dutton}, although the goodness-of-fit changes; we will explicitly fit for these parameters in a future paper to ascertain the extent to which one model prefers a given set of parameters over another."
 We also allow for an offset between the model stellar mass and the SPS stellar mass. equivalent to a imasdependent TIME. (referred. to as the “Free? TAIF inodel).," We also allow for an offset between the model stellar mass and the SPS stellar mass, equivalent to a mass-dependent IMF (referred to as the `Free' IMF model)."
 The offset has the form ΕΝ... here a is equivalent to that introduced by Treuetal.(32010) ancl is equivaleut to the slope of Equation 5 in ," The offset has the form = -, where $\alpha$ is equivalent to that introduced by \citet{t10} and $\eta$ is equivalent to the slope of Equation 5 in \citet{t10}."
We adopt the MBPS derived by Augerctal. (2009).. although our results are robust with respect to changes in the SPS models aud choice of priors (Treuetal.2010) due to the relative simplicity of the older stellar populations of the SLACS ETGs.," We adopt the $_*^{\rm SPS}$ derived by \citet{slacsix}, although our results are robust with respect to changes in the SPS models and choice of priors \citep{t10} due to the relative simplicity of the older stellar populations of the SLACS ETGs."
 Uncertaiuties in SPS uniodels of evolved galaxies (Conroy&Camu2010:Alarastouetal.2009) iuaw load to systematic offsets iu ADS of order ~0.05 dex (o.@..Treuetal.2010).," Uncertainties in SPS models of evolved galaxies \citep{conroy,maraston} may lead to systematic offsets in $_*^{\rm SPS}$ of order $\sim 0.05$ dex \citep[e.g.,][]{t10}."
". We perform a Markov Chain Monte Carlo simulation to fit the models to the data: we propose ALi, for cach lens aud use Equation 1 (and 2 for the Free IME model) to predict a stellar mass.", We perform a Markov Chain Monte Carlo simulation to fit the models to the data; we propose $_{\rm vir}$ for each lens and use Equation 1 (and 2 for the Free IMF model) to predict a stellar mass.
 These masses normalize the halo and bulec. aud we use these profiles to predict the stellar kinematics. SL. aud WL signals for cach galaxy.," These masses normalize the halo and bulge, and we use these profiles to predict the stellar kinematics, SL, and WL signals for each galaxy."
 We conrpare these quantities aud the model stellar masses with the observed quantities ona ealaxy-by-ealaxy basis (i.c.. we do not bin or average measurements. mchudius the WL) to optimize the free model parameters and determine eooduess-offit values.," We compare these quantities and the model stellar masses with the observed quantities on a galaxy-by-galaxy basis (i.e., we do not bin or average measurements, including the WL) to optimize the free model parameters and determine goodness-of-fit values."
 The best-fit poarznueters and 47 for each model are listed am Table 1.. and we show the best-fit uass distributions for a characteristic leas system εν= LOLS) in Figure 1.," The best-fit parameters and $\chi^2$ for each model are listed in Table \ref{T_results}, and we show the best-fit mass distributions for a characteristic lens system $L_V = 10^{10.85} L_\odot$ ) in Figure \ref{F_bulgehalo}."
 This illustrates that the low normalization of the Chabrier IME requires a more massive CDM halo to increase the central projected mass and be consistent with the mass required by the SL data (this is true for the NEW. AC models)., This illustrates that the low normalization of the Chabrier IMF requires a more massive CDM halo to increase the central projected mass and be consistent with the mass required by the SL data (this is true for the NFW AC models).
 However. increasing the halo wiass also inercascs the scale radius because the virial radius increases aud the concentration becomes smaller (6.9..Maccióetal. 2008): therefore. the halo mass erows at a faster rate than the projected dark," However, increasing the halo mass also increases the scale radius because the virial radius increases and the concentration becomes smaller \citep[e.g.,][]{maccio}; ; therefore, the halo mass grows at a faster rate than the projected dark"
detection of the FISC requires detection of compact τοpau ciuission in a deuse core with no enüssiouat AS21 yan.,detection of the FHSC requires detection of compact $70~\micron$ emission in a dense core with no emissionat $\lambda \lesssim 24~\micron$ .
 The 70jan scusitivity limit of the ¢2d survey was approximately 75 1η (So). somewhat higher than the fux expectedfor a typical FIISC (Boss&Yorke1995:Young&Evans 2005).," The $70 ~\micron$ sensitivity limit of the c2d survey was approximately 75 mJy $5\sigma$ ), somewhat higher than the flux expectedfor a typical FHSC \citep{by95,ye05}."
.. Thus auv FUSCs appeared as starless cores in our ceusus. aud deeper 70421 follow-up observations were necessary fo reveal We selected targets frou our census of starless cores (Enochctal.2008).. which iucludes 108 cores and is complete to Moore20.2ML.. based on three criteria. (," Thus any FHSCs appeared as starless cores in our census, and deeper $70~\micron$ follow-up observations were necessary to reveal We selected targets from our census of starless cores \citep{enoch08}, which includes 108 cores and is complete to $\mcore \gtrsim 0.2 \msun$, based on three criteria. ("
1) Visual inspection of the ¢2d 70.pau maps at starless core positions. looking for faint cunissiou below the formal 5o detection linüt. (,"1) Visual inspection of the c2d $70~\micron$ maps at starless core positions, looking for faint emission below the formal $5\sigma$ detection limit. ("
2) Starless cores with mass-to-size ratios above the critical value for stable Bounor-Ehert spheres (central to surface density ratio ορ21l. or M AL.J/PWOAL >31.9 for T=10 I).,"2) Starless cores with mass-to-size ratios above the critical value for stable Bonnor-Ebert spheres (central to surface density ratio $\rho_c / \rho_0 > 14$, or M $\msun$ ]/FWHM [pc] $> 31.9$ for $T=10$ K)."
 These cores are likely to be [pe]collapsing. but the lack of uecar- or muid-infrared cussion indicates that they have not vet formed a true protostar. (," These cores are likely to be collapsing, but the lack of near- or mid-infrared emission indicates that they have not yet formed a true protostar. ("
"3) Spatially compact starless cores. either muresolved bv the 31"" Bolocam beam or colmpact iu follow-up. CARAIA 2.9 nua interferometric maps (Enoch et ab.","3) Spatially compact starless cores, either unresolved by the $31\arcsec$ Bolocam beam or compact in follow-up CARMA 2.9 mm interferometric maps (Enoch et al.,"
 in prep)., in prep).
 Again. these cores have high central density aud may be collapsing.," Again, these cores have high central density and may be collapsing."
 Iu the three clouds. 7 starless cores met one or more of the above criteria.," In the three clouds, 7 starless cores met one or more of the above criteria."
 This is by uo means a complete or nubiased sample: rather. we used these criteria to identity interesting objects for deeper follow-up observations.," This is by no means a complete or unbiased sample; rather, we used these criteria to identify interesting objects for deeper follow-up observations."
 Deep 70jan maps of our 7 targets were obtained with the Multi-Daud ΠασάςPhotometer for Spitzer (MIPS) ou the (Wernerctal.2001) during 2008 October 20 - October 27., Deep $70~\micron$ maps of our 7 targets were obtained with the Multi-Band ImagingPhotometer for Spitzer (MIPS) on the \citep{werner04} during 2008 October 20 - October 27.
 Iutegration times ranged between P1002000 8 pixel|. depending ou the brightuess of the background/— estimated from c2d maps.," Integration times ranged between $1100-2000$ s $^{-1}$, depending on the brightness of the background estimated from c2d maps."
 The resulting Ὁσ detection limit of ~7.5 uJ was chosen based ou the models of Οι(2007)... to detect FIISC's with mass =(L015AL. at the distance of Perseus aud Serpeus (200260 pe).," The resulting $5\sigma$ detection limit of $\sim 7.5$ mJy was chosen based on the models of \citet{omukai07}, to detect FHSCs with mass $\gtrsim 0.015~\msun$ at the distance of Perseus and Serpens $250-260$ pc)."
 Data were reduced. using the ce2d. pipeline (Evans.ctal. 2007)., Data were reduced using the c2d pipeline \citep{evans07}.
. One source was clearly detected at 70sau: [EY Bolo 55 (RA=03 29 25.7. Dec=|31 28 16.3). hereafter C2006]Per-Bolo 58. a dense core identified as starless by both Enochetal.(2006). aud Watchelletal.(2007).," One source was clearly detected at $70~\micron$: [EYG2006] Bolo 58 (RA=03 29 25.7, Dec=+31 28 16.3), hereafter Per-Bolo 58, a dense core identified as starless by both \citet{enoch06} and \citet{hatch07}."
. Per-Bolo 58 was included i our target Hist based on faint enuission in the c2d 70µια map and a spatially compact 1.1 uni core., Per-Bolo 58 was included in our target list based on faint emission in the c2d $70~\micron$ map and a spatially compact 1.1 mm core.
 It is located in the outskirts of the NGC 1333 vouug cluster (see Figure 2.. right).," It is located in the outskirts of the NGC 1333 young cluster (see Figure \ref{mapfig}, , right)."
 The 70jau flux of Per-Bolo 58 is 65+6 mJ., The $70~\micron$ flux of Per-Bolo 58 is $65\pm6$ mJy.
 A subsequent search of the c2d deep catalogs (Evansotal.2007) revealed a low signal-to-noise 21jan detection at the position of Per-Bolo 5s. with a fux of 0.582:0.2 Lindy.," A subsequent search of the c2d deep catalogs \citep{evans07} revealed a low signal-to-noise $24~\micron$ detection at the position of Per-Bolo 58, with a flux of $0.88\pm0.24$ mJy."
" This source was not included in the “high reliability” or ""η caucidate” e2d catalogs. as it did not mect the τσ signal-to-noiseοσο requirement."," This source was not included in the “high reliability” or “YSO candidate” c2d catalogs, as it did not meet the $7\sigma$ signal-to-noise requirement."
 Caven the good positional correspoudence and point-like emission. we believe the 2|gan source is associated with Per-Bolo 58.," Given the good positional correspondence and point-like emission, we believe the $24~\micron$ source is associated with Per-Bolo 58."
" There is also a nearby low signal-to-noise [.5pan detection (0.02[2E0.007 11Jv) that we do not believe is associated. as the onlv poiut-like οσοι is offset bv 10"" from the 21jan and TOjuu sources."," There is also a nearby low signal-to-noise $4.5~\micron$ detection $0.024\pm0.007$ mJy) that we do not believe is associated, as the only point-like emission is offset by $10\arcsec$ from the $24~\micron$ and $70~\micron$ sources."
 We treat the [5jan point as an upper limit., We treat the $4.5~\micron$ point as an upper limit.
" We estimate a 160ju flux of 2,8 Jv from the c2d MIPS 160jan map: this value has a luwee uncertaity (~50%)) due to nou-uniform extended enission in NCC 1333.", We estimate a $160~\micron$ flux of 2.8 Jy from the c2d MIPS $160~\micron$ map; this value has a large uncertainty $\sim50$ ) due to non-uniform extended emission in NGC 1333.
 naages at 1.5.21.70. aud 160pau are shown in Figure 2..," images at $4.5, 24, 70$ , and $160~\micron$ are shown in Figure \ref{mapfig}."
 All known coutinuuu fiuxes for Per-Bolo 58 are eiven in Table 1.. aud the observed SED is shown iu Figure 1..," All known continuum fluxes for Per-Bolo 58 are given in Table \ref{fluxtab}, and the observed SED is shown in Figure \ref{predictfig}."
 Tn addition to the data described above. we include SITARC. II: 350sau (Dunham et al.," In addition to the data described above, we include SHARC II $350~\micron$ (Dunham et al.,"
 in prep). SCUBA 850jnu (Iatchelletal.2005).. and Bolocam 1.1 nuu (Enochoetal.2006) fluxes. as well as the CARMA 2.9 nua map from Schueectal.(2010).," in prep), SCUBA $850~\micron$ \citep{hatch05}, and Bolocam 1.1 mm \citep{enoch06} fluxes, as well as the CARMA 2.9 mm map from \citet{schnee10}."
.. The interferometric CARALTA 2.9 nuu map provides both a long wavelength flux measuremoeut aud a mieasure of the resolved radial iuteusitv profile., The interferometric CARMA 2.9 mm map provides both a long wavelength flux measurement and a measure of the resolved radial intensity profile.
 The 2.9 nuu visibility auplitudes versus ac-distance are shown iu Figures 3 and [.., The 2.9 mm visibility amplitudes versus $uv$ -distance are shown in Figures \ref{fit1fig} and \ref{fit2fig}.
 Ow τοjun detection of Per-Bolo 58 indicates the preseuce of an internal luminosity source in this dense core. with a flux consistent with that expected for a FUSC.," Our $70~\micron$ detection of Per-Bolo 58 indicates the presence of an internal luminosity source in this dense core, with a flux consistent with that expected for a FHSC."
 The detection at 21jan is inconsistent with a FIISC surrounded by a spherically svuuimetric envelope. however (Figure 1)). leaving open the possibility that Por-Dolo 58 may be a very low ποτν protostar.," The detection at $24~\micron$ is inconsistent with a FHSC surrounded by a spherically symmetric envelope, however (Figure \ref{predictfig}) ), leaving open the possibility that Per-Bolo 58 may be a very low luminosity protostar."
 We discuss these possibilities below., We discuss these possibilities below.
 We use RADAIC. a ftio-dineusional Moute Carlo radiative transfer code (Dullemoud&Domini2001) to model a FIISC aud low hunuinositv protostar.," We use RADMC, a two-dimensional Monte Carlo radiative transfer code \citep{dd04} to model a FHSC and low luminosity protostar."
" We adopt a deusitv profile that inclides a rotating iufalliug euvelope characterized by total massAL... ceutrifueal radiusA54,4;. and outer radius citep|e.g.||| ubich67.. aud outflow cavity characterized by full opening angle o."," We adopt a density profile that includes a rotating infalling envelope characterized by total mass, centrifugal radius, and outer radius \\citep[e.g.][]{ulrich67}, and outflow cavity characterized by full opening angle $\ang$."
" The euvelope mass is held fixed at AM,=1.2AL... as caleulated from the 1.1 nuu flux (see Enochetal. 2006)) assuming a dust teniperature of T—8 EK. consistent with radiative trausfer models of such a low hunuiositv source."," The envelope mass is held fixed at $\menv=1.2~ \msun$, as calculated from the 1.1 mm flux (see \citealt{enoch06}) ) assuming a dust temperature of $T=8$ K, consistent with radiative transfer models of such a low luminosity source."
 Details of the deusitv profile aud modeling are as in Enochetal.(20095)., Details of the density profile and modeling are as in \citet{enoch09b}.
 Dust opacities are from Table 1. column 5 of OsseukopfWenning(1991). for dust erains with thin ice mantles. including scattering.," Dust opacities are from Table 1, column 5 of \citet{oh94} for dust grains with thin ice mantles, including scattering."
 Following Dunhametal.(2010).. we remove the effects of excess back-scatteriug of the interstellar radiation ficld (SRE) by subtracting the fiux for a starless core model at short waveleneths.," Following \citet{dunham10}, we remove the effects of excess back-scattering of the interstellar radiation field (ISRF) by subtracting the flux for a starless core model at short wavelengths."
 Dunhametal.(2010) make this correction for A<10 jan: herewe additionally include the simall plateau of ISRF cuissiou at A=10.LO gan., \citet{dunham10} make this correction for $\lambda\le10~\micron$ ; herewe additionally include the small plateau of ISRF emission at $\lambda=10-40 ~\micron$ .
 While this correction affects the shape of our iiodol SEDs atshort wavelengths. the correction at 21jnu ds small. so it does not significantly affect the model fits.," While this correction affects the shape of our model SEDs atshort wavelengths, the correction at $24~\micron$ is small, so it does not significantly affect the model fits."
 To match the peak of the SED at SOyan we include an ISRF that is a factor of 3 strongerthan the “Black-Draine” ISRF (Evansetal.2001:Blacks 1975)..," To match the peak of the SED at $80~\micron < \lambda < 1000~\micron$ we include an ISRF that is a factor of 3 strongerthan the “Black-Draine” ISRF \citep{evans01,black94,draine78}. ."
 A more inteiuse radiation field may be expected in NGC 1333. where there is siguificaut," A more intense radiation field may be expected in NGC 1333, where there is significant"
from the center of mass is. using equation (6). If Mie=αλ«<Mg. Chen litle mass is accreted onto the star Adj.,"from the center of mass is, using equation (6), If $M_{b2} =q M_{b1} \ll M_{b1}$, then little mass is accreted onto the star $M_{b2}$."
 In addition. as will be seen later in Chis section. (he specific angular momentum of the mass accreted onto (he massive component Mj will be low. reducing the probability of accretion disk formation (ihe limit of «1 gives the single accreting star case).," In addition, as will be seen later in this section, the specific angular momentum of the mass accreted onto the massive component $M_{b1}$ will be low, reducing the probability of accretion disk formation (the limit of $q \ll 1$ gives the single accreting star case)."
 The exact conditions for the formation of accretion disks around one or two stars must wail for [ull 2-dimensional numerical simulations., The exact conditions for the formation of accretion disks around one or two stars must wait for full 3-dimensional numerical simulations.
 In (he present section I will estimate (his value., In the present section I will estimate this value.
 I will assume namely. the radius of the accreting star is much smaller than the width of the accretion column at the binary location wer=(425). which is smaller than orbital separation of the accreting binary svstem. which is much smaller than the accretion radius of the binary svstem (eq.," I will assume namely, the radius of the accreting star is much smaller than the width of the accretion column at the binary location $w(x=a_{12})$, which is smaller than orbital separation of the accreting binary system, which is much smaller than the accretion radius of the binary system (eq."
 1). which is smaller (han the orbital separation with (the mass losing star.," 1), which is smaller than the orbital separation with the mass losing star."
 For example. 30AUSaj300AU. 10AUXRaeS30AU. 10.20.05AUXaySLAU. with wir=d45) somewhat smaller. and (OLR.H4ELR..," For example, $30 \AU
\lesssim a_0 \lesssim 300 \AU$, $10 \AU \lesssim R_{\rm acc}
\lesssim 30 AU$, $10 R_\odot \simeq 0.05 \AU \lesssim a_{12}
\lesssim 1 AU$, with $w(x=a_{12})$ somewhat smaller, and $0.01
R_\odot \lesssim R_1 \lesssim 1 R_\odot$."
 As mentioned in section 1. the flow onto the binary svstem will not be exactly axisvinmetrical as depicted in figure 2. but rather the accreted mass will posses some angular momentum.," As mentioned in section 1, the flow onto the binary system will not be exactly axisymmetrical as depicted in figure 2, but rather the accreted mass will posses some angular momentum."
 This excess angular momentum causes the bending of the accretion column and shock wave., This excess angular momentum causes the bending of the accretion column and shock wave.
 However. for the large orbital separation cj assumed here. (his angular momentum and departure [rom axisvinnmetry of the flow at large distance from the accreting binary svstem can be neglected relative to the specific angular momentum of the accreted mass within the binary svstem. as I now show.," However, for the large orbital separation $a_0$ assumed here, this angular momentum and departure from axisymmetry of the flow at large distance from the accreting binary system can be neglected relative to the specific angular momentum of the accreted mass within the binary system, as I now show."
 ] start wilh examining the situation lor perpendicular orbital planes., I start with examining the situation for perpendicular orbital planes.
 In the ealeulations to follow I assume that most accreted mass reaches the binary svstem via (he accretion column. such that the characteristic width of the accretion column wr). fulfills the condition «e(r=ayo)Sdy.," In the calculations to follow I assume that most accreted mass reaches the binary system via the accretion column, such that the characteristic width of the accretion column $w(x)$, fulfills the condition $w(x=a_{12}) \lesssim a_{12}$."
 When wr=ays)Zay the caleulations below do not hold anv more., When $w(x=a_{12}) \gtrsim a_{12}$ the calculations below do not hold any more.
 The binary svstem will be inside the accretion column. the accretion process will be more complicated. and I expect that the Formation of an accretion disk will be less [avorable.," The binary system will be inside the accretion column, the accretion process will be more complicated, and I expect that the formation of an accretion disk will be less favorable."
 In this case @=90° (see fig., In this case $\theta=90^\circ$ (see fig.
 1)., 1).
 Most of the accreted mass flows toward the center of, Most of the accreted mass flows toward the center of
Now we shall COMPALC the spatial distribution of superclusters based ou the Abell and APM cluster samples.,Now we shall compare the spatial distribution of superclusters based on the Abell and APM cluster samples.
 We restrict our comparison to the area of the sky covered by the APM sample., We restrict our comparison to the area of the sky covered by the APM sample.
 We use supercluster catalogs coustructed for all clusters (ucline clusters with estimated redshifts)., We use supercluster catalogs constructed for all clusters (including clusters with estimated redshifts).
 We show the distribution of clusters oulv in superclusters of richuess Land ligher., We show the distribution of clusters only in superclusters of richness 4 and higher.
 This has two reasons., This has two reasons.
 First. the skeleton of the supercluster-void network is essentiallv deerimimned by rcli superclusters: poorer svstenis only add sole fiue details to he structure (E97d).," First, the skeleton of the supercluster-void network is essentially determined by rich superclusters; poorer systems only add some fine details to the structure (E97d)."
 Redshift errors do uot affect the sky positions of the clusters mt have an effect ou their radial «istance onlv., Redshift errors do not affect the sky positions of the clusters but have an effect on their radial distance only.
 Redslüft errors have the sale effect for superclusters as the velocity «ispersion of ealaxies in clusters: they elougate the superclusters iu the racial «irection., Redshift errors have the same effect for superclusters as the velocity dispersion of galaxies in clusters: they elongate the superclusters in the radial direction.
 This effect may destroy real juperclusters. if the separation between individual clusters im the radial direction. due to errors. is too large.," This effect may destroy real superclusters, if the separation between individual clusters in the radial direction, due to errors, is too large."
 Usine rich aud very rich superclusters we have a Neher probability to select real superclusters since it is very unlikely that errors bring isolatec clusters together to form compact rich svstenis., Using rich and very rich superclusters we have a higher probability to select real superclusters since it is very unlikely that errors bring isolated clusters together to form compact rich systems.
 Figure 23 shows that in fus region of the sky the distribution of clusters das two asia. located at cstaucc Sorc200 and 300 ffrom 1s," Figure 3 shows that in this region of the sky the distribution of clusters has two maxima, located at distances $r \approx 200$ and 300 from us."
 The APM supereaster catalog contains no rich suverchisters at distances less than 150Mpc. SO Wwe ot iu Fieure 6 clusers du rich and very rich superchsters in two lavers. 150rox250 Ape. aud 25οπ350\pe.. which correspoud to imaxina of the cluster sample density.," The APM supercluster catalog contains no rich superclusters at distances less than 150, so we plot in Figure 6 clusters in rich and very rich superclusters in two layers, $150 \le r
< 250$ , and $250 \le r < 350$, which correspond to maxima of the cluster sample density."
" oThe richest superchsters in the nearer sheet (lower paucls in Figure 6) are the Pisces-Cetus. Caus-Indus. Horologiuu-Roeticului (see also Park Lee 1998) superclusters located at RA xOlrad.sin(DEC)~QE: RA0.6.sin(DEC)zm 0.5. and RA0.9.sin(DEC)5-0.5. respectively,"," The richest superclusters in the nearer sheet (lower panels in Figure 6) are the Pisces-Cetus, Grus-Indus, Horologium-Reticulum (see also Park Lee 1998) superclusters located at RA $\approx 0.1\,rad, ~\sin(DEC) \approx -0.4$; $\approx -0.6,
~\sin(DEC) \approx -0.8$ , and $\approx 0.9, ~\sin(DEC) \approx
-0.8$, respectively."
" The richest superclusters in the farther sheet (upper pancls) are the Sculptor supercluster at RAs0.1.sindc0.5, acl the Cachun and Forunax-Exidauus supercluster complex at RÀ Id.sind=0.6."," The richest superclusters in the farther sheet (upper panels) are the Sculptor supercluster at $ \approx 0.1, ~\sin
\delta \approx -0.5$, and the Caelum and Fornax-Eridanus supercluster complex at RA $ \approx 1.1, ~\sin \delta \approx -0.6$."
 Very rich superclusters are seelin both (the Abell aud APM) supercluster catalogs base ou all clusters. although individual clusters differ.," Very rich superclusters are seen in both (the Abell and APM) supercluster catalogs based on all clusters, although individual clusters differ."
 Extra-laree open circles in Figure 6 mark the ceuters of very rich superclusters as determined from positious of t1011) member clusters in the APALRS supercluster catalog., Extra-large open circles in Figure 6 mark the centers of very rich superclusters as determined from positions of their member clusters in the APM.R8 supercluster catalog.
 Tn the nearer sheet only the Worologimm-Reticuhim supercluster complex has a sufficient number of 1nemboers with meastred redshifts. and is split iuto two subsystems. separatec by 25 Hn D space.," In the nearer sheet only the Horologium-Reticulum supercluster complex has a sufficient number of members with measured redshifts, and is split into two subsystems, separated by 25 in 3-D space."
 In the farther sheet we see the Sculptor supercluster which is also split into two subsvstcdus (separated by 50 Hn 3-D space. see Figure 9).," In the farther sheet we see the Sculptor supercluster which is also split into two subsystems (separated by 50 in 3-D space, see Figure 9)."
 We conclude that the whole ADPM supoercluster catalog. derived ou the basis of nneasured redshitts. is donuuated by two very rich supoercluster complexes. the Uorologituu-Reticuhun and the Seuotor superclusters.," We conclude that the whole APM supercluster catalog, derived on the basis of measured redshifts, is dominated by two very rich supercluster complexes, the Horologium-Reticulum and the Sculptor superclusters."
 Other rich superclusters in the APM sample have oulv a few. if any. clusters with measured redshifts.," Other rich superclusters in the APM sample have only a few, if any, clusters with measured redshifts."
 These dominant rich superclusters surround the Sculptor void. oue of the largest voids known with a ciameter| of about 185Mpe.," These dominant rich superclusters surround the Sculptor void, one of the largest voids known with a diameter of about 185."
. Ou the borders of this void Were are several APM clusters forming a cluster filament., On the borders of this void there are several APM clusters forming a cluster filament.
 The spatial density of this flameut is. however. lower than the lumii density below which for claritv we omit clusters in low-deusity enviromnents in Figure 6G.," The spatial density of this filament is, however, lower than the minimum density below which for clarity we omit clusters in low-density environments in Figure 6."
 The sample of all APA clusters (APALAL) reaches its iixiual spatial deusity at a distance r2:300 ((see Figure 3)., The sample of all APM clusters (APM.A1) reaches its maximal spatial density at a distance $r \approx 300$ (see Figure 3).
 Itis interesting to note that clusters which, Itis interesting to note that clusters which
Tard X-ray observations are the most efficient wav of tracing cuuission due to accretion mcchanisias. such in Active Galactic Nuclei (AGN).,"Hard X-ray observations are the most efficient way of tracing emission due to accretion mechanisms, such in Active Galactic Nuclei (AGN)."
 Tard N-ray selection is less affected by strong biases present at other wavelenetls., Hard X-ray selection is less affected by strong biases present at other wavelengths.
 For example. a colin of a few times 1072 ? has ucelieible effect in the 5-10 keV baud. while it reduces by ~100 times unclear emission below 2 keV. Soft X-ray surveys (0.8. Hasimger et al.," For example, a column of a few times $10^{22}$ $^{-2}$ has negligible effect in the 5-10 keV band, while it reduces by $\sim100$ times nuclear emission below 2 keV. Soft X-ray surveys (e.g. Hasinger et al."
 1998. Schinidt et al.," 1998, Schmidt et al."
 1998) are also often coutaminated by nou uuclear components. like cussion from binaries and/or from optically thin plasmas iu stir-forimation regions surroundius the uncleus.," 1998) are also often contaminated by non nuclear components, like emission from binaries and/or from optically thin plasmas in star-formation regions surrounding the nucleus."
 Optical and UW color selection is biased against objects with even modest extinction or an intrinsically που cussion spectrum, Optical and UV color selection is biased against objects with even modest extinction or an intrinsically `red' emission spectrum
larger to the large-scale structures.,larger to the large-scale structures.
 In this description we can suspect. that the small scale structures are almost isotropic.," In this description we can suspect, that the small scale structures are almost isotropic."
 The degree of anisotropy grows with the scale. and for the large scale structures it exceeds somewhat the GS95 anisotropy.," The degree of anisotropy grows with the scale, and for the large scale structures it exceeds somewhat the GS95 anisotropy."
 In the case of a weak external magnetic field the degree of anisotropy is uniform over the whole range of scales. but also here we notice very good agreement with GS95 anisotropy.," In the case of a weak external magnetic field the degree of anisotropy is uniform over the whole range of scales, but also here we notice very good agreement with GS95 anisotropy."
 The compressible component of the velocity field behaves differently., The compressible component of the velocity field behaves differently.
 For subAfvénnic turbulence it is almost perfectly isotropic and independent of the sonic Mach number., For subAfvénnic turbulence it is almost perfectly isotropic and independent of the sonic Mach number.
 In the superAlfvénnic turbulence. however. we see. that structure of subsonic potential field is isotropic on average. although there are relatively large time departures from the isotropy.," In the superAlfvénnic turbulence, however, we see, that structure of subsonic potential field is isotropic on average, although there are relatively large time departures from the isotropy."
 On contrary. in the case of supersonic turbulence. the structure of the potential field contains a higher amount of anisotropic structures and the degree of anisotropy aims to the GS95-]aw.," On contrary, in the case of supersonic turbulence, the structure of the potential field contains a higher amount of anisotropic structures and the degree of anisotropy aims to the GS95-law."
 In Figure 5. we show the anisotropy of structures for Alfvénn. slow and fast modes of velocity.," In Figure \ref{fig:mode_anisotropy} we show the anisotropy of structures for Alfvénn, slow and fast modes of velocity."
 The separation of the modes is performed with the respect of the local mean magnetic field., The separation of the modes is performed with the respect of the local mean magnetic field.
 The first comparison of these plots shows. that the degree of anisotropy does not depend on the sonic Mach number for the Alfvénn and slow modes.," The first comparison of these plots shows, that the degree of anisotropy does not depend on the sonic Mach number for the Alfvénn and slow modes."
 The anisotropy of fast mode. similarly to the case of the compressible part of the velocity. depends on the regime of sonic motions.," The anisotropy of fast mode, similarly to the case of the compressible part of the velocity, depends on the regime of sonic motions."
 If the fluid motions are subsonie the structures of the fast waves is more isotropic., If the fluid motions are subsonic the structures of the fast waves is more isotropic.
 When the fluid motions are supersonic. more structures become anisotropic.," When the fluid motions are supersonic, more structures become anisotropic."
 We should note here. that speaking about the anisotropy of structures we understand a mean anisotropy of all structures of the analyzed field.," We should note here, that speaking about the anisotropy of structures we understand a mean anisotropy of all structures of the analyzed field."
 Individual structures in the turbulence evolve and yield at different deformations resulting in a constant change of the degree of the anisotropy., Individual structures in the turbulence evolve and yield at different deformations resulting in a constant change of the degree of the anisotropy.
 In this situation. speaking about an anisotropy. we are relating to the mean dominant anisotropy of all structures in the system.," In this situation, speaking about an anisotropy, we are relating to the mean dominant anisotropy of all structures in the system."
 The two remaining modes. the Alfvénn and slow. show very similar anisotropies to those observed in velocity and its incompressible part.," The two remaining modes, the Alfvénn and slow, show very similar anisotropies to those observed in velocity and its incompressible part."
 We see that basically the degree of anisotropy of structures in these two modes ts closer to GS95., We see that basically the degree of anisotropy of structures in these two modes is closer to GS95.
 In addition. the presence of a strong magnetic field results in a much stronger bending of the curves in right plots of Figure 5 signifying changes of the anisotropy with the scale of the structure.," In addition, the presence of a strong magnetic field results in a much stronger bending of the curves in right plots of Figure \ref{fig:mode_anisotropy} signifying changes of the anisotropy with the scale of the structure."
" The curves in plots for superAlfvénnic models are almost straight and independent of the values of sl,", The curves in plots for superAlfvénnic models are almost straight and independent of the values of $S^{\parallel}_2$.
 Intermittency is an essential property of astrophysical fluids., Intermittency is an essential property of astrophysical fluids.
 As intermittency violates self-similarity of motions. it is impossible to naively extrapolate the properties of fluids obtained computationally with a relatively low resolution to the actual astrophysical situations.," As intermittency violates self-similarity of motions, it is impossible to naively extrapolate the properties of fluids obtained computationally with a relatively low resolution to the actual astrophysical situations."
 In astrophysics the intermittency affects turbulent heating. momentum transfer. interaction with cosmic rays. radio waves and many more essential processes.," In astrophysics the intermittency affects turbulent heating, momentum transfer, interaction with cosmic rays, radio waves and many more essential processes."
 Physical interpretation of intermittency started after the work by Kolmogorov. but the first successful model was presented by(1994).," Physical interpretation of intermittency started after the work by Kolmogorov, but the first successful model was presented by."
". The scaling relations suggested by relate ς(ϱ) to the scaling of the velocity v~/'/*. the energy cascade rate tj!~ΓΑ, and the co-dimension of the dissipative structures C: Parameter C is related to the dimension of the dissipative structures D through relation C23—D 2000)."," The scaling relations suggested by relate $\zeta(p)$ to the scaling of the velocity $v_l
\sim l^{1/g}$, the energy cascade rate $t^{-1}_l \sim l^{-x}$, and the co-dimension of the dissipative structures $C$: Parameter $C$ is related to the dimension of the dissipative structures $D$ through relation $C = 3 - D$ ."
 In hydrodynamical turbulence. according to the Kolmogorov scaling. we have g23 and x22/3.," In hydrodynamical turbulence, according to the Kolmogorov scaling, we have $g = 3$ and $x = 2/3$."
 Vortex filaments. which are one-dimensional structures. correspond to C22 (D=1).," Vortex filaments, which are one-dimensional structures, correspond to $C=2$ $D=1$ )."
 In the MHD turbulence we also observe current sheets. which are two-dimensional dissipative structures and correspond to C=| (D=2).," In the MHD turbulence we also observe current sheets, which are two-dimensional dissipative structures and correspond to $C=1$ $D=2$ )."
 For these two types of dissipative structures we obtain two different scaling relations (substituting g23 and x22/3): and Relation (17)) is often called the She Lévvéqque scaling1994).. while relation (18)). the Miilller-Biskamp scaling2000).," For these two types of dissipative structures we obtain two different scaling relations (substituting $g = 3$ and $x = 2/3$ ): and Relation \ref{eqn:sl}) ) is often called the She Lévvêqque scaling, while relation \ref{eqn:mb}) ), the Mülller-Biskamp scaling."
. There are theoretical areuments against the model1994).. but so far the She Lévvéqque scaling is the best for reproducing the intermittency of incompressible hydrodynamic turbulence.," There are theoretical arguments against the model, but so far the She Lévvêqque scaling is the best for reproducing the intermittency of incompressible hydrodynamic turbulence."
 Structure functions can be calculated with respect to the global or local reference frames., Structure functions can be calculated with respect to the global or local reference frames.
 By the scaling exponents calculated in the global reference frame we understand the scaling exponents calculated from the structure functions averaged over all directions., By the scaling exponents calculated in the global reference frame we understand the scaling exponents calculated from the structure functions averaged over all directions.
 In the local reference we distinguish between directions parallel and perpendicular to the local mean magnetic field., In the local reference we distinguish between directions parallel and perpendicular to the local mean magnetic field.
 In this way we define the reference frame locally by the direction of the local magnetic field., In this way we define the reference frame locally by the direction of the local magnetic field.
 In Figures 6 and 7 we show scaling exponents for the velocity and all its parts and waves calculated in the global reference frame., In Figures \ref{fig:expons_parts} and \ref{fig:expons_waves} we show scaling exponents for the velocity and all its parts and waves calculated in the global reference frame.
 In the top left plot of Figure 6 we see that for the subAlfvénnic turbulence the scaling exponents of velocity follow the She-Lévvéqque (S-L) scaling with D=I., In the top left plot of Figure \ref{fig:expons_parts} we see that for the subAlfvénnic turbulence the scaling exponents of velocity follow the She-Lévvêqque (S-L) scaling with $D=1$.
 Supported by the theoretical considerations we can say that most of the dissipative structures are one-dimensional., Supported by the theoretical considerations we can say that most of the dissipative structures are one-dimensional.
 Even though the scalings are not perfectly independent of the value of M. since we see somewhat lower values of à for higher p. the differences between these values for models with different sonic Mach numbers are within their error bars. thus it ts relatively difficult to state that the sealings are completely independent or only weakly dependent of the values of τς.," Even though the scalings are not perfectly independent of the value of ${\cal M}_s$, since we see somewhat lower values of $\zeta$ for higher $p$, the differences between these values for models with different sonic Mach numbers are within their error bars, thus it is relatively difficult to state that the scalings are completely independent or only weakly dependent of the values of ${\cal M}_s$."
 Looking in the corresponding plot for models with à weak magnetic field we clearly see that the spread of curves for different sonic Mach numbers ts much higher than in the previous case., Looking in the corresponding plot for models with a weak magnetic field we clearly see that the spread of curves for different sonic Mach numbers is much higher than in the previous case.
 For subsonic model the scaling exponents of velocity follow very well the theoretical curve defined by the S-L scaling with parameter D corresponding to one-dimensional structures., For subsonic model the scaling exponents of velocity follow very well the theoretical curve defined by the S-L scaling with parameter $D$ corresponding to one-dimensional structures.
" The model with Vi,~2.3. however. follows perfectly the S-L scaling with D=2 corresponding to the two-dimensional dissipative structures."," The model with ${\cal M}_s \sim 2.3$, however, follows perfectly the S-L scaling with $D=2$ corresponding to the two-dimensional dissipative structures."
 Moreover. models with even higher values of the sonic Mach number have the scaling exponents for p=>3 somewhat below the S-L scaling with D=2.," Moreover, models with even higher values of the sonic Mach number have the scaling exponents for $p>3$ somewhat below the S-L scaling with $D=2$."
 These observations suggest that the scaling exponents of the velocity change with the sonic Mach number but only in the case of weak magnetic field turbulence., These observations suggest that the scaling exponents of the velocity change with the sonic Mach number but only in the case of weak magnetic field turbulence.
 The presence of a strong magnetic field significantly reduces these changes and preserves the generation of the dissipative structures of higher thanone dimensions., The presence of a strong magnetic field significantly reduces these changes and preserves the generation of the dissipative structures of higher thanone dimensions.
 After the decomposition of velocity into its incompressible and compressible parts we also calculate their scaling exponents., After the decomposition of velocity into its incompressible and compressible parts we also calculate their scaling exponents.
 In the middle and right columns of Figure 6 , In the middle and right columns of Figure \ref{fig:expons_parts} 
lleinz Begclman (2002). this choice of parameters results in a cocoon in which the Ixelvin-Helmholtz. growth: rates approximately match those caleulatecl for real systems.,"Heinz Begelman (2002), this choice of parameters results in a cocoon in which the Kelvin-Helmholtz growth rates approximately match those calculated for real systems."
 Code units are related. to physical units by parameters fixed to the background. medium., Code units are related to physical units by parameters fixed to the background medium.
" For most runs we use the following canonical conversion factors: ry=100kpc. ear=1000kis no=O.01cm and 4j=44.10° Ix. This gives a total kinetic luminosity of the jets of 9.3eres making the simulation relevant. for. powerful radio sources located in rich galaxy clusters. like Cvgnus A. We will refer to this as our ""large cluster case”."," For most runs we use the following canonical conversion factors: $r_{0}=100\,{\rm kpc}$ , $c_{\rm
ISM}=1000\,{\rm km}\,{\rm s}^{-1}$ , $n_{0}=0.01\,{\rm cm}^{-3}$ and $T_{0}=4.4\times 10^{7}$ K. This gives a total kinetic luminosity of the jets of $9.3\times 10^{45}\,{\rm erg}\,{\rm s}^{-1}$ , making the simulation relevant for powerful radio sources located in rich galaxy clusters, like Cygnus A. We will refer to this as our “large cluster case”."
" Since many radio galaxies are. however. found in smaller and cooler galaxy clusters. we also consider a cdillerent. scaling (ο=dOkpe. eis=500kms|. ny=O.05em* and Ty=LlO° Ix) which results in a total kinetic source luminosity of 1.16l0ergs| and is closer to objects like Hvedra A and. Vireo A. We will refer to this as our ""spiall cluster case”."," Since many radio galaxies are, however, found in smaller and cooler galaxy clusters, we also consider a different scaling $r_{0}=10\,{\rm kpc}$, $c_{ISM}=500\,{\rm km}\,{\rm s}^{-1}$, $n_{0}=0.05\,{\rm cm}^{-3}$ and $T_{0}=1.1\times 10^{7}$ K) which results in a total kinetic source luminosity of $1.16\times
10^{44}\,{\rm erg}\,{\rm s}^{-1}$ and is closer to objects like Hydra A and Virgo A. We will refer to this as our “small cluster case”."
 In the first case. one code unit of time corresponds to 50 Myr.," In the first case, one code unit of time corresponds to 50 Myr."
 The total length of the simulation is 240 Myr. the jet is turned olf after 50 Myr.," The total length of the simulation is 240 Myr, the jet is turned off after 50 Myr."
 In the second case. one code unit of time is LO Myr. the total length is about 50. Myr. and the jet is turned off after LO Myr.," In the second case, one code unit of time is 10 Myr, the total length is about 50 Myr, and the jet is turned off after 10 Myr."
 For more cletailed information on the simulations. see Itevnolcds. lleinz 3ceclman (2002).," For more detailed information on the simulations, see Reynolds, Heinz Begelman (2002)."
 As an example of the simulation results. a density map is given in Fig.," As an example of the simulation results, a density map is given in Fig."
 1., 1.
 Lt is important to note that since this simulations were performed. Chandra observations have revealed several cases where the radio lobes of the central source are surrounded bv a significant amount of cold gas.," It is important to note that since this simulations were performed, Chandra observations have revealed several cases where the radio lobes of the central source are surrounded by a significant amount of cold gas."
 The ICM structures surrounding Perscus-A provide an excellent. example. (sec Fabian et al., The ICM structures surrounding Perseus-A provide an excellent example (see Fabian et al.
 2003a: Schmidt. Fabian Sanders 2002).," 2003a; Schmidt, Fabian Sanders 2002)."
 At the current time. the physical mechanisms underlving these cool rims are not understood aad they cannot be reproduced in any first principles simulations.," At the current time, the physical mechanisms underlying these cool rims are not understood and they cannot be reproduced in any first principles simulations."
 ln systems. containing such rims. the observed oxvgen OVILE and OVILE absorption lines towards the central AGN are likely to be dominated by the contribution from gas in the rim: according to Schmidt. Fabian Sanders (2002). 106 temperature and density in the cold gas rim of the Perseus cluster are 3.6«LO Ix and 0.12em.7 respectively.," In systems containing such rims, the observed oxygen OVII and OVIII absorption lines towards the central AGN are likely to be dominated by the contribution from gas in the rim; according to Schmidt, Fabian Sanders (2002), the temperature and density in the cold gas rim of the Perseus cluster are $3.6\times10^{7}$ K and $0.12\,{\rm cm}^{-3}$ respectively."
" Using simple caleulations based on Shull van Steenberg (1982). this gives an OVILE equivalent width of ουδ), larger than that predicted from the entire rest ofthe ICM atmosphere."," Using simple calculations based on Shull van Steenberg (1982), this gives an OVIII equivalent width of eV, larger than that predicted from the entire rest of the ICM atmosphere."
 Pherefore. our results presented here are valid only for systems that do not possess these cold rims. or for lines of sight that do not intercept such rims.," Therefore, our results presented here are valid only for systems that do not possess these cold rims, or for lines of sight that do not intercept such rims."
 The hwdrodysnamic simulation gives us the density. temperature ancl velocity of the ICAL at. every. point. in space as a function of time throughout the modelled radio-galaxyLOCAL interaction.," The hydrodynamic simulation gives us the density, temperature and velocity of the ICM at every point in space as a function of time throughout the modelled radio-galaxy/ICM interaction."
 To simulate the absorption line spectrum that one would see towards the central racio galaxy. we used the photoionization code οΕΛ (developed bv Tim Ixallman: see Kallman Bautista 2001) to calculate the opacity of the IC'M along radial rays that originate at the AGN.," To simulate the absorption line spectrum that one would see towards the central radio galaxy, we used the photoionization code XSTAR (developed by Tim Kallman; see Kallman Bautista 2001) to calculate the opacity of the ICM along radial rays that originate at the AGN."
 We assume a collisionally-ionized plasma in thermal and ionization equilibrium., We assume a collisionally-ionized plasma in thermal and ionization equilibrium.
 These assumptions are reacilv verified for the ions of interest., These assumptions are readily verified for the ions of interest.
 We take the temperature. clensity ancl velocity directly from the hydrodynamic simulation.," We take the temperature, density and velocity directly from the hydrodynamic simulation."
 The frequencey-dependent opacity is calculated including both thermal line broadening ancl line-oFsight Doppler shifts., The frequency-dependent opacity is calculated including both thermal line broadening and line-of-sight Doppler shifts.
 We choose an energy resolution of Adefie=3.291891. givingNEM a velocity. resolution of about 100 km/s. ΝΤΑ includes 19 elements. and. the elemental abundances are important for the resulting line spectrum.," We choose an energy resolution of $\Delta E/E =
3.29\times 10^{-4}$, giving a velocity resolution of about 100 km/s. XSTAR includes 13 elements, and the elemental abundances are important for the resulting line spectrum."
 We chose abuncances measured by Allen Fabian (1998). Loewenstein Mushotzky (1996) and Baumgartner ct al. (," We chose abundances measured by Allen Fabian (1998), Loewenstein Mushotzky (1996) and Baumgartner et al. ("
2005). which we summarize in Table 1.,"2005), which we summarize in Table 1."
" ALL publicly available photoionization codes have some problems dealing. with very low densities anc very high temperatures (210"" A: NSTAR is no exception."," All publicly available photoionization codes have some problems dealing with very low densities and very high temperatures $>10^{9}
K$ ); XSTAR is no exception."
 However. such high temperatures only occur in carly times of the simulations. when the jet is still on. and then only in the innermost eric cells. very close to the source.," However, such high temperatures only occur in early times of the simulations, when the jet is still on, and then only in the innermost grid cells, very close to the source."
 We can safely ignore those cells (through. the inclusion of a. threshold temperature above which the simulation cell is skipped) since very high temperature gas has no line absorption anyway., We can safely ignore those cells (through the inclusion of a threshold temperature above which the simulation cell is skipped) since very high temperature gas has no line absorption anyway.
 Phe only case where a considerable number of grid cells is alfected is along the jet-axis @<1., The only case where a considerable number of grid cells is affected is along the jet-axis $\theta<1^\circ$.
" Due to the forced. svmmetrv of the simulations the region 8<1 is unphysical because material gets ""stacked"" on the jet axis.", Due to the forced symmetry of the simulations the region $\theta<1^\circ$ is unphysical because material gets “stacked” on the jet axis.
 The minimum line-ol-si¢ht angle that we choose to analyze is @=5.5°. appropriate for blazar-tvpe sources.," The minimum line-of-sight angle that we choose to analyze is $\theta=5.5^\circ$, appropriate for blazar-type sources."
 In this section. we discuss the impactthat theradio-galaxy activity. has on the observed. X-ray absorption lines in our simulated system.," In this section, we discuss the impactthat theradio-galaxy activity has on the observed X-ray absorption lines in our simulated system."
 The astrophysical relevance of these, The astrophysical relevance of these
Donnéees astronomique,Donnéees astronomique.
(CDS)'.. Table 1. summarises the available photometry., Table \ref{tbl_phot} summarises the available photometry.
 The V band photometry is estimated from our flux calibrated spectroscopy., The $V$ band photometry is estimated from our flux calibrated spectroscopy.
 We estimated the uncertainty in the V band to be 0.3 mag. which includes systematic effects of varying seeing and data reduction procedures.," We estimated the uncertainty in the $V$ band to be 0.3 mag, which includes systematic effects of varying seeing and data reduction procedures."
 The2MASS Ky band is uncertain and is not considered further., The $K_S$ band is uncertain and is not considered further.
 We first fitted the observed Balmer lines with a grid of pure hydrogen synthetic spectra that extends from 5000 to 7000 K (n steps of 500 K). 8000 to 000 Κ (in steps of 1000 K). 000 to 000 K (in steps of 2000 K). and 000 to 000 K (in steps of 4000 K) at logg=6.0 to 9.5 (in steps of 0.25 dex).," We first fitted the observed Balmer lines with a grid of pure hydrogen synthetic spectra that extends from 5000 to 7000 K (in steps of 500 K), 8000 to 000 K (in steps of 1000 K), 000 to 000 K (in steps of 2000 K), and 000 to 000 K (in steps of 4000 K) at $\log{g} = 6.0$ to $9.5$ (in steps of 0.25 dex)."
 Following initial parameter estimations for NLTT 11748. we calculated a denser grid with temperatures ranging from 8000 K to 9000 K (in steps of 100 K) and surface gravities ranging from logg=6.0 to 7.0 (in steps of 0.25 dex).," Following initial parameter estimations for NLTT 11748, we calculated a denser grid with temperatures ranging from 8000 K to 9000 K (in steps of 100 K) and surface gravities ranging from $\log{g} = 6.0$ to $7.0$ (in steps of 0.25 dex)."
 Figure | shows the observed Balmer line profiles compared to the best-fit model spectrum., Figure \ref{fig_fit} shows the observed Balmer line profiles compared to the best-fit model spectrum.
 This model spectrum corresponds to an effective temperature of Tj=8540+50 K and a surface gravity of logg=6.20x0.15., This model spectrum corresponds to an effective temperature of $T_{\rm eff} = 8540\pm50$ K and a surface gravity of $\log{g} = 6.20\pm0.15$.
 The quoted uncertainties are only statistical (1c7)., The quoted uncertainties are only statistical $1\sigma$ ).
 In general surface gravity measurements may be effected by a systematic offset between the assumed and the actual spectral resolution., In general surface gravity measurements may be effected by a systematic offset between the assumed and the actual spectral resolution.
 At low surface gravity. a ooffset would correspond to =0.08dex shift in surface gravity or 0.007M...," At low surface gravity, a offset would correspond to $\approx 0.08$dex shift in surface gravity or $0.007\ M_\odot$."
 Using the evolutionary mass-radius relations of Serenellietal. (2002).. we obtain a mass of 0.167+0.005M.. and a cooling age of 4.0—6.3 Gyrs.," Using the evolutionary mass-radius relations of \citet{ser2002}, we obtain a mass of $0.167\pm0.005\ M_\odot$ and a cooling age of $4.0 - 6.3$ Gyrs."
 The use of low-metallicity models from Serenellietal.(2002) is justified by the old age of the star., The use of low-metallicity models from \citet{ser2002} is justified by the old age of the star.
 We also calculated the absolute magnitude (My) using these models., We also calculated the absolute magnitude $M_V$ ) using these models.
 Forcomparison we also calculated a mass of 0.164+0.005M. using the solar metallicity models of Serenellietal.(2001).. which is very similar to the low-metallicity mass determination.," Forcomparison we also calculated a mass of $0.164\pm0.005\ M_\odot$ using the solar metallicity models of \citet{ser2001}, which is very similar to the low-metallicity mass determination."
 However. the cooling age of the system assuming solar metallicity would be between | and 4 Gyrs. which is significantly lower than the cooling age predicted by the low-metallicity models.," However, the cooling age of the system assuming solar metallicity would be between 1 and 4 Gyrs, which is significantly lower than the cooling age predicted by the low-metallicity models."
 The difference in the cooling age can be attributed to theeffect of metallicity on the residual hydrogen burning. which is expected for white dwarfs with such a low mass with solar or low-metallicity progenitors (Serenellietal..2002).," The difference in the cooling age can be attributed to theeffect of metallicity on the residual hydrogen burning, which is expected for white dwarfs with such a low mass with solar or low-metallicity progenitors \citep{ser2002}."
.. Figure 2shows the effective temperature and surface gravity of NLTT 11748 compared to the evolutionary mass-radius relations of Serenelli and Serenellietal. (2002).., Figure \ref{fig_evol} shows the effective temperature and surface gravity of NLTT 11748 compared to the evolutionary mass-radius relations of \citet{ser2001} and \citet{ser2002}. .
 Table 2. lists our adopted values for Tay. log g. M. and My.," Table \ref{tbl_prop} lists our adopted values for $T_{\rm eff}$ , $\log{g}$ , $M$ , and $M_V$ ."
bby bboxes).,by boxes).
 We compared this calibration to that obtained using spectroscopic standard stars and found it to be consistent to within +20% (with no significant systematic offset) for the alignment stars., We compared this calibration to that obtained using spectroscopic standard stars and found it to be consistent to within $\pm 20$ (with no significant systematic offset) for the alignment stars.
" Using the flux calibration, we computed our survey sensitivity as a function of wavelength."," Using the flux calibration, we computed our survey sensitivity as a function of wavelength."
" The 5c limiting line flux is 3.12-0.5x10715 erg cm""? s! a range of Lya line widths typical of our LBG (assumingsamples), implying that we should be able to detect Lyo with frame equivalent widths of greater than 20+3 Που i’-drops with zgso~27."," The $\sigma$ limiting line flux is $\pm
0.5 \times$ $^{-18}$ erg $^{-2}$ $^{-1}$ (assuming a range of $\alpha$ line widths typical of our LBG samples), implying that we should be able to detect $\alpha$ with rest-frame equivalent widths of greater than $20\pm 3$ for $i'$ -drops with $z_{850}\simeq 27$."
 We searched for Lya emission at the spatial position of the targeted LBGs in the Keck spectra., We searched for $\alpha$ emission at the spatial position of the targeted LBGs in the Keck spectra.
 Line fluxes and EWs were calculated following the procedures discussed in Paper I. We account for the effects of line contamination and Lya forest absorption (estimated using relations presented in Meiksin 2006)) on the observed zggo-band fluxes., Line fluxes and EWs were calculated following the procedures discussed in Paper I. We account for the effects of line contamination and $\alpha$ forest absorption (estimated using relations presented in \citealt{Meiksin06}) ) on the observed $_{850}$ -band fluxes.
" Of the 23 i'-band dropouts for which we obtained ultra-deep spectra, 11 show Lya emission, while 2 of the 7 i'-drops for which we obtained 3.67 hour integrations show Lyo (Figure 1))."," Of the 23 $i'$ -band dropouts for which we obtained ultra-deep spectra, 11 show $\alpha$ emission, while 2 of the 7 $i'$ -drops for which we obtained 3.67 hour integrations show $\alpha$ (Figure \ref{fig:lya_idrops}) )."
 These results imply a large fraction of i’-drops have prominent Lyo emission., These results imply a large fraction of $i'$ -drops have prominent $\alpha$ emission.
 The rest-frame EWs for the ;'-drops range between 9.4 and 350A., The rest-frame EWs for the $i'$ -drops range between 9.4 and 350.
. The vast majority of the emission lines are detected with high significance., The vast majority of the emission lines are detected with high significance.
" Even so, we take a conservative cut, limiting our analysis to those sources with rest-frame EWs greater than 25 aand S/N>7."," Even so, we take a conservative cut, limiting our analysis to those sources with rest-frame EWs greater than 25 and $S/N>7$."
 This excises the Lyo detection in the middle right bottom panel of Figure 1.., This excises the $\alpha$ detection in the middle right bottom panel of Figure \ref{fig:lya_idrops}.
" As a result, even among the faintest sources, the emission lines used in our analysis are very confidently detected (<S/N >=18), removing concern regarding spurious features."," As a result, even among the faintest sources, the emission lines used in our analysis are very confidently detected $<S/N>$ =18), removing concern regarding spurious features."
" As in Paper I, we determine the completeness of our Lya detections as a function of wavelength by adding and recovering fake emission lines at random positions across the 2-D spectra."," As in Paper I, we determine the completeness of our $\alpha$ detections as a function of wavelength by adding and recovering fake emission lines at random positions across the 2-D spectra."
 We compute the Lya recovery rate as a function of absolute magnitude and wavelength for all masks observed (including those in Paper I) and make appropriate corrections., We compute the $\alpha$ recovery rate as a function of absolute magnitude and wavelength for all masks observed (including those in Paper I) and make appropriate corrections.
 This test demonstrates that in our deep 12.5 hour mask we are >90% complete to lines with Wryao>50 even for the faintest i'-drops on our mask (zassoc 27).," This test demonstrates that in our deep 12.5 hour mask we are $>$ complete to lines with $_{Ly\alpha,0}>50$ even for the faintest $i'$ -drops on our mask $z_{850}\simeq 27$ )."
" For lines with Wryao~ A,, the completeness implied by our simulations is ~75—80% for sources in the faintest magnitude bin covered by our z~6 spectra."," For lines with $_{Ly\alpha,0}\simeq 20$ , the completeness implied by our simulations is $\simeq 75-80$ for sources in the faintest magnitude bin covered by our $z\simeq 6$ spectra."
" The completeness is of course lower on the DEIMOS mask observed for only 3.67 hours, reaching below ~50% for faint sources with"," The completeness is of course lower on the DEIMOS mask observed for only 3.67 hours, reaching below $\simeq 50$ for faint sources with"
along the orbit of NGC 1404 and. some fraction of them can be clearly seen within the NGC 1899's optical raclius (defined as 5 2. of NGC 1399 in this paper).,along the orbit of NGC 1404 and some fraction of them can be clearly seen within the NGC 1399's optical radius (defined as 5 $R_{\rm e}$ of NGC 1399 in this paper).
 Since the orbit ofNGC 1404 in the adopted spherical gravitational potential forms a rosette. the stripped GC's can be finally distributed widely throughout the cluster scale radius ro (~ 55 kpc).," Since the orbit of NGC 1404 in the adopted spherical gravitational potential forms a rosette, the stripped GCs can be finally distributed widely throughout the cluster scale radius $r_{\rm
s}$ $\sim$ 55 kpc)."
" The overall (projected) distribution. of the stripped. CC's within the central LOO kpe of the cluster can be regarded as rather elongated. which reflects the fact that the NGC 1404s orbit is highly eccentric (orbital eccentricity e, = 0.76)."," The overall (projected) distribution of the stripped GCs within the central 100 kpc of the cluster can be regarded as rather elongated, which reflects the fact that the NGC 1404's orbit is highly eccentric (orbital eccentricity $e_{p}$ = 0.76)."
 The overall distributions also show appreciable inhomogeneity. in particular. for the regions with > 50 kpc.," The overall distributions also show appreciable inhomogeneity, in particular, for the regions with $R$ $>$ 50 kpc."
 During dvnamical evolution of NGC 1404 within the cluster. the GC's are more strongly inlluenced by the cluster idal field than the main stellar component of Νέας 1404. »ecause the GC system is initially less compact.," During dynamical evolution of NGC 1404 within the cluster, the GCs are more strongly influenced by the cluster tidal field than the main stellar component of NGC 1404, because the GC system is initially less compact."
 Fig., Fig.
 2 shows iow this dillerence in the effectiveness of tidal stripping )etween the stellar component and GCs can cause the decrease of Sx in NGC 1404., 2 shows how this difference in the effectiveness of tidal stripping between the stellar component and GCs can cause the decrease of $S_{\rm N}$ in NGC 1404.
 After the first. pericenter ossage. the Sy is reduced [rom 5 to 2.86 (E = OAT). and the reduction rate is nearly proportional to the number of the stripped GC's.," After the first pericenter passage, the $S_{\rm N}$ is reduced from 5 to 2.86 $T$ = 0.47), and the reduction rate is nearly proportional to the number of the stripped GCs."
 GC's can be successively stripped everytime GC. 1404. passes through the cluster core. though the otal number of the stripped ones become smaller in the second and the third. pericenter. passages compared: with dose stripped in the first pericenter. passage (Sy = 2.32 nel 1.58 for 7 = 0.94 ancl 142. respectively).," GCs can be successively stripped everytime NGC 1404 passes through the cluster core, though the total number of the stripped ones become smaller in the second and the third pericenter passages compared with those stripped in the first pericenter passage $S_{N}$ = 2.32 and 1.88 for $T$ = 0.94 and 1.42, respectively)."
 These results early suggest that the origin of the observed smaller Sv ( 2) of NGC 1404 GCs may be closely associated with idal stripping of GC's by the cluster tidal field., These results clearly suggest that the origin of the observed smaller $S_{N}$ $\sim$ 2) of NGC 1404 GCs may be closely associated with tidal stripping of GCs by the cluster tidal field.
 FinalA 1 dynamical evolution. of these stripped GCs becomes eravitationally inlluenced. by the cluster potential rather man by NGC 1404 and therefore can be identified as drifting intracluster GC's.," Finally, the dynamical evolution of these stripped GCs becomes gravitationally influenced by the cluster potential rather than by NGC 1404 and therefore can be identified as drifting intracluster GCs."
 The NGC 1404s tical radius (44). within which collisionless components (dark matter. stellar components. and CC's) o not sulfer so severely from tidal stripping by the cluster. can be analvtically estimated to be 3.72 2. for the fiducial model with e; = 0.76 and the total mass of NGC 1404 equal to 2.64 104 A£..," The NGC 1404's tidal radius $R_{\rm t}$ ), within which collisionless components (dark matter, stellar components, and GCs) do not suffer so severely from tidal stripping by the cluster, can be analytically estimated to be 3.7 $R_{\rm e}$ for the fiducial model with $e_{\rm p}$ = 0.76 and the total mass of NGC 1404 equal to 2.64 $\times$ $10^{11}$ $M_{\odot}$."
 The initial total number of GCs within Z2 is 288 for the fiducial mocel. and accordingly 1062 GCs are expected to be tidally stripped from. NGC 1404.," The initial total number of GCs within $R_{\rm t}$ is 288 for the fiducial model, and accordingly 1062 GCs are expected to be tidally stripped from NGC 1404."
 The final number of GCs within /4 is 253 (T = 1.42 Gyr) in the simulation. and accordingly LOOT GCs are tically stripped from NGC 1401.," The final number of GCs within $R_{\rm t}$ is 253 $T$ = 1.42 Gyr) in the simulation, and accordingly 1097 GCs are tidally stripped from NGC 1404."