ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 star., star.3 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."4 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.5 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."6 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.7 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."8 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.9 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).10 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.11 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.12 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.13 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.14 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 $\tau_{\rm15 thermal}$ is the secondaries thermal timescale."16 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.17 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.18 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."19 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}."20. We consider the secondary in more detail in Section 2.3., We consider the secondary in more detail in Section 2.3.21 We list our models that match 5? Velorum in Table I.., We list our models that match $\gamma^{2}$ Velorum in Table \ref{initialp}.22 The range of initial primary masses and initial separations are narrow., The range of initial primary masses and initial separations are narrow.23 “Phe initial separations are similar to the observed: separation., The initial separations are similar to the observed separation.24 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.25 For each of the svstems a binary interaction mus jwe occurred., For each of the systems a binary interaction must have occurred.26 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.27 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."28 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.29 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 0.6.30 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.31 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"32reveals a line profile very similar to that first discovered by T95.,reveals a line profile very similar to that first discovered by T95.33 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."34 1998. 1999).," 1998, 1999)."35 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},"36 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.37" 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 $R_{\rm in} = 3.438R_{\rm S}$ ) 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 $\propto39R^{-\alpha}$ of the line."40" 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 $\chi^2 \, / d.o.f =4142/47$ , and Fe $\alpha$ flux $(1.37 \pm 0.29) \times 10^{-4} \rm \,42ph \, cm^{-2} \, s^{-1}$, with equivalent width $W_{\rm K\alpha} \sim43295 \pm 118$ eV. This is in good agreement with the broad iron line measured using the data which has $W_{\rm K\alpha} = 398 \pm4458$ eV, comparable to previous (e.g. Lee et al."45 1998. 1999) and measurements of this source (e.g. Iwasawa et al.," 1998, 1999) and measurements of this source (e.g. Iwasawa et al."46 1996. 1999).," 1996, 1999)."47" The apparent ""sharp drop’ at 6.7 keV (Fig. 3))", The apparent `sharp drop' at 6.7 keV (Fig. \ref{fig-hegasca}) )48 appears resolved (AE70 eV from peak to drop)., appears resolved $\Delta E > 70$ eV from peak to drop).49" 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 $E_{\rm K\alpha(obs)} = 6.21 \pm500.08$ keV, Gaussian width $\sigma = 0.66 \pm 0.11$ (FWMH $\sim5173,000$ ), and iron line flux $I_{\rm K\alpha} = (2.13 \pm 0.31)52\times 10^{-4} \, \rm \ph \, cm^{-2} \, s^{-1}$."53 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.54 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 $\sim 20-30 \rm \55cts\, bin^{-1}$ (where bin =, or an ACIS pixel) and therefore sufficient for statistically meaningful errors derived from spectral fitting."56 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}) ).57 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."58 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). ("59The 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 $> 99.9$ .)60 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)."61 Wilms et al. (, Wilms et al. (62"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. (632002) attribute this to the blue wing of the disk line.,2002) attribute this to the blue wing of the disk line.64 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."65 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}) )"66 of the time averaged 125 ks CoacoivomdsIOergem?s y observation.," of the time averaged 125 ks $f_{\, \rm (0.4-10~keV)} = 674 \times 10^{-11} \, \rm erg\,cm^{-2}\,s^{-1}$ ) observation."68 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}."69 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.70 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.71 This is depicted in Fig., This is depicted in Fig.72 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."73 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 $\approxlt 274\sigma$ ) significant."75Active 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.76 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}."77 X-ray emission is especially a characteristic property of AGNs., X-ray emission is especially a characteristic property of AGNs.78 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.79 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}."80 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."81 Colour selection is a powerful technique in extracting AGN candidates., Colour selection is a powerful technique in extracting AGN candidates.82 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."83 Richardsetal.(2002). selected quasars via their nonstellar colours using SDSS photometry., \citet{Richards2002-AJ} selected quasars via their nonstellar colours using SDSS photometry.84 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}. ."85Colour 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.86 extracted obscured AGN candidates using the colour selection of (JAvy)2.0., extracted obscured AGN candidates using the colour selection of $(J-K_\textnormal{\tiny S})>2.0$.87 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}."88 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)."89 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."90 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."91 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).92 This enables us to extract AGN candidates using only near-infrared photometry., This enables us to extract AGN candidates using only near-infrared photometry.93 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}."94 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."95 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"96areal 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.97 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)."98 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."99 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."100 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."101 If the break mass reflected. sav. the local Jeans mass or some," If the break mass reflected, say, the local Jeans mass or some"102Dust 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.103 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."104 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}."105.. 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."106 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."107 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."108 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}. ."109 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."110 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.111 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."112 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."113 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."114 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.115 In addition. deposition JJOCCSSesS COILIC Lact as filters.," In addition, deposition processes could act as filters."116 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."117 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."118 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, $R_{V} = 119A_{V} / E(B - V)$ 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}."120. 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.121"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."122" 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."123 1991)., 1991).124" These peaks correspond successively to singly, doubly (etc) scattered photons (see Figs."," These peaks correspond successively to singly, doubly (etc) scattered photons (see Figs."125 4a-4d in Loeb et al., 4a-4d in Loeb et al.126 1991)., 1991).127" 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."128 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}) )129" decreases strongly and, therefore, the Eq. (5))"," decreases strongly and, therefore, the Eq. \ref{G}) )"130" 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 $G(x,131T_{\mathrm{e}})$ is given by The sub-exponential function $f(s)=-3 s-x \exp(-s)$ has a maximum at the frequency shift $s_{\mathrm{max}}=\ln(x/3)$."132 Since exp(xexp(—Smax))>>1 the approximate expression for the generalized spectral function is valid and we calculate the integral in Eq. (22))," Since $\exp(x133\exp(-s_{\mathrm{max}}))\gg1$ the approximate expression for the generalized spectral function is valid and we calculate the integral in Eq. \ref{Ghf}) )"134 by the Laplace’s method., by the Laplace's method.135 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. (136"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."137" 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."138" 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 $s$."139" 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 $P_1 (s,140T_{\mathrm{e}})$ at the temperature of $k_{\mathrm{b}}141T_{\mathrm{e}}=500$ keV is shown in Fig. \ref{P1}."142" 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)."143" 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 $G(x,144k_{\mathrm{b}} T_{\mathrm{e}}=500$ keV) on the dimensionless frequency $x$ is illustrated in Fig. \ref{500}."145 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.146" 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."147 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.148" 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."149" 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.,"150 2005)., 2005).151" 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 $/$ "152birthplace indicating a NS age ~5«LO? ves (Walter&Lattimer2002:Kaplanetal. 2002).,"birthplace indicating a NS age $\sim5\times10^5$ yrs \citep{walter02, kaplan02}."153. 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.154 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}."155. 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}."156. 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.157 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."158 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.159 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."160" 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."161" 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 $v_m$."162 We fixed ο from the observed stellar proper motion at 200 kant., We fixed $v_m$ from the observed stellar proper motion at 200 $^{-1}$.163 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. [164Thermal 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.165" 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).166 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]."167 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.168" 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}) )."169 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}.170". 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 $\dot{M}=\pi171R_m^2 v_0 \rho_m = 5\times10^7R_{m,12}^2 v_{m,7} n_m$ $^{-1}$."172Llowever. 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}."173 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}.174. 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}."175. 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}.176. 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}."177.. 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}.178. 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}."179. A similar situation exists for magnetized heavy clement atmospheres (Rajagopaletal.1997)., A similar situation exists for magnetized heavy element atmospheres \citep{rajagopal97}.180. 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}.181. 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}."182. 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."183 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}."184. 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.185 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 $\rho\sim186560Z^{-3/5}B_{12}^{6/5}$ $^{-3}$, with $B=10^{12}B_{12}$ G) has essentially no emissivity for O-mode X-rays but is near blackbody for E-mode ones."187⋟∪↥⋅↑↕∐∖↥⋅⋜∥∐⋜↧↑↕∐∶↴⋁↴∖↴↿∐⋅↕⋟⋜↧↸⊳↸∖↕↴∖↴∿∶≩∩⋅↱↗∩⋰⋰⊒∶∕ ≼∐∖↻↸∖∐≼∐∐∶↴⋁∪∐↻∐∪↑∪∐↸∖∐↸∖↥⋅∶↴∙⊾⋅↖↽⋜⋯≼↧↕⊔⋜↧∶↴∙⊾∐↸∖↑↕↸⊳∐↸∖↕≼↧∶↴∙⊾↸∖∪∐∐∖⊓⋅⋅↖↽ (Zaueetal.2003)., Then the average emissivity for the radiating surface is $\sim$ depending on photon energy and magnetic field geometry \citep{zane03}.188. 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}.189 Will we find other isolated NSs in the propeller phase?, Will we find other isolated NSs in the propeller phase?190 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 $\tau_{csd}$ ).191 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. ("192Alternatively. 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.)"193" 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 $n_m \gg 1$ $^{-3}$ )."194 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.195 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).196 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}. ."197 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 after198and are likely also in lower mass svstems.,and are likely also in lower mass systems.199 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."200 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."201 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.202 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."203lima 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 $\lim_{r\to \infty}f^-_n(r,\nu^2)=0$."204"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."205tions of , This is the case for potentials of finite range.206summation ancltaking, But then again we have to sum over $n$ and this sum will be logarithmically divergent.207 limit do notcommute., The renormalization for this case as forthe Green' s function method is discussed insubsection \ref{renorm}. .208or TESS.,or HESS.209 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}."210. 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}.211. 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}."212 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 $\gtrsim 10^{39}\,\rm erg\, s^{-1}$."213 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}214 from the direction of the GCBH, and is unlikely to be related to Sgr $^\ast$."215 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}."216. 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.217 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."218 (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."219 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}) )."220 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."221 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}."222. 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.223 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."224 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}."225 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}))."226 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"227refer 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.228 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.229 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.230 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."231" 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."232 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.233 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."234 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."235" 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 $\vec{\nabla} \kappa236=\hat{D}\gamma237\equiv \vec{u}_{\gamma}$, 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)."238 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 $\gamma(\vec{\theta})$."239 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."240 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., $\left<\kappa\right>=\int\,d^2\theta'\,241W(|\vec{\theta}-\vec{\theta}'|;\vartheta)\,\kappa(\vec{\theta}')$, where $W(\theta;\vartheta)$ is a smooth, continuous window function with a characteristic scale of $\vartheta$ ."242 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 $\gamma$ : $\vec{\nabla}\left<\kappa\right>=\hat{D}\left<\gamma\right>243\equiv \vec{u}_{\left<\gamma\right>}$; $(\triangle\left<\kappa\right>,0)=244\hat{D}^2\left<\gamma\right>=\hat{D}\vec{u}_{\left<\gamma\right>}245=({\rm div}\vec{u}_{\left<\gamma\right>},246{\rm rot}\vec{u}_{\left<\gamma\right>})$ ."247 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."248 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 $\vec{u}_{\left<\epsilon\right>}249:=\hat{D}\left<\epsilon\right>$, 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)."250 On the other hand. τονc measures the ‘pure’ noises," On the other hand, ${\rm rot}\vec{u}_{\left<\epsilon\right>}$ measures the `pure' noise."251" 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 $\hat{D}\vec{u}_{\left<\epsilon\right>}=252({\rm div}\vec{u}_{\left<\epsilon\right>},253 {\rm rot}\vec{u}_{\left<\epsilon\right>})=254-n_{\rm g}^{-1}\sum_{m=1}^{N_{\rm255g}}p(|\vec{\theta}-\vec{\theta}_m|;\vartheta)\,256[\epsilon_{\rm t}(\vec{\theta}_m;\vec{\theta}),257 \epsilon_{\rm r}(\vec{\theta}_m;\vec{\theta})] 258$ , 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 $\epsilon_{\rm r}(\vec{\theta};\vec{\theta}_0):=-259\Im[\epsilon(\vec{\theta})e^{-2i\phi}]$ with $\phi= {\rm Arg}(\vec{\theta}-\vec{\theta}_0)$ and $p(\theta;\vartheta)=W''(\theta;\vartheta)-W'(\theta;\vartheta)/\theta$."260" 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 $\sigma^2_{ij}:={\rm E}[261\hat{D}^2\left<\epsilon^{\rm (s)}\right>\,262\hat{D}^2\left<\epsilon^{\rm (s)}\right>]_{ij}$ : where $\sigma_{\epsilon}$ is the dispersion of the intrinsic source ellipticities and we have assumed ${\rm E}[\epsilon^{\rm (s)}(\vec{\theta}_m)263\epsilon^{\rm (s)}(\vec{\theta}_n)]_{ij}=\delta_{ij}\delta_{mn}264\sigma_{\epsilon}^2/2$ ."265 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.266 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}] ])267 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.268 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 $W_{\rm269G}(\theta;\vartheta)=\exp(-\theta^2/\vartheta^2)/\pi\vartheta^{2}$, in which case $p(\theta;\vartheta)=4(\theta/\vartheta)^2270\exp(-\theta^2/\vartheta^2)/\pi\vartheta^4\equiv p_{\rm271G}(\theta;\vartheta)$; $p_{\rm G}(\theta;\vartheta)$ has its maximum at $\theta =\vartheta$ and falls off rapidly at $\theta>\vartheta$."272" 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."273 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}) )."274 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."275 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."276 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 $\left<z\right>=z_0\Gamma(4/\beta)/\Gamma(3/\beta)$."277 Tn the presen Letter. we consider ouly the case (625.7)=(1.0.1.0).," In the present Letter, we consider only the case $(\left<z\right>, \beta)=(1.0, 1.0)$."278 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 $\vec{u}_{\gamma_{\infty}}=\hat{D}\gamma_{\infty}$ ."279 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"280result does not come from a circular argument.,result does not come from a circular argument.281" 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."282 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."283 Ht is also worth noting. as in Fig. 11((," It is also worth noting, as in Fig. \ref{hz}( ("284"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."285 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."286 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."287 The latter argument is perhaps (he most likely., The latter argument is perhaps the most likely.288 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."289 2009)., 2009).290 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.291 2007: et al., 2007; Men{\'e}nndez-Delmestre et al.292 2009)., 2009).293 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."294 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."295 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."296 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.297 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.298 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"299 2=D ;z20/—230. o ~10°AL. 26. τοL. à Z32 1.Syan 210. ~1000 arcmin7. Lya , $z=5$ $z\approx 20-30$ $\sigma$ $\sim 10^5~{\rm M_\odot}$ $z\gsim 6$ $z\gsim 4$ $\alpha$ $\gsim 32$ $1-5\mu$ $z\gsim 10$ $\sim1000$ $^{-2}$ $\alpha$ 300Ry 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.301 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."302 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."303 We would like to thank the anonymous referee for useful suggestions., We would like to thank the anonymous referee for useful suggestions.304ancl 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."305 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."306 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)."307 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)."308 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 $r_{2500}$."309 “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."310 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."311 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.312 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."313 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."314 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).315 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."316 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."317 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 $r318\sim 0.2 r_{2500}$ Allen 2001a,b,c; David 2001; Schmidt 2001)."319 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}.320 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."321 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."322 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.323 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, $\Omega_{\rm m}$ ."324 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.325 SWA and ACE acknowledge the support. of the Itoval Society., SWA and ACF acknowledge the support of the Royal Society.326distribution modified by cluster disruption processes.,distribution modified by cluster disruption processes.327 We recall that based on our magnitude-Iimited sample analysis (Fig., We recall that based on our magnitude-limited sample analysis (Fig.328 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, $\alpha =329-2$, expected for young star cluster systems."330 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)."331 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."332 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.333 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."334" 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."335 2005: Mengel et al., 2005; Mengel et al.336 2005: see also Tremonti ct al., 2005; see also Tremonti et al.337 2001)., 2001).338 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}."339 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."340 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."341 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."342 2005)., 2005).343 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."344 S. break down., \ref{clfs2.fig} break down.345 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."346 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."347 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."348 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."349 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.350 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."351 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."352" ""Thus. in addition to our conclusion in de Cirijs et al. ("," Thus, in addition to our conclusion in de Grijs et al. ("3532005) 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 $\Delta \langle \log( {\rm Age / yr} ) \rangle354\le 0.35$ 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."355 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)."356 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."357 An alternative triggering mechanism for the voung(er) clusters may be neeceel., An alternative triggering mechanism for the young(er) clusters may be needed.358" 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 $\log(t_4^{\rm dis}/{\rm yr}) =3599.9 \pm 0.1$ , where $t_{\rm dis} = t_4^{\rm dis} (M_{\rm cl}/10^4 {\rm360M}_\odot)^{0.62}$ , for the LMC cluster system."361 This is consistent with earlier. preliminary," This is consistent with earlier, preliminary"362however it is double-valued at low abundances.,however it is double-valued at low abundances.363" 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."364 This provides a poorer correction than the Ha to Hf ratio., This provides a poorer correction than the $H\alpha$ to $H\beta$ ratio.365" 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}."366 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.367" This is a first approximation, not a choice dictated by theory."," This is a first approximation, not a choice dictated by theory."368" 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."369 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}) ).370 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.371" 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."372 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.373 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.374" 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."375" 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."376" 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."377 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.378" 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."379" 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}."380" 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."381 '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.382 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.383 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.384" 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."385" 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."386 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.387" 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."388 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.389" 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}) )."390" 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}."391 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.392" 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."393" 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."394" 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."395" 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."396 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.397 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.398" ?) 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}."399" 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}."400" 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."401" 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."402" 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."403 Are the polynomial fits really better than the traditional straight line fits?, Are the polynomial fits really better than the traditional straight line fits?404" 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."405" 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."406" 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."407" 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"408"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 $\AA$.409 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."410 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."411" This task is beyond the scope of this, rather qualitative, analysis."," This task is beyond the scope of this, rather qualitative, analysis."412" 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."413 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.414 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.415 From Fig. 7..," From Fig. \ref{Ha_orig},"416 we estimate that spectroscopic eclipse began around HID 2454000., we estimate that spectroscopic eclipse began around HJD 2454000.417 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 $\sim$ $^\circ$.418 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.419 In Fig. 12..," In Fig. \ref{Roche},"420" 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},"421 and the sightlines at the onset of spectroscopic and photometric eclipse., and the sightlines at the onset of spectroscopic and photometric eclipse.422 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.423 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."424" 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}."425 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.426 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."427 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."428 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},"429 also Fig. 2))., also Fig. \ref{ha_fig1}) ).430" 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."431 By Mar 2007. the final spectrum of this series," By Mar 2007, the final spectrum of this series"432"Notice, though. the three separately tageed UCDs iu Figure 6..","Notice, though, the three separately tagged UCDs in Figure \ref{fig:amr}."433 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."434 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.435 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.436 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."437 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}.438 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."439 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).440 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}) )."441" 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."442 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."443"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)."444 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."445 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."446 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."447" 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."448 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.449 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."450 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."451 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."452 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.453 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."454 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."455 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."456 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."457" 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}) )."458 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}) )."459" We refer to these as merged GCs,", We refer to these as merged GCs.460 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}) )."461 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}) )."462" 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}."463 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."464 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."465 We restrict the suuple to objects with distances confined either bv, We restrict the sample to objects with distances confirmed either by466simulations.,simulations.467 The low mass loss rate simulation has Alyw=5OAL.ve +., The low mass loss rate simulation has $\Mdot_{\rm W} = 5 \times 10^{-5} \Msol \pyr$ .468 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."469 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.470 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 $n_{0}= 10 \pcc$."471 We shall only consider simulatedROSAP PSPC (Position Sensitive Proportional Counter) data., We shall only consider simulated PSPC (Position Sensitive Proportional Counter) data.472 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."473 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).474 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."475 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.476 There are three definable stages of bubble growth. seen in the simulation., There are three definable stages of bubble growth seen in the simulation.477 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 $T \sim 10^{4}478\K$ ) shell."479 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).480 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 {\rm481(K)} \ltsimm 6.0 $ ."482 The shell is thick (see Figs., The shell is thick (see Figs.483 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}) )."484 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."485 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."486 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 $t \approx 6300487\yr$ as the shell begins to cool."488 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."489 The peak UV-IZUVM. luminosity of 9.110?eres briefly exceeds the wind energy input.," The peak UV-EUV luminosity of $9.1 \times 10^{37}490\ergps$ briefly exceeds the wind energy input."491 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."492 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 $T \approx49310^{4} \K$."494 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."495 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)."496 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."497 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}-498-2 and the radial profiles X Pie. 4.., \ref{fig:temp_4t} and the radial profiles of Fig. \ref{fig:weaver}.499 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.500 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}) ),"501 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."502 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."503 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."504 The instabilities of the shell seen in Figs., The instabilities of the shell seen in Figs.505 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."506 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)."507 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."508towards the west.,towards the west.509 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."510 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."511 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}.512. 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.513 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.514 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)."515 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."516" 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 $\sim4\arcsec-7\arcsec$."517 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.518" 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."519 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."520 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}."521. 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}) ).522" 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."523 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}) ).524 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."525 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 $\Gamma_0$."526 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."527 Assuming acdiabatic expansion. the evolution of each jet is governed by:," Assuming adiabatic expansion, the evolution of each jet is governed by:"528removes 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."529 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."530 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."531 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}) )."532 Thus. high. A can help to extract more RV information.," Thus, high $R$ can help to extract more RV information."533 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}) )."534 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.535 (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."536 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}) )."537 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)."538 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 $R$ .539 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}.540 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 $R$."541 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."542 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 $R$."543 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."544 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), $Q_{\rm{DFDI}}\propto R^{0.63}$."545" 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."546" 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."547" 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."548"Therefore. 6v,x10.630.5Hk113 for DEDL","Therefore, $\delta v_{rms}\propto R^{-0.63-0.5}=R^{-1.13}$ for DFDI."549 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."550 The power law is, The power law is551the 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.552 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.553 1998: Fender et al., 1998; Fender et al.554 1999 and references therein)., 1999 and references therein).555 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-1.556 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.557 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."558 ‘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 11-20.559 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.560 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."561" 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."562 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."563 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.564 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)."565 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."566 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."567 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.568 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."569 LOOT: Fender Pooley 1998)., 1997; Fender Pooley 1998).570" 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.571 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.572 s pointed out by Cotton et al. (, As pointed out by Cotton et al. (573"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.574 Alarscher Gear (1985) and. O'Dell οἱ al. (, Marscher Gear (1985) and O'Dell et al. (5751988) 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.576 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).577 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."578 Alternatively. Wang ct al. (," Alternatively, Wang et al. ("5791997) 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.580 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"581Studies 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.582 2002)., 2002).583 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.584 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."585" 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."586 1998)., 1998).587 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."588 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.589" 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.590" 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.)"591 have been applied to the sample definition., have been applied to the sample definition.592"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."593 The above list can be regarded as a set ofrequirements., The above list can be regarded as a set of.594 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."595 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."596 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."597 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)."598 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."599 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."600 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.601 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 $0.05 \mu \rm m$."602 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."603 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."604 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"605where 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.606.. However. the function e depends on the solution itself.," However, the function v depends on the solution itself."607 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."608 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."609 Clearly. means u4.," Clearly, means."610. 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 ( )."611 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."612 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}."613 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}614 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."615 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.)616 is weak., is weak.617 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"618is 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.619 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."620" 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)."621 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."622 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.623" 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)."624 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)."625 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}."626. 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)."627 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."628 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."629 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}."630 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."631 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.632 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}."633. 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."634 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}."635- 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."636 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}."637. 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."638 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."639 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."640 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}."641. 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."642 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."643eraius 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}."644. 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)."645 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}.646. 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}."647. 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 $S$ )."648" 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."649 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."650 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."651 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."652 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."653" 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 $2\beta_b654/ 3$, respectively."655 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."656" There is the obvious constraint that 2,/3+295/53=1.", There is the obvious constraint that $\beta_a/3 + 2\beta_b / 3 = 1$.657" 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 $A_b = 8.78 \times 10^{-7}$ $^{-1}$ $2p_{3/2}$ $2s_{1/2}$ )."658 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 $2p_{3/2}$ ; $g_{2s} = 2$ )."659" 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."660 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 $A_{21}$ .661 The, The662"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."663" 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."664 Therefore. we show only oue representative example in this Letter.," Therefore, we show only one representative example in this Letter."665" 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."666" 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}."667 Our example would imply (hat. fe is close to the photosphere., Our example would imply that $R$ is close to the photosphere.668 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 $\gtrsim E_{\rm sh}/(20 \pi R^3 m_{\rm p} c^2)$."669 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.670 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.671 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."672 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 $K^0_{\rm S}$."673 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."674 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."675 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.676 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."677 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}. .678 Svncehrotron and inverse Compton emissions are taken into account., Synchrotron and inverse Compton emissions are taken into account.679" 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 $\pi^+ \pi^- \pi^0$ )."680 In 3. total neutrino spectra emitted from this example are shown.," In Figure 3, total neutrino spectra emitted from this example are shown."681" 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"682is complex (with five components).,is complex (with five components).683 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."684 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."685 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."686 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.687" 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.688 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."689" 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."690 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."691 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."692 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.693 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 $S/N>30$."694 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."695 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."696 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.697 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.698 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.699 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 are700its one dav alias.),its one day alias.)701 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."702 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.703 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 $^{-1}$. ("704Nb.,Nb.705 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.)706 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.70734... 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"708or each giveu mass accretion rate until we achieved a best-fit.,for each given mass accretion rate until we achieved a best-fit.709" 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."710 The 4Z for his fit was 3.11., The $\chi^{2}_{\nu}$ for this fit was 3.41.711 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).712 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."713 Therefore. again. we rejected this resut.," Therefore, again, we rejected this result."714 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."715 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."716 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.717 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."718 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.719 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.720 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.721 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."722 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.723 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.724 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.725 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 8.726 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.727 2003: Sion et al., 2003; Sion et al.728 20()3)., 2003).729 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."730 The white dwarf appears to have a rotatioua| velocity of GOO4100 s5., The white dwarf appears to have a rotational velocity of $600 \pm 100$ $~$ $^{-1}$.731 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.732 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."733 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."734 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."735 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 rejected736sun )) can be made with Lipparcos measurements based on motions of Cepheicls.,Sun ) can be made with Hipparcos measurements based on motions of Cepheids.737 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)."738 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 $29.45\pm0.15$."739. 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.740 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."741 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."742" 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}. ."743" 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 $\approx3$.744 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.745. 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)."746 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.747 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."748. 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"749"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."750" 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."751" 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 $S^{10\mu{\rm m}}_{{\rm752 peak}}$, supports the wide usage of $S^{10\mu{\rm m}}_{{\rm753 peak}}$ as a proxy for dust size in literature \citep{VB03,KE06,PA09}."754". 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)."755" 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."756 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."757" 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."758 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).759"As discussed by Oliveira et al.,"," As discussed by Oliveira et al.,"760 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.761 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.762" 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."763" 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}."764" 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'."765" 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}."766". 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."767" 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}."768. 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).769" 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."770" 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."771" 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)."772" 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."773 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 disks774"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."775 2007)., 2007).776 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."777 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."778 2002)., 2002).779 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."780 2007: Cain et al., 2007; Cain et al.781 2008)., 2008).782 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.783 Torri et al. (, Torri et al. (7842004) 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 $\sim785 5$."786 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)."787 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."788 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."789 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.790 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."791 This contrasts with the results of Gladders et al. (, This contrasts with the results of Gladders et al. (7922003) 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.793 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."794" 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."795 Thus. the large Einstein radius of Abell 1689 is probably not unique.," Thus, the large Einstein radius of Abell 1689 is probably not unique."796 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.797 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."798 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."799 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."800 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.801 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.802 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."803 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."804 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.805 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."806 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."807 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."808 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 cosmological809are also strong sources at COAL?PEL energies (von Montigny οἱ al.,are also strong sources at COMPTEL energies (von Montigny et al.810 1995. et al.," 1995, Mukherjee et al."811 1997)., 1997).812 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."813 1997 ancl references therein)., 1997 and references therein).814 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.815 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.816production 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)817of sight velocities to velocities relative to the Local Group (Yahiletal.1977).,of sight velocities to velocities relative to the Local Group \cite{yah77}.818.. 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. (8191990) and de Blok et al. (,1990) and de Blok et al. (8201996) 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}.821. 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.822 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.823 The mean correction was 0.2 magnitudes., The mean correction was 0.2 magnitudes.824 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.825 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 $(M_{HI}/L_{B})_{\odot}>0.3$ ).826 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}.827 1n figure 7 we show (Mgi/Lg). plotted against Alp., In figure \ref{fig:mvml} we show $(M_{HI}/L_{B})_{\odot}$ plotted against $M_{B}$.828 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 $(M_{HI}/L_{B})_{\odot}$.829 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. (8301992) 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."831 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 $(M_{HI}/L_{B})_{\odot}$.832 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."833The 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.834 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}.835 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."836 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."837 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}."838 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."839 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.840 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.841 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."842 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}."843 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"844oue 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.845 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.846 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."847 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.848 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.849 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.850 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.851 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.852 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}{ (8532002) 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.854 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 - $\alpha^1({\bf x}) = \gamma \left[ x^1\cos(2\theta_\gamma) +855x^2\sin(2\theta_\gamma)\right]$, $\alpha^2({\bf x}) = \gamma \left[856x^1\sin(2\theta_\gamma) - x^2\cos(2\theta_\gamma) \right]$."857 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.858 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.859 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.860 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}861 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."862" 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 $\Sigma_c = 2.65\times 10^9863h_{65}\msun\kpc^{-2}$."864 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."865 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.866 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."867 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."868 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 $\gamma=0.102$.869 This velocity dispersion is not unusual for a leus galaxy., This velocity dispersion is not unusual for a lens galaxy.870 The estimated circular velocity is Vane=V20 yar.," The estimated circular velocity is $V_{\rm871circ}=\sqrt{2}\sigma_{\rm host}$ ."872 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.873 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.874 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.875 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.876 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."877 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."878Observations 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.879 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."880 ‘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)."881 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).882 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."883 2004: Coorayv οἱ al., 2004; Cooray et al.884 2004). is still possible.," 2004), is still possible."885 ]xashlinsky et al. (, Kashlinsky et al. (8862005) 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.887 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."888 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)."889 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.890 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.891 2005)., 2005).892 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."893 The residual large scale signal has been ascrixxb by Washlinsky ct al. (, The residual large scale signal has been ascribed by Kashlinsky et al. (8942005) 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 $^{-2}$ $^{-1}$.895 ‘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.896 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. ("8972005).,2005).898 Section 4 summarizes our conclusions., Section 4 summarizes our conclusions.899"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."900" The subscript ""IV in ή denotes relativistic Bremsstrallune."," The subscript “R” in $Q_{\rm ff,R}$ denotes relativistic Bremsstrahlung."901" 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 $c_s^2\approx 2c_{sp}^2$ ."902":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."903 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."904 We therefore neglect svnchrotron emission in our analvsis., We therefore neglect synchrotron emission in our analysis.905 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)."906 For simplicity. we neglect also thermal conduction.," For simplicity, we neglect also thermal conduction."907 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."908 Our simplified hydrodynamic model is. therefore. very instructive.," Our simplified hydrodynamic model is, therefore, very instructive."909 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."910 Oncethe basic, Oncethe basic911this 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."912 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."913 2002). weakening the case for a very steep continuum.," 2002), weakening the case for a very steep continuum."914 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."915 Combining previous ASCA and BeppoSAX results (Collinge Brandt 2000; Iwasawa et al., Combining previous ASCA and BeppoSAX results (Collinge Brandt 2000; Iwasawa et al.916 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."917 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.918 1998)., 1998).919 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).920 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.921 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."922 2001a: Fiore et al., 2001a; Fiore et al.923 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."924 1997)., 1997).925 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."926 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."927 2001)., 2001).928 Let us call this matter. for simplicity. the ‘torus’.," Let us call this matter, for simplicity, the `torus'."929 There are several possible solutions to this problem (see Matolino et al., There are several possible solutions to this problem (see Maiolino et al.930 2001b for a discussion)., 2001b for a discussion).931 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."932 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."933 2001 and references therein)., 2001 and references therein).934 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."935 1998)., 1998).936 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."937 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."938 1997)., 1997).939 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."940 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."941 An alternative solution. proposed by Matolino et al. (," An alternative solution, proposed by Maiolino et al. ("9422001b). 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."943 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."944 200109)., 2001b).945 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."946 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."947 2000) and therefore at least part of them must be reflected., 2000) and therefore at least part of them must be reflected.948 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.949" 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."950 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."951 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."952 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."953 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."954 2000 and Iwasawa et al., 2000 and Iwasawa et al.955 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."956 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.957 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."958 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."959We 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.960 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."961 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."962 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."963 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.964 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."965 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.966" 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."967 MIST. acknowledges support from NSERC aud a Premiers Research Excellence Award., MJH acknowledges support from NSERC and a Premier's Research Excellence Award.968To 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 $\xi$ ).969 For this estimation we have used the standard method., For this estimation we have used the standard method.970 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.971 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.972 We tried different & values to fulfill this requirement., We tried different $\xi$ values to fulfill this requirement.973 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.974 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.975 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."976 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}.977. 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.978 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 $<$ 0.01.979. 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}."980 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.981 To estimate errors for our WIDTH metallicities we consider the following facts., To estimate errors for our WIDTH metallicities we consider the following facts.982 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."983 We assume a error due to the continuum level determination., We assume a error due to the continuum level determination.984 This translates into maximum uncertainties in the metallicity estimation., This translates into maximum uncertainties in the metallicity estimation.985 The atomic constants may also have uncertainties., The atomic constants may also have uncertainties.986 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.987" 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."988 We derived a median difference of 0.20 dex., We derived a median difference of 0.20 dex.989 The largest difference corresponds to HD 28978 (0.55 dex)., The largest difference corresponds to HD 28978 (0.55 dex).990 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.991 For each object. we need to identify and measure many spectral lines.," For each object, we need to identify and measure many spectral lines."992 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}.993. 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.994 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 $\xi$ )."995 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"996be a precious information for people dealing with stellar evolution and. nucleosynthesis.,be a precious information for people dealing with stellar evolution and nucleosynthesis.997 Finally. we address the problem of +°O) destruction in stars.," Finally, we address the problem of $^{17}$ O destruction in stars."998 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)."999 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."1000 In Fig., In Fig.1001 9a.b we compare results from Mocellines) to those [rom Aloclelfines}.," 9a,b we compare results from Model to those from Model."1002 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."1003 A realistic situation should be probably something in between the two., A realistic situation should be probably something in between the two.1004 Assuming that all the IO present in the gasout, Assuming that all the $^{17}$ O present in the gasout1005"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."1006 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 $\alpha_{t_0}$."1007 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."1008 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."1009 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."1010 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}) )."1011 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 $n_{n t}$ )."1012" 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."1013" 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."1014 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.1015 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."1016" 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."1017" 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."1018 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 $\alpha_{t_0}$."1019 Additionally. it is easy to," Additionally, it is easy to"1020sub-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.1021" 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."1022 These results are summarized in Appendix A., These results are summarized in Appendix \ref{sec:appA}.1023 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.1024 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.1025 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)."1026 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."1027 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.1028 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)."1029 Our findings can be summarized as follows., Our findings can be summarized as follows.1030 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 BP82.1031 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)."1032 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/).1033 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."1034 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.1035 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."1036 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. (1037Note. 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}. .)"1038 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."1039 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."1040 Our treatment has been deliberately, Our treatment has been deliberately1041To 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.1042 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."1043 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.1044. 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.1045 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.1046 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.1047 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."1048For 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."1049 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.1050" 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."1051. The enission of a soft photon is a quasi-classical process., The emission of a soft photon is a quasi-classical process.1052 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 $."1053 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.1054 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 $.1055 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 $%1056d\sigma _{0 describing the scattering of quarks."1057 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."1058Irfirstorder1/.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 $%10591/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 $\sigma _{0}(s)\sim 1/s$ ."1060. The cross section decreases with increasing energv., The cross section decreases with increasing energy.1061 For a more general paranmetrization of (he, For a more general parametrization of the1062"]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).1063 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."1064: Lt is: à comparable size: to observational. errors (Bower. Lucey Ellis 1992).," It is a comparable size to observational errors (Bower, Lucey Ellis 1992)."1065 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."1066 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)."1067 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."1068 Colour of cach galaxy, Colour of each galaxy1069"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 $0.4''$.1070" 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."1071 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).1072 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."1073 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)."1074 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 2.1075 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)."1076 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.)."1077 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 MM2.1078 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)."1079" 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 $0.6''$."1080 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).1081 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)."1082 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."1083 Line identification was. performed following the methodology described below., Line identification was performed following the methodology described below.1084 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."1085 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."1086 3-left)., 3-left).1087 Second. we computed rotational diagrams (Fig.," Second, we computed rotational diagrams (Fig."1088 4) with the ttransitions. allowing us to determine the gas temperature," 4) with the transitions, allowing us to determine the gas temperature"1089"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."1090" 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}."1091" 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}) )."1092" 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."1093 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.1094 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.1095" 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."1096 These numbers agree well with ~0.02% measured directly from our dataset., These numbers agree well with $\sim$ measured directly from our dataset.1097" We conclude that bright flares are rare, but not singular events."," We conclude that bright flares are rare, but not singular events."1098 No separate process is required for their explanation., No separate process is required for their explanation.1099" 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."1100" 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."1101 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.1102" 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."1103" 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."1104" 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} $.1105 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.1106" 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.1107 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.1108" 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."1109" 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."1110 All three bands show the same behavior., All three bands show the same behavior.1111" 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."1112 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.1113" 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."1114" 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}. ."1115 A clumpy wind is also supported by observations., A clumpy wind is also supported by observations.1116" Based on data,"," Based on data,"1117(see. e.g.. De Marco et 2005 and Mathieu Geller 2009).,"(see, e.g., De Marco et 2005 and Mathieu Geller 2009)."1118 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."1119 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.1120 The two binaries. ΙΟΤΕ and NOLS]. may be examples.," The two binaries, KOI-74 and KOI-81, may be examples."1121 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."1122 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.1123 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."1124 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.1125 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."1126 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."1127 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.1128 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.1129 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."1130 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)."1131 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.1132 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."1133 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.1134 Massive cwarls have radii comparable to the radius of the Earth., Massive s have radii comparable to the radius of the Earth.1135 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.1136 The rracdius increases with decreasing mass (see., The radius increases with decreasing mass (see.1137 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."1138 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.1139 We therefore, We therefore114055C [lux is stronger than the IC/CMD flux in 5-avs.,SSC flux is stronger than the IC/CMB flux in $\gamma$ -rays.1141 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."1142" 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 $p_2 = 2.5$."1143 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}."1144 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."1145" 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."1146 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."1147 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}."1148.. This difference of the magnetic field strength can be explained as follows., This difference of the magnetic field strength can be explained as follows.1149 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.1150 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.1151 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."1152 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."1153" 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 $B_{\rm now} = 85\mu \rm{G}$ ."1154velocity as the ejecta expand.,velocity as the ejecta expand.1155 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."1156 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.1157 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)."1158 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.1159 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)."1160 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 3.1161 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."1162 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."1163 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."1164 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.1165 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."1166 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."1167" 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."1168 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."1169 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."1170 We address the possible physical origin of this leature in 85., We address the possible physical origin of this feature in 5.1171appears 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).1172 ]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)."1173 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+226.1174" 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."1175 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.)"1176 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.1177 2004 who reach a similar conclusion)., 2004 who reach a similar conclusion).1178 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.1179 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."1180 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."1181 2002)., 2002).1182 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."1183 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."1184 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.1185 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.1186 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."1187 However the data of Nagata et al. (, However the data of Nagata et al. (11882003). 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.1189 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."1190 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.)."1191 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 $-$ 318.1192 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.1193 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."1194 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."1195 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.1196 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.1197 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.1198 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."1199 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.1200 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.

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