diff --git "a/batch_s000031.csv" "b/batch_s000031.csv"
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+source,target
+" For our study, we take 500 yr."," For our study, we take $t_{sed}=500$ yr."
+ We assume the initial shock velocity uo to be 10? cm/s. This gives CR escape times from the SNRs in the range of 10°) yr and the corresponding escape radii as Rese pc., We assume the initial shock velocity $u_0$ to be $10^9$ cm/s. This gives CR escape times from the SNRs in the range of $t_{esc}=(500-10^5)$ yr and the corresponding escape radii as $R_{esc}\sim(5-100)$ pc.
+" Finally, we treat the escape parameter o and the injection efficiency of the protons (helium) hereafter denoted by ερ(μα) as free parameters."," Finally, we treat the escape parameter $\alpha$ and the injection efficiency of the protons (helium) hereafter denoted by $\epsilon_{p(he)}$ as free parameters."
+" For our calculations, we will assume that all the parameters mentioned above are same for all the SNRs."," For our calculations, we will assume that all the parameters mentioned above are same for all the SNRs."
+" Because of lack of precise informations on the values of a and we perform calculations for several of their randomly chosen €p(ne),combinations."," Because of lack of precise informations on the values of $\alpha$ and $\epsilon_{p(he)}$, we perform calculations for several of their randomly chosen combinations."
+ We choose the escape parameter in the range of a=(1— 3)., We choose the escape parameter in the range of $\alpha=(1-3)$ .
+ This range approximately covers the a values given in some available literatures., This range approximately covers the $\alpha$ values given in some available literatures.
+" Studies based on non-linear diffusive shock acceleration theories which takes into account the modification of the shock structure by the CRs give α~0.8 (e.g., Ptuskin & Zirakashvili 2005)."," Studies based on non-linear diffusive shock acceleration theories which takes into account the modification of the shock structure by the CRs give $\alpha\sim 0.8$ (e.g., Ptuskin $\&$ Zirakashvili 2005)."
+ Blasi & Amato 2011 adopted a~3.2 in their study of the effect of random nature of SNRs on the CR spectrum., Blasi $\&$ Amato 2011 adopted $\alpha\sim 3.2$ in their study of the effect of random nature of SNRs on the CR spectrum.
+ Investigations of y-ray emissions from molecular clouds interacting with nearby SNRs adopt values in the range of a=(2.4—2.6) (Gabici et al., Investigations of $\gamma$ -ray emissions from molecular clouds interacting with nearby SNRs adopt values in the range of $\alpha=(2.4-2.6)$ (Gabici et al.
+" 2009, Ohira et al."," 2009, Ohira et al."
+ 2011)., 2011).
+" We consider the CR injection efficiencies in the range of ερ=(5—25)% for protons and ej.=(1—5)96 for helium, where the values are in units of 10?! ergs."," We consider the CR injection efficiencies in the range of $\epsilon_p=(5-25)\%$ for protons and $\epsilon_{he}=(1-5)\%$ for helium, where the values are in units of $10^{51}$ ergs."
+ The averaged proton to helium injection ratio of 5 which we consider here is less than the observed proton to helium flux ratio of ~(20—13) in the energy range of ~(20—200) GeV/n (Yoon et al., The averaged proton to helium injection ratio of 5 which we consider here is less than the observed proton to helium flux ratio of $\sim (20-13)$ in the energy range of $\sim (20-200)$ GeV/n (Yoon et al.
+ 2011)., 2011).
+" But, our wide range of efficiencies for both the species well cover the observed flux ratios."," But, our wide range of efficiencies for both the species well cover the observed flux ratios."
+ It should be understood that the observed flux ratios may not necessarily represent the injection efficiency ratios from the source., It should be understood that the observed flux ratios may not necessarily represent the injection efficiency ratios from the source.
+ Effects during the propagation inthe Galaxy suchas due tospallation(which are different for different nuclear species depending on their interaction cross-sections) may change the composition ratios produced by thesource., Effects during the propagation inthe Galaxy suchas due tospallation(which are different for different nuclear species depending on their interaction cross-sections) may change the composition ratios produced by thesource.
+" In addition, propagation of CRs is charged dependent."," In addition, propagation of CRs is charged dependent."
+ Those which undergo faster diffusion will escape more, Those which undergo faster diffusion will escape more
+eedee.,edge.
+ Consequently. this unique energv distribution of XUV lines relative to the energy locations of these photoionization edges indicates that the NUY emission should produce different. effects on the phosphorous aud sulfur ionization equilibria.," Consequently, this unique energy distribution of XUV lines relative to the energy locations of these photoionization edges indicates that the XUV emission should produce different effects on the phosphorous and sulfur ionization equilibria."
+ This is illustrated in Figure 3 which shows that the pphotoionizalion rate is 22 10 times larger than the rrate al (he temperature of maximum XUV emission (see Fig., This is illustrated in Figure 3 which shows that the photoionization rate is $\approx$ 10 times larger than the rate at the temperature of maximum XUV emission (see Fig.
+ 1)., 1).
+ Figure 3 includes the rrate Lor comparison. and also shows the expected photoionization rates using only the X-ray energv band (> 124 eV).," Figure 3 includes the rate for comparison, and also shows the expected photoionization rates using only the X-ray energy band $\ge$ 124 eV)."
+ As evident from the displacement of these curves. the neglect of XUV radiation leads to underestimates of (these rates which are substantial al temperatures < 2 MK.," As evident from the displacement of these curves, the neglect of XUV radiation leads to underestimates of these rates which are substantial at temperatures $<$ 2 MK."
+ This implies that alone. the radiation [rom the X-ray energy band is not expected to have a significant impact on the ionization equilibrium of phosphorous. as was recently demonstrated by Ixrticka Ixubat (2009) in their study of the Anger effect on the fractional abundance ofv.," This implies that alone, the radiation from the X-ray energy band is not expected to have a significant impact on the ionization equilibrium of phosphorous, as was recently demonstrated by Krticka Kubat (2009) in their study of the Auger effect on the fractional abundance of."
+. Our graphical avguments illustrate that (he XUV line emission lor temperatures between 0.5 (o 2 MIX is expected to have a major impact on the ionization structure of phosphorous but a relatively minor effect on sulfur., Our graphical arguments illustrate that the XUV line emission for temperatures between 0.5 to 2 MK is expected to have a major impact on the ionization structure of phosphorous but a relatively minor effect on sulfur.
+ This implies that the sulfur surrogate argument needs to be re-examined (see 83)., This implies that the sulfur surrogate argument needs to be re-examined (see 3).
+ In fact. since the pphotoionization edge is almost identical to the eedge (see Fig.," In fact, since the photoionization edge is almost identical to the edge (see Fig."
+ 2). lis a more appropriate surrogate forv.," 2), is a more appropriate surrogate for."
+. The main difference is that the pphotoionizatlion rate is larger than the rrate (see Fig., The main difference is that the photoionization rate is larger than the rate (see Fig.
+ 3) due to the differences in their cross sections., 3) due to the differences in their cross sections.
+ The stellar wind ionization equilibria of phosphorous ancl sulfur are calculated in a straightforward way (o determine the level of NUVX-ray radiation required to affect the fractional ionization abundances. ie. aandv).," The stellar wind ionization equilibria of phosphorous and sulfur are calculated in a straightforward way to determine the level of XUV+X-ray radiation required to affect the fractional ionization abundances, i.e., and."
+. We consider a stellar effective temperature 5;)) range from 27500 to 45000 K which covers the O and early D spectral range where hhas been used (o studyM., We consider a stellar effective temperature ) range from 27500 to 45000 K which covers the O and early B spectral range where has been used to study.
+. The ionization equilibrium is determined by adopting an ionization/recombination rate balance approach similar to the one used in FASTWIND (Puls et al., The ionization equilibrium is determined by adopting an ionization/recombination rate balance approach similar to the one used in FASTWIND (Puls et al.
+ 2005)., 2005).
+ The main dillerences are: 1) we use the photoionization cross sections of Verner Yakovlev (1995): 2) a wind diffuse field as prescribed by Drew (1989). and: 3) we assume a radial power law dependent wind temperature (Z4) which is adjusted to produce a phosphorous wind ionization structure similar (o that of Puls," The main differences are: 1) we use the photoionization cross sections of Verner Yakovlev (1995); 2) a wind diffuse field as prescribed by Drew (1989), and; 3) we assume a radial power law dependent wind temperature $T_W$ ) which is adjusted to produce a phosphorous wind ionization structure similar to that of Puls"
+ol N127245. 5101000. and S103759 determined [rom the Lya emission line are 4.42+0.01. 4.612:0.01. and 4.8240.01.respectivelv.,"of N127245, S101900, and S103759 determined from the $\alpha$ emission line are $4.42\pm0.01$, $4.61 \pm 0.01$, and $4.82 \pm 0.01$,."
+. The other identified LBG. N141368 that is chosen from the faint LBG sample (z=25.32 mag). shows a single emission line al7478A.," The other identified LBG, N141368 that is chosen from the faint LBG sample $z' =25.32$ mag), shows a single emission line at."
+. If the emission line is wjJA5007 or 112. its redshift is 0.49 or 0.54. respectively.," If the emission line is $\lambda$ 5007 or $\beta$, its redshift is 0.49 or 0.54, respectively."
+ In this case. a strong Io. emission line is expected to come to orÀ.. respectively.," In this case, a strong $\alpha$ emission line is expected to come to or, respectively."
+ However. no significant emission is seen at this wavelength. though the signal-to-noise ratio (S/N) in such à red region is very low.," However, no significant emission is seen at this wavelength, though the signal-to-noise ratio (S/N) in such a red region is very low."
+ If the emission line is A5007. the 14]A4959 should be seen al with a 1/3 flux. which is not seen.," If the emission line is $\lambda$ 5007, the $\lambda$ 4959 should be seen at with a 1/3 flux, which is not seen."
+ Another possibility is an identification of u|A3727., Another possibility is an identification of $\lambda$ 3727.
+ IE this is the case. an IL2 emission line should come toO754A.," If this is the case, an $\beta$ emission line should come to."
+. Again no significant emission line is seen., Again no significant emission line is seen.
+ Since the continuum feature longward of the emission line is slightly seen in the wavelength regions where the sky emission is weak and there seems to be a break around αἱ74604... we identilied this emission line as Ίνα.," Since the continuum feature longward of the emission line is slightly seen in the wavelength regions where the sky emission is weak and there seems to be a break around at, we identified this emission line as $\alpha$."
+ The redshift determined [rom the Lya emission is 5.152:0.01., The redshift determined from the $\alpha$ emission is $5.15 \pm 0.01$.
+ The properties of the LBGs identilied are summarized in Table BiQe., The properties of the LBGs identified are summarized in Table \ref{lbg2}.
+ Among 18 remaining bright LBC candidates. we concluded one object. (N111905) is a foregrouncl contamination as described in section 3.2..," Among 18 remaining bright LBG candidates, we concluded one object (N111905) is a foreground contamination as described in section \ref{foregroundobjects}."
+ The remaining objects were not identified because of low S/N in their spectra or no spectral feature in the observed wavelength coverage., The remaining objects were not identified because of low S/N in their spectra or no spectral feature in the observed wavelength coverage.
+ Combining with our previous results bv Ancoetal.(2004.2007).. the total number of ourspectroscopic LBG sample is 16.," Combining with our previous results by \citet{and04,and07}, the total number of ourspectroscopic LBG sample is 16."
+ We examined the identification rates of the spectroscopic sample including the results bv Andoetal.(2004.2007).," We examined the identification rates of the spectroscopic sample including the results by \citet{and04,and07}."
+. Targets with the higher average surface brightness tend to be identified elficientlv: in z/<25 mag. among targets with average surface brightness higher (han 20.8 mag 7? the identification rate is ~ GOW... while it is ~LOW among targets with fainter average surface brightness.," Targets with the higher average surface brightness tend to be identified efficiently; in $z'<25$ mag, among targets with average surface brightness higher than 29.8 mag $^{-2}$ the identification rate is $\sim 60$ , while it is $\sim 10$ among targets with fainter average surface brightness."
+ Among targets with z/<25 mag. highly concentrated objects tend to be identified: ~50% are identified among objects with C<2.5. while only ~20*4. for those with C>2.5.," Among targets with $z'<25$ mag, highly concentrated objects tend to be identified; $\sim 50$ are identified among objects with $C<2.5$, while only $\sim 20$ for those with $C>2.5$."
+ These trends. however. can not be seen among targets with />25 mae: identification rate is higher in more extended objects. though the subsample size is smaller.," These trends, however, can not be seen among targets with $z'>25$ mag; identification rate is higher in more extended objects, though the subsample size is smaller."
+ It should be worth mentioning that the identified LDGs with z>25 mag show strong Lya emission. while most of the identified LBGS with z'«25 mag show no or verv weak Lya emission.," It should be worth mentioning that the identified LBGs with $z'>25$ mag show strong $\alpha$ emission, while most of the identified LBGS with $z'<25$ mag show no or very weak $\alpha$ emission."
+ The trend that luminous LDGs do not show strong Lya emission prevents us from achieving a high identification rate even in briehtLDG sample., The trend that luminous LBGs do not show strong $\alpha$ emission prevents us from achieving a high identification rate even in brightLBG sample.
+ The trend and its physical cause will be discussed in 54.3., The trend and its physical cause will be discussed in $\S 4.3$ .
+effects.,effects.
+ For example. oscillatory trend. will a period of 1 νου is probably due to orbital motion of the Earth and SOMO satellite around the Sun.," For example, oscillatory trend with a period of 1 year is probably due to orbital motion of the Earth and SOHO satellite around the Sun."
+ Similarly. (he sharp variation seen between 1998.4 and 1999.4. (depending on degree () is most likely a result. of changes in insirumental characteristics during the recovery of (he SOIIO satellite.," Similarly, the sharp variation seen between 1998.4 and 1999.4, (depending on degree $\ell$ ) is most likely a result of changes in instrumental characteristics during the recovery of the SOHO satellite."
+ These two effects can account for all claims of radius variation made earlier., These two effects can account for all claims of radius variation made earlier.
+ All these instrumental errors need to be eliminated before any claim can be made about the cause of frequency. variation., All these instrumental errors need to be eliminated before any claim can be made about the cause of frequency variation.
+ From the results presented above it is clear that after eliminating these instrumental effects there is no significant. variation in the solar radius as determined by [node frequencies., From the results presented above it is clear that after eliminating these instrumental effects there is no significant variation in the solar radius as determined by f-mode frequencies.
+ Similar conclusion was obtained by Antia et al. (, Similar conclusion was obtained by Antia et al. (
+2001) using a more detailed analvsis of both GONG and MIDI data.,2001) using a more detailed analysis of both GONG and MDI data.
+ The svstematie error between MDI data sets before ancl after recovery also manilests in other studies (Antia 2002: Basu Antia 2002: Antia. Chitre Thompson 2003).," The systematic error between MDI data sets before and after recovery also manifests in other studies (Antia 2002; Basu Antia 2002; Antia, Chitre Thompson 2003)."
+ In particular. it is found. that this svstematie error is mostly confined to modes with (>120. which is consistent with the results in this work.," In particular, it is found that this systematic error is mostly confined to modes with $\ell>120$, which is consistent with the results in this work."
+ When these modes are neglected no radius variation is [ound. while if these are included then we find varving amount of radius variation around 1999.," When these modes are neglected no radius variation is found, while if these are included then we find varying amount of radius variation around 1999."
+ If the radius variation is real. ib cannot depend on (," If the radius variation is real, it cannot depend on $\ell$."
+" Lt is quite likely (hat svslenmatic errors are present in all MDI data sets. but their magnitude has changed during recovery,"," It is quite likely that systematic errors are present in all MDI data sets, but their magnitude has changed during recovery."
+ If we assume that (he inferred radius variation between 1998.4 ancl 1999.4 is of instrumental origin. (hen we can put some limits on radius variation.," If we assume that the inferred radius variation between 1998.4 and 1999.4 is of instrumental origin, then we can put some limits on radius variation."
+ From Table 1 it can be seen Chat the maxinmun variation between (he 3 points before the data gap is 1.3 kim. while that in the 4 points alter the gap is 1.1 km.," From Table 1 it can be seen that the maximum variation between the 3 points before the data gap is 1.3 km, while that in the 4 points after the gap is 1.1 km."
+ Considering an error of about 0.6 km in each data point. (his variation is consistent with no radius variation.," Considering an error of about 0.6 km in each data point, this variation is consistent with no radius variation."
+ This would suggest an upper limit comparable to error bars in each point on radius variation during half of the solar cycle., This would suggest an upper limit comparable to error bars in each point on radius variation during half of the solar cycle.
+ Similar conclusion can be obtained from Fig., Similar conclusion can be obtained from Fig.
+ 4. which shows the results using only 6«120 mocles which are not expected to be affected bv the svstematic error in MDI data.," 4, which shows the results using only $\ell<120$ modes which are not expected to be affected by the systematic error in MDI data."
+ From these results we can put a conservative upper limit of 2 km on radius variations during (he last 6 years., From these results we can put a conservative upper limit of 2 km on radius variations during the last 6 years.
+" This would vield AR/R«3x10"" and Wo<0.003 as the ratio of radius to luminosity variation.", This would yield $\Delta R/R < 3\times10^{-6}$ and $W<0.003$ as the ratio of radius to luminosity variation.
+ Such a small value should favor models involving changes in (he outer lavers (o explain the observed Iuminosity. variations (Gough 1981: Dapppen 1983: Dalinforth. Gough Merrvlield 1996: Gough 2001).," Such a small value should favor models involving changes in the outer layers to explain the observed luminosity variations (Gough 1981; Däpppen 1983; Balmforth, Gough Merryfield 1996; Gough 2001)."
+ Of course. the value of WW in these models is determined by radius variation at the photosphere. while [-modes are sensitive to variations at cepths of about 1.12 Mm.," Of course, the value of $W$ in these models is determined by radius variation at the photosphere, while f-modes are sensitive to variations at depths of about 1–12 Mm."
+ But if the photospheric radius variations are much larger (han those inferred by the modes. then the cause is most likely to be near the surface.," But if the photospheric radius variations are much larger than those inferred by the f-modes, then the cause is most likely to be near the surface."
+ Emilio et al. (, Emilio et al. (
+2000) find a much larger increase in photospheric radius by about. 15z2 km during the solar cycle.,2000) find a much larger increase in photospheric radius by about $15\pm2$ km during the solar cycle.
+ These direct measurements [rom ALDI are also affected by a number of svstematic errors ancl as (hey have pointed out this value should be regarded as an upper Init to radius variation., These direct measurements from MDI are also affected by a number of systematic errors and as they have pointed out this value should be regarded as an upper limit to radius variation.
+"Amalgamating our results with previous CSS data, our key findings are as follows:","Amalgamating our results with previous CSS data, our key findings are as follows:"
+2007).,.
+. Choosing au appropriate kinetic temperature estimate was difficult for this region because of the inclusion of several compact sources witlin our beam., Choosing an appropriate kinetic temperature estimate was difficult for this region because of the inclusion of several compact sources within our beam.
+ We have asstuned arange of 10.100 IX based ou the estimates of Broganctal.(2007) and Codellaetal.(2005) for the ambient cloud velocity of ~ 10.5st., We have assumed a range of 40–100 K based on the estimates of \citet{Bro07} and \citet{Cod05} for the ambient cloud velocity of $\sim$ 10.5.
+.. Codellaotal.(2005) also find densities between in the region using SO and SiO. The brightest of three compact sources discovered. by Martin(1973).. which contains a zero-age lai sequence star of spectral type earlier than O7.5 (Lugoctal.2001) clubedded in a molecular cloud with an inner shell structure πάτοςας the star (Pratapetal.1997).," \citet{Cod05} also find densities between in the region using SO and SiO. The brightest of three compact sources discovered by \citet{Mar73}, which contains a zero-age main sequence star of spectral type earlier than O7.5 \citep{Lug04} embedded in a molecular cloud with an inner shell structure surrounding the star \citep{Pra97}."
+.. A systemic velocity of around -57 was observed in several transitions of between 365170 GIIz by vanderTaketal.(2000) with an absorption feature near -60 having been detected in (Wilsonetal. 1983).. both of which are found in our spectra.," A systemic velocity of around -57 was observed in several transitions of between 365–470 GHz by \citet{Tak00} with an absorption feature near -60 having been detected in \citep{Wil83}, both of which are found in our spectra."
+ The region is also notable for the rare occurrence of 6 cni (1490-1413) inasers (Botsetal.1981:Hoffinan 2003).," The region is also notable for the rare occurrence of 6 cm ) masers \citep{Rot81,Hof03}."
+. The absorption component we observe was also detected by Woffinanetal.(2003). in the 2c (244-249)) transition ofC'O., The absorption component we observe was also detected by \citet{Hof03} in the 2cm ) transition of.
+. A temperature of — 220 EK was found using by Mauersbergeretal.(1988)... while a cooler temperature of 176 IX was found by AGtchelletal.(1990). using 15CC€0 and a warmer one of 215 K by Qiuetal.(2011) ming CILCCN.," A temperature of $\sim$ 220 K was found using by \citet{Mauer88}, while a cooler temperature of 176 K was found by \citet{Mit90} using CO and a warmer one of 245 K by \citet{Qiu11} using CN."
+ The Mitchell study. also iudicates a cold gas component of 25 I from which the absorption conrponeut in our J=3 spectrun likely arises;, The Mitchell study also indicates a cold gas component of 25 K from which the absorption component in our $J=3$ spectrum likely arises.
+ From 15CC0 observaions. Qiuefal.(2011) estimate the deusitv to be ~ for a region about half the size of our beam. while Mitchelletal.(1990) indicate that the deusity is >σι," From CO observations, \citet{Qiu11} estimate the density to be $\sim$ for a region about half the size of our beam, while \citet{Mit90} indicate that the density is $>$."
+", Hoffiiauetal.(2003) note that o excite the masers in the region. lower densities in the rause of are required."," \citet{Hof03}
+ note that to excite the masers in the region, lower densities in the range of are required."
+ The masers arise feo a region directly iu front IRS 1 rot coincident with the hot core region examined by the xeviouslv mentioned authors., The masers arise from a region directly in front IRS 1 not coincident with the hot core region examined by the previously mentioned authors.
+ Since our observations are of ligher excitation ransitions than the masers in the Woffinan study. it is likely that we areο sampling material from voth the hot core and maser regions. and thus the intermediate density we find is justified.," Since our observations are of higher excitation transitions than the masers in the Hoffman study, it is likely that we are sampling material from both the hot core and maser regions, and thus the intermediate density we find is justified."
+ However. this rotion also calls iuto question the validity of our kinetic eniperature assunrption. so it should be noted that cluperatures <<120 [I would result in densities > iu the LVG approximation.," However, this notion also calls into question the validity of our kinetic temperature assumption, so it should be noted that temperatures $<120$ K would result in densities $>$ in the LVG approximation."
+ That said. due to he excitation requirements of the aud ransitious compared to that of the J=1 A--doublet observe bv Uoffinanetal. (2003).. if is more Likely hat our measurements are biased toward the higher eniperatures of the hot core.," That said, due to the excitation requirements of the and transitions compared to that of the $J=1$ -doublet observed by \citet{Hof03}, , it is more likely that our measurements are biased toward the higher temperatures of the hot core."
+ Fortunately. the LVG models are relatively independent of for temperatures 2100 TS (see mp) ).," Fortunately, the LVG models are relatively independent of for temperatures $\gtrsim100$ K (see \\ref{temp}) )."
+ Deeply eiibedded cold IR source ~ 50” south of IRS 1 that is associated with a large reflection nebula (Werneretal. 1979).., Deeply embedded cold IR source $\sim$ $\arcsec$ south of IRS 1 that is associated with a large reflection nebula \citep{Wer79}. .
+ The radial velocity of the IRS 9 cloud core from HCCN. CS. aud data is about -57 (Sandelletal.2005:vanderTak2000).. consistent with our observations.," The radial velocity of the IRS 9 cloud core from CN, CS, and data is about -57 \citep{San05,Tak00}, consistent with our observations."
+ Sinele-cish JCAIT and2oy.. 218 CIIz) spectra sugecst a easο temperature of = 60 I& (Sandelletal.2005).. while Ποσο.UPCCO!.. and PCCO observations suggest a lower temperature of ~ 20 I& (ITasegawa&Mitchell1995:Alitehelletal.1990).," Single-dish JCMT and, 218 GHz) spectra suggest a gas temperature of $\geq$ 60 K \citep{San05}, while , and CO observations suggest a lower temperature of $\sim$ 30 K \citep{HM95,Mit90}."
+. \litchelletal.(1990). also detect the presence of warn gas (180 IK) aud sugeest that the deusity of this material is >cm?., \citet{Mit90} also detect the presence of warm gas (180 K) and suggest that the density of this material is $>$.
+", Qur study is concerned primarily with the application of a previously unused deusitomietrv technique and. as such. well-studied objects were chosen to assess the iuiethod's performance."," Our study is concerned primarily with the application of a previously unused densitometry technique and, as such, well-studied objects were chosen to assess the method's performance."
+ This discussion is focused ou limitations to the technique that require claboration., This discussion is focused on limitations to the technique that require elaboration.
+ These effects παν also help explain the problem of the anomalously low transition ratios [ΤουBide ο that precluded LVG modcling in a few sources (soe Table 1))., These effects may also help explain the problem of the anomalously low transition ratios $\int{}T_{mb}(\jthree)d\nu$ $\int{}T_{mb}(\jfour)d\nu$ ] that precluded LVG modeling in a few sources (see Table \ref{tab:LVGDetections}) ).
+" The LVC model for the 210. transitions is solnewhat dependent on kinetic temperature ΓΕΝ),", The LVG model for the and transitions is somewhat dependent on kinetic temperature ).
+ As previously mentioned in rofLVG.. this ameant that ai careful selection οἳ appropriate kinetic teniperatures from previous studies was necessary.," As previously mentioned in \\ref{LVG}, this meant that a careful selection of appropriate kinetic temperatures from previous studies was necessary."
+ Wherever possible. temperatures derived from measurements were used to eusiure coupling to the eas traced by our observations.," Wherever possible, temperatures derived from measurements were used to ensure coupling to the gas traced by our observations."
+ Estimates taken fron analyses of other dense eas tracers such as or CIT4CCN are otherwise preferable., Estimates taken from analyses of other dense gas tracers such as or CN are otherwise preferable.
+ That said. anv time a temperature derived from measurements of disparate molecules (or even measurements of significantly different cxcitation requiremeutz) is adopted. the question of whether or not this estimate cau be associated with the gas sampled by our beam iust be radsed.," That said, any time a temperature derived from measurements of disparate molecules (or even measurements of significantly different excitation requirements) is adopted, the question of whether or not this estimate can be associated with the gas sampled by our beam must be raised."
+ In at least oue case (NCC) 7538 IRS 1. refu7538i1l)) there appears to be a distinct possibility that the temperature we adopted may uot be wholly appropriate for the material traced by our observatious.," In at least one case (NGC 7538 IRS 1, \\ref{n7538i1}) ) there appears to be a distinct possibility that the temperature we adopted may not be wholly appropriate for the material traced by our observations."
+ Iu Table L. we elected to include the spatial deusities derived for a range of kinetic temperatures based on the error estimates in our assumed values.," In Table \ref{tab:LVGDetections}, we elected to include the spatial densities derived for a range of kinetic temperatures based on the error estimates in our assumed values."
+ This is because the temperature dependence of our LVC models is not linear aud itselfdepends on the raugeοof plivsical paralucters being studied., This is because the temperature dependence of our LVG models is not linear and itselfdepends on the rangeof physical parameters being studied.
+ Table 7 provides a stuuimation of this cepeudence by dividing the rauge of kinetic, Table \ref{tab:Tdepend} provides a summation of this dependence by dividing the range of kinetic
+winds of massive stars on stellar laminosity. offering a polentially more accurate distance indicator.,"winds of massive stars on stellar luminosity, offering a potentially more accurate distance indicator."
+ While hot O stars provide so [ur the best calibration of the WLR (Pulsetal.1996.. Pulsetal. 2002)). it is the visually brightest late B and early A supergiants CM:27—9) which offer (he largest potential as extragalactic standard candles (Ixudiritzki1993.. 1999)).," While hot O stars provide so far the best calibration of the WLR \citealt{puls96}, , \citealt{puls02}) ), it is the visually brightest late B and early A supergiants $M_V\simeq-9$ ) which offer the largest potential as extragalactic standard candles \citealt{kud98}, \citealt{kud99}) )."
+ This can now be investigated [or the nearby galaxies (D<10 Mpc) with multiobject spectroscopy al 82m class telescopes. as the exploratory work of Bresolinet has shown.," This can now be investigated for the nearby galaxies $D<10$ Mpc) with multiobject spectroscopy at 8-m class telescopes, as the exploratory work of \citet{bresolin01,
+bresolin02} has shown."
+ Quantitative spectroscopy of individual BA supergiants leads to the determination of evavilies. temperatures. metallicities and stellar wind parameters (based on (he wind emission in Hoa). which are then combined to provide distances.," Quantitative spectroscopy of individual BA supergiants leads to the determination of gravities, temperatures, metallicities and stellar wind parameters (based on the wind emission in $\alpha$ ), which are then combined to provide distances."
+ In this Letter we suggest a novel method. based on the absorption strengths of the hieher Balmer lines formed in the photosphere.," In this Letter we suggest a novel method, based on the absorption strengths of the higher Balmer lines formed in the photosphere."
+ The concept is not entirely new. since a relationship between the equivalent width of Ils and Αν for Galactic and. Magellanic Clouds BA supergiants. together with a temperature dependence. has already been discussed bv several authors (Petrie1965... Crampton 1979.. Tully&Wolff 1934... Hilletal.1986).," The concept is not entirely new, since a relationship between the equivalent width of $\gamma$ and $M_V$ for Galactic and Magellanic Clouds BA supergiants, together with a temperature dependence, has already been discussed by several authors \citealt{petrie65}, \citealt{crampton79}, \citealt{tully84}, \citealt{hill86}) )."
+ The recent improvements in the modelling of A supergiant atmospheres in NLTE and the development of new diagnostic tools (Venn1995a.b.1999:Vennetal.2001.. Przvbillaetal.POOL:Przvbilla&Butler2001: 2002)) allows us to determine stellar parameters wilh unprecedented accuracy aud reliability.," The recent improvements in the modelling of A supergiant atmospheres in NLTE and the development of new diagnostic tools \citealt{venn95a,venn95b,venn99,venn01}, \citealt{przybilla01a,przybilla01b,przybilla02}) ) allows us to determine stellar parameters with unprecedented accuracy and reliability."
+ Based on these achievements. we discuss here a promising application on a set of hieh-to-intermediate resolution spectra obtained for a sample of A supergiants in the Milky Wav. the Magellanic Clouds. NGC 300 and NGC 3621. and [few additional objects in a handful of nearby galaxies.," Based on these achievements, we discuss here a promising application on a set of high-to-intermediate resolution spectra obtained for a sample of A supergiants in the Milky Way, the Magellanic Clouds, NGC 300 and NGC 3621, and few additional objects in a handful of nearby galaxies."
+ The theoretical concept is explaimed in 82. and we present the observational tests in 83 and 84.," The theoretical concept is explained in 2, and we present the observational tests in 3 and 4."
+ Future work on supereiants in nearby galaxies is briefly. summarized in 35., Future work on supergiants in nearby galaxies is briefly summarized in 5.
+ Massive stars during their evolution towards the Bed Supergiant stage pass through ihe phase of late B and early A supergiants quickly aud with roughly constant mass and luminosity (Mevnet&Maeder2000:Meynetetal.1994:HegerLanger2000)).," Massive stars during their evolution towards the Red Supergiant stage pass through the phase of late B and early A supergiants quickly and with roughly constant mass and luminosity \citealt{meynet00,meynet94,heger00}))."
+" This means that in (his phase the stellar gravity g and effective temperature are coupled through the condition ο =const, We call Τη the ""Hux-weighted eravitv.", This means that in this phase the stellar gravity $g$ and effective temperature are coupled through the condition $g$ $= const.$ We call $g$ the “flux-weighted gravity”.
+" Assuming that mass and liminositv follow the usual relation Lx MP"". (a~ 3). we derive arelationship between absolute bolometric magnitude Mj aud the fIux-weighted gravity of the form"," Assuming that mass and luminosity follow the usual relation $L \propto M^{\alpha}$ , $\alpha \sim 3$ ), we derive arelationship between absolute bolometric magnitude $M_{bol}$ and the flux-weighted gravity of the form"
+discovery observations.,discovery observations.
+ New observational data yielding statistically sound detections are required before PO4s claim that #11916 is the most distant galaxy yet discovered may be resurrected., New observational data yielding statistically sound detections are required before P04's claim that 1916 is the most distant galaxy yet discovered may be resurrected.
+ We consider this unlikely. but we hope to be surprised by Pelló et al.," We consider this unlikely, but we hope to be surprised by Pelló et al."
+"'s forthcoming observations,",'s forthcoming observations.
+ We first discuss possible reasons for the difference between our non-detection of #11916 and PO4s —3-4o near-infrared detections., We first discuss possible reasons for the difference between our non-detection of 1916 and P04's ${\sim}3{-}4{\sigma}$ near-infrared detections.
+ We single out two data reduction steps refisaac for discussion — the efficiency of bad pixel rejection and the conservation) of noise properties., We single out two data reduction steps \\ref{isaac}) ) for discussion – the efficiency of bad pixel rejection and the conservation of noise properties.
+ The efficiency with which bad pixels are identified can affect estimates of the signal coming from faint sources — 1f some bad pixels are not identified. they could enhance the flux detected by SExtractor.," The efficiency with which bad pixels are identified can affect estimates of the signal coming from faint sources – if some bad pixels are not identified, they could enhance the flux detected by SExtractor."
+ In refisaac we used two different methods to identify bad pixels. finding a small but potentially important difference between the two methods.," In \\ref{isaac} we used two different methods to identify bad pixels, finding a small but potentially important difference between the two methods."
+ To assess the impact of reduced efficiency of bad pixel identification we made à mask image that contained the 12.373 bad pixels identified only in the bad pixel mask generated from the sky-flats.," To assess the impact of reduced efficiency of bad pixel identification we made a mask image that contained the 12,373 bad pixels identified only in the bad pixel mask generated from the sky-flats."
+ We then made one copy of the mask per science frame and integer-pixel shifted them to match the observed dither pattern., We then made one copy of the mask per science frame and integer-pixel shifted them to match the observed dither pattern.
+ Finally we took the weighted average of these aligned mask frames and summed the pixel counts ina 1.5” diameter aperture centered on #11916., Finally we took the weighted average of these aligned mask frames and summed the pixel counts in a $1.5''$ diameter aperture centered on 1916.
+ From this we concluded that 7 bad pixels that are only identified in our “sky-flat™ bad pixel maps fall within the final photometric aperture., From this we concluded that 7 bad pixels that are only identified in our “sky-flat” bad pixel maps fall within the final photometric aperture.
+ We estimate that if not identified and excluded from the analysis. these pixels could increase the flux estimates by several tenths of a magnitude.," We estimate that if not identified and excluded from the analysis, these pixels could increase the flux estimates by several tenths of a magnitude."
+ The approach to re-sampling (or not) the individual frames during the data reduction process. especially the alignment of individual frames affects the noise properties of the final stacked frame.," The approach to re-sampling (or not) the individual frames during the data reduction process, especially the alignment of individual frames affects the noise properties of the final stacked frame."
+ If the individual frames are re-sampled. for example by sub-integer pixel aligning them immediately prior to producing the final stacked frame. then the noise m the stacked frame is correlated.," If the individual frames are re-sampled, for example by sub-integer pixel aligning them immediately prior to producing the final stacked frame, then the noise in the stacked frame is correlated."
+ Such pixel-to-pixel correlations are generally absent from integer-pixel aligned data. thus simplifying the noise properties of the final frame.," Such pixel-to-pixel correlations are generally absent from integer-pixel aligned data, thus simplifying the noise properties of the final frame."
+ Neither sub-integer nor integer pixel alignment is intrinsically correct., Neither sub-integer nor integer pixel alignment is intrinsically correct.
+ The relevant issue is correct measurement of the noise in each case — this is critical to assess accurately the statistical significance of sources detected close to the sensitivity limit of the data., The relevant issue is correct measurement of the noise in each case – this is critical to assess accurately the statistical significance of sources detected close to the sensitivity limit of the data.
+ Specifically. if the pixel-to-pixel correlations in sub-integer pixel aligned. data are not included in the error analysis. then the noise is under-estimated and the statistical significance of the detection over-estimated (Casertano et 22000).," Specifically, if the pixel-to-pixel correlations in sub-integer pixel aligned data are not included in the error analysis, then the noise is under-estimated and the statistical significance of the detection over-estimated (Casertano et 2000)."
+ We integer pixel aligned the individual frames in refisaac in order to simplify the error analysis., We integer pixel aligned the individual frames in \\ref{isaac} in order to simplify the error analysis.
+ We now estimate by how much we would have under-estimated the noise if we had sub-integer pixel aligned the individual frames and then ignored the pixel-to-pixel correlations when calculating noise level., We now estimate by how much we would have under-estimated the noise if we had sub-integer pixel aligned the individual frames and then ignored the pixel-to-pixel correlations when calculating noise level.
+ This is achieved by simply sub-integer aligning the individual frames and re-combining them using a weighted average., This is achieved by simply sub-integer aligning the individual frames and re-combining them using a weighted average.
+ lgnoring any resulting pixel-to-pixel correlations. we obtain a 30 threshold of H=25.2. which ts slightly fainter than our threshold of H=25.0.," Ignoring any resulting pixel-to-pixel correlations, we obtain a $3{\sigma}$ threshold of $H{=}25.2$, which is slightly fainter than our threshold of $H{=}25.0$."
+ In summary. it is plausible that the difference between our near-infrared non-detection and PO4’s detection of #11916 using the same raw data can at least in part be explained by the efficiency of bad pixel identification and treatment of correlated noise.," In summary, it is plausible that the difference between our near-infrared non-detection and P04's detection of 1916 using the same raw data can at least in part be explained by the efficiency of bad pixel identification and treatment of correlated noise."
+ We now discuss the implications of our results for future work. focusing on the feasibility of searches for gravitationally magnified stellar systems at 7 using ground-based near-infrared data.," We now discuss the implications of our results for future work, focusing on the feasibility of searches for gravitationally magnified stellar systems at $z{\gs}7$ using ground-based near-infrared data."
+ We begin by noting that. had their photometry been reliable. PO4’s original zz10 interpretation of #11916 was plausible. based solely on the photometric data.," We begin by noting that, had their photometry been reliable, P04's original $z{=}10$ interpretation of 1916 was plausible, based solely on the photometric data."
+ We therefore adopt PO4's optical and near-infrared photometry as being representative of what may be expected from similar future experiments — ooptical non-detections, We therefore adopt P04's optical and near-infrared photometry as being representative of what may be expected from similar future experiments – optical non-detections
+The model of Bernstein mode (BAL model) is he first one to interpret the formation of ZP structure.,The model of Bernstein mode (BM model) is the first one to interpret the formation of ZP structure.
+ The emission mechauisui is the nonlinear coupling between two Bernstein modes. or Bernstein mode and other clectrostatic upper hvbrid waves.," The emission mechanism is the nonlinear coupling between two Bernstein modes, or Bernstein mode and other electrostatic upper hybrid waves."
+ The elecrons with non-equilibriun distribution over velocities perpendicular to the magnetic field are located iu a “μα source. where the plasma is weakly magnetized (fue9 fo) and the magnetic field is uniform.," The electrons with non-equilibrium distribution over velocities perpendicular to the magnetic field are located in a small source, where the plasma is weakly magnetized $f_{pe}\gg f_{ce}$ ), and the magnetic field is uniform."
+" Tjese. electrons excite longitudinal electrostatic waves at frequency of the sum of so-called. Berustei modes frequency sf. aud the upper hybrid frequency ""m.", These electrons excite longitudinal electrostatic waves at frequency of the sum of so-called Bernstein modes frequency $sf_{ce}$ and the upper hybrid frequency $f_{uh}$.
+ The BM excitation occurs in relatively narrow frecποιος baud., The BM excitation occurs in relatively narrow frequency band.
+ The emission frequency (Roscubere. 1972: Chinderi et al. 1973: Zaitsev aud Stepanov. 1983) is: Tere. fj. i the electron plaziua frequeucy.. f the clectrou gvro-frequeney. s is harmonics uber.," The emission frequency (Rosenberg, 1972; Chiuderi et al, 1973; Zaitsev and Stepanov, 1983) is: Here, $f_{pe}$ is the electron plasma frequency, $f_{ce}$ the electron gyro-frequency, $s$ is harmonics number."
+ This model preseuts the frequency separation between the adjaceutzebra stripes just as the electron Af=f (Zhelezuvakov and Zlotuik. 1975). (," This model presents the frequency separation between the adjacentzebra stripes just as the electron gyro-frequency: $\Delta f=f_{ce}$ (Zheleznyakov and Zlotnik, 1975). ("
+2) Ueterogenous models. which proposed that zebra stripes iu a structure are generated frou some extended source regious iu the maenetic iux tube (Ixuijpors. 1975: Fouichey aud Fiiushteiu. 19581: Mollswo. 1983: Ledeuev. Yau. aud Fu. 2001: Chernov et al. 2005: AÁltvutsev et al. 2005).,"2) Heterogenous models, which proposed that zebra stripes in a structure are generated from some extended source regions in the magnetic flux tube (Kuijpers, 1975; Fomichev and Fainshtein, 1981; Mollwo, 1983; Ledenev, Yan, and Fu, 2001; Chernov et al, 2005; Altyntsev et al, 2005)."
+ One of the important heterogenous model is owed ou the plasima waves interact with whistler waves (Chernov. 1996. 2006). called as whistler wave mode LAVW inodel).," One of the important heterogenous model is based on the plasma waves interact with whistler waves (Chernov, 1996, 2006), called as whistler wave model (WW model)."
+ The coupling of plasina wave aud whistler wave can operate in different conditions: when whistlers generate at the normal Doppler evclotron resonance they can escape along the magnetic loop and vicl« fiber bursts: when whistlers generate at the anomalous Doppler cyclotron resonance uuder large augles to the magnetic field they may form standing wave packets in frout of the shock wave. and when the eroup veocity of whistlers is approximated to the shock velocity. a ZP structure with slow oscillating frequency dif will appear.," The coupling of plasma wave and whistler wave can operate in different conditions: when whistlers generate at the normal Doppler cyclotron resonance they can escape along the magnetic loop and yield fiber bursts; when whistlers generate at the anomalous Doppler cyclotron resonance under large angles to the magnetic field they may form standing wave packets in front of the shock wave, and when the group velocity of whistlers is approximated to the shock velocity, a ZP structure with slow oscillating frequency drift will appear."
+ The whistler wave eroup velocity peaks at whistler frequency focO.25fF...," The whistler wave group velocity peaks at whistler frequency $f_{w}\sim
+0.25f_{ce}$."
+" The frequency xparatiou Af between adjacent zebra stripes is about 2 times of whistler frequency: Af~2f£,. aud then we mav obtain: f:24Af."," The frequency separation $\Delta f$ between adjacent zebra stripes is about 2 times of whistler frequency: $\Delta
+f\sim 2f_{w}$, and then we may obtain: $f_{ce}\sim2\Delta f$."
+ The most developed heterogenous model for ZP structure generation is called double plasma resonance model (DPR model). which explain ZP structure in a uaural way (Pearlstein. et al. 1966: Zhelezuvalsov Zlotuik. 1975: Berney Benz. LOTS: Winelee Dulk. 1986: Zlotuilk et al. 2003: Yasnov Warlicky. 20M: Wuzuctsov Tsap. 2007).," The most developed heterogenous model for ZP structure generation is called double plasma resonance model (DPR model), which explain ZP structure in a natural way (Pearlstein, et al, 1966; Zheleznyakov Zlotnik, 1975; Berney Benz, 1978; Winglee Dulk, 1986; Zlotnik et al, 2003; Yasnov Karlicky, 2004; Kuznetsov Tsap, 2007)."
+ This model proposed that euliauced excitation of plasiia waves occurs at some resonance levels where the upper hybrid frequeney coincides with the harmonics of clectron evro-frequeucy 11i the inhomegencous flix tube: The emüssiou frequency is dominated not only bv the electron evro-frequency. but also by plasma frequency.," This model proposed that enhanced excitation of plasma waves occurs at some resonance levels where the upper hybrid frequency coincides with the harmonics of electron gyro-frequency in the inhomogeneous flux tube: The emission frequency is dominated not only by the electron gyro-frequency, but also by plasma frequency."
+" When the emission ecuerates from the coalesceuce of two excited plasina waves. the polarization mnmav be very weak. the cussion frequency is fm2f,~-- 25f.,. and the frequency separation between the adjacent⊳ zebra stripes⋅ is⋅ Af⊽⊐⋅⋅⊓∐↴Way: Ico."," When the emission generates from the coalescence of two excited plasma waves, the polarization may be very weak, the emission frequency is $f\approx
+2f_{pe}\approx 2sf_{ce}$ , and the frequency separation between the adjacent zebra stripes is $\Delta
+f=\frac{2sf_{ce}H_{b}}{|sH_{b}-(s+1)H_{p}|}$."
+" I,=fr(dfp⊽fdr)!Βία)* and FoGPdr)EDndn,di)+= "," Here, $H_{b}=f_{B}(df_{B}/dr)^{-1}=B(dB/dr)^{-1}$ and $H_{p}=f_{pe}(df_{pe}/dr)^{-1}=2n_{e}(dn_{e}/dr)^{-1}=2H_{n}$ ."
+"21H21 aud IL, ave the scale heiehts of maguetie field B aud the plasiia density à», in the source regions. respectively,"," $H_{b}$ and $H_{n}$ are the scale heights of magnetic field $B$ and the plasma density $n_{e}$ in the source regions, respectively."
+" For fie«fj, aud s2» 1. we mayget:"," For $f_{ce}\ll f_{pe}$ and $s\gg 1$ , we mayget:"
+averaged out when summing over the images of a night (see Sect. 4)).,averaged out when summing over the images of a night (see Sect. \ref{section:meanimage}) ).
+ Changes in observational conditions (atmospheric absorption and background flux) are taken into account with a global correction relative to the reference image., Changes in observational conditions (atmospheric absorption and background flux) are taken into account with a global correction relative to the reference image.
+ We assume that a linear correction is sufficient: where corrected ANd draw are the pixel fluxes after and before correction respectively., We assume that a linear correction is sufficient: where $\phi_{\rm corrected}$ and $\phi_{raw}$ are the pixel fluxes after and before correction respectively.
+" The absorption factor a is estimated for each image with a PEIDA procedure, based on the comparison of star fluxes between this image and the reference frame (Ansari, 1994))."," The absorption factor $a$ is estimated for each image with a PEIDA procedure, based on the comparison of star fluxes between this image and the reference frame (Ansari, \nocite{Ansari:1994}) )."
+ A sky background excess is supposed to affect pixel fluxes by an additional term b which differs from one image to another., A sky background excess is supposed to affect pixel fluxes by an additional term $b$ which differs from one image to another.
+" In Fig. 2,,"," In Fig. \ref{fig:fdc},"
+ we plot the absorption factor (top) and the sky background (bottom) estimated for each image with respect to the reference image as a function of time., we plot the absorption factor (top) and the sky background (bottom) estimated for each image with respect to the reference image as a function of time.
+ The absorption can vary by as much as a factor 2 within a single night., The absorption can vary by as much as a factor $2$ within a single night.
+" During full moon periods, the background flux can be up to 20 times higher than during moonless nights, increasing the statistical fluctuations by a factor up to 4.5."," During full moon periods, the background flux can be up to 20 times higher than during moonless nights, increasing the statistical fluctuations by a factor up to 4.5."
+" However, this high level of noise concerns very few images (see Fig. 2)),"," However, this high level of noise concerns very few images (see Fig. \ref{fig:fdc}) ),"
+ and only about of the images more than double their statistical fluctuations., and only about of the images more than double their statistical fluctuations.
+" Despite their large noise, full moon images improve the time sampling, and at the end of the whole treatment, the error bars associated with these points are not significantly larger than those corresponding to new moon periods, except for a few nights."," Despite their large noise, full moon images improve the time sampling, and at the end of the whole treatment, the error bars associated with these points are not significantly larger than those corresponding to new moon periods, except for a few nights."
+ We note the presence of a variable spatial pattern particularly important during full moon periods., We note the presence of a variable spatial pattern particularly important during full moon periods.
+" This residual effect, probably due to reflected light, can be eliminated with a procedure similar to that applied to the AGAPE data, as described by Ansari et al. ("," This residual effect, probably due to reflected light, can be eliminated with a procedure similar to that applied to the AGAPE data, as described by Ansari et al. ("
+1997).,1997).
+ We calculate a median image with a sliding window of 9x pixels on the difference between each image and the reference image., We calculate a median image with a sliding window of $9\times 9$ pixels on the difference between each image and the reference image.
+" It is important to work on the difference in order to eliminate the disturbing contributions of stars, and to get a median that retains only large-scale spatial variations."," It is important to work on the difference in order to eliminate the disturbing contributions of stars, and to get a median that retains only large-scale spatial variations."
+" We then subtract the corresponding median from each image, to filter out large-scale spatial variations."," We then subtract the corresponding median from each image, to filter out large-scale spatial variations."
+" In Fig. 3,,"," In Fig. \ref{fig:CLavap},"
+ we show a light curve before and after this correction., we show a light curve before and after this correction.
+" Above, the pixel light curve presents important systematic effects during full moon periods, effects which have disappeared below, after correcting for these large-scale variations."," Above, the pixel light curve presents important systematic effects during full moon periods, effects which have disappeared below, after correcting for these large-scale variations."
+" After these alignments, we eliminate images whose parameters lie in extreme ranges."," After these alignments, we eliminate images whose parameters lie in extreme ranges."
+ We keep images which have no obvious defects and parameters in the following range:, We keep images which have no obvious defects and parameters in the following range:
+reproduce the properties of off-uuclear ULNs.,reproduce the properties of off-nuclear ULXs.
+ We trace he orbits of each BIT and examine the frequency with which they cross the disk region. which we defiue as three disk scale leugths (7).," We trace the orbits of each BH and examine the frequency with which they cross the disk region, which we define as three disk scale lengths $r_s$ )."
+ The last few Covr of orbital history or the DIIs in our simulated galaxy. 1239 are shown iu Figure 3.., The last few Gyr of orbital history for the BHs in our simulated galaxy h239 are shown in Figure \ref{fig:orbit}.
+" Several of the DIIS pass through the 37, limit. which we denote as a disk passage."," Several of the BHs pass through the $3r_s$ limit, which we denote as a disk passage."
+ The average BIT dis- owsaece rate for all of the simulatious combined is 10.6 oor Civi: however. this rate increases with time (13.3 per Cyr for the last few billion years of galaxy evolution) duc o the increased population of wandering BITs in the halo (since more mergers have occurred).," The average BH disk passage rate for all of the simulations combined is 10.6 per Gyr; however, this rate increases with time (13.3 per Gyr for the last few billion years of galaxy evolution) due to the increased population of wandering BHs in the halo (since more mergers have occurred)."
+ The likehhood of observing such a disk passage depends ou the column density of gas in the disk. the nuuber of passages made. aud the velocity of the BUs.," The likelihood of observing such a disk passage depends on the column density of gas in the disk, the number of passages made, and the velocity of the BHs."
+ We use results from ΤΗ measurements of the surface densities of local disk galaxies (Leroyctal.2008) to estimate an average disk column deusitv of 10 pe7., We use results from HI measurements of the surface densities of local disk galaxies \citep{Leroy08} to estimate an average disk column density of 10 $^{-2}$.
+ The ability of the passing DII to attract a sutiicicut amount of gas to be observed as a ULX depends iost stronely on its velocity. which ranges between 151 - 662 km s and on average is ~ LTO law | for BIS at periceuter.," The ability of the passing BH to attract a sufficient amount of gas to be observed as a ULX depends most strongly on its velocity, which ranges between 154 - 662 km $^{-1}$ , and on average is $\sim$ 470 km $^{-1}$ for BHs at pericenter."
+ To estimate a DIIS luniuositv. we developed a simple scenario where the passing DII is able to collect the eas within a radius of influence determined by its mwas and periceuter velocity: Mion=NugantGAMpgeg).," To estimate a BH's luminosity, we developed a simple scenario where the passing BH is able to collect the gas within a radius of influence determined by its mass and pericenter velocity: $M_{coll} = \Sigma_{disk} \pi (GM_{BH}/v_{BH}^2)^2$."
+ We perform a Monte Carlo simulation taking iuto account the observed ranges of DII velocities. and estimate that the amount of mass collected by a BIT during its disk passage ranges from 10?i M.," We perform a Monte Carlo simulation taking into account the observed ranges of BH velocities, and estimate that the amount of mass collected by a BH during its disk passage ranges from $10^{-3} - 10^{-4}$ ."
+" Assundug au accretion hundnuositv {τνyale? with jj=0 aud the canonical Iuninosity of Eepy=107 coves Ἡ, the DIE will acerete and radiate in a high state for a mean duration of 2000 years. though a slower BIL may radiate for up to 10 vears."," Assuming an accretion luminosity $L \sim \eta \dot M c^2$ with $\eta =
+0.1$ and the canonical luminosity of $L_{ULX} = 10^{39}$ ergs $^{-1}$, the BH will accrete and radiate in a high state for a mean duration of 2000 years, though a slower BH may radiate for up to $10^5$ years."
+ However. if the ULN transitions from a high state to a low state at anv point. the accretion rate would drastically decrease and the observed cussion could continue for nich longer.," However, if the ULX transitions from a high state to a low state at any point, the accretion rate would drastically decrease and the observed emission could continue for much longer."
+ This simple scenario. however. is not sufficicut to reproduce the properties of IILAX-1l. which is at least l1 kpe away from the plane of its host galaxy.," This simple scenario, however, is not sufficient to reproduce the properties of HLX-1, which is at least 1 kpc away from the plane of its host galaxy."
+ A DII passing through the disk moving at a few πάτος lan + will oulv travel a distauce of ~ tens of parsecs during the duration of the predicted accretion event., A BH passing through the disk moving at a few hundred km $^{-1}$ will only travel a distance of $\sim$ tens of parsecs during the duration of the predicted accretion event.
+ Such an event would appear to be a disk ULX source. possibly iudistinguishable from those cospatial with star formation regions.," Such an event would appear to be a disk ULX source, possibly indistinguishable from those cospatial with star formation regions."
+ HILX-1l is also 1000 times more huumous aud LOO times less massive (Wiersenia et al., HLX-1 is also 1000 times more luminous and 100 times less massive (Wiersema et al.
+ 2010..," 2010.,"
+ in prep.), in prep.)
+ than the assmnued values iu our calculation above., than the assumed values in our calculation above.
+ Taking all of these factors iuto account. it is unlikely that an isolated swandermg DII passing through a ealaxy disk can reproduce the properties of off-nuclear ULXs.," Taking all of these factors into account, it is unlikely that an isolated wandering BH passing through a galaxy disk can reproduce the properties of off-nuclear ULXs."
+ ILowever. our wandering BIT scenario would be feasible if IILN-1 is not au isolated DIT. but is traveling with a bound chump of gas and stars.," However, our wandering BH scenario would be feasible if HLX-1 is not an isolated BH, but is traveling with a bound clump of gas and stars."
+ The magnitude aud colors of the detected optical counterpart of ITELN-1 are consistent with a elobular cluster (Soriaetal.2009)., The magnitude and colors of the detected optical counterpart of HLX-1 are consistent with a globular cluster \citep{Soria09}.
+. This cluster may in fact be the stellar reminauts of a stripped dwarf ealaxy core that is still bouud to the DIL, This cluster may in fact be the stellar remnants of a stripped dwarf galaxy core that is still bound to the BH.
+ IILX-1l could be similar to the object Gl im MOI. a elobular cluster that may actually be the uucleus of a stripped dwarf galaxy (Alevlanetal.2001) and harbors the most promising Local Group caudidate for an IMDBII (Ulvestadetal.2007).," HLX-1 could be similar to the object G1 in M31, a globular cluster that may actually be the nucleus of a stripped dwarf galaxy \citep{Meylan01} and harbors the most promising Local Group candidate for an IMBH \citep{Ulvestad07}."
+. If we presume that the wandering DII in our simulations retain a nuclear star cluster aud eas reservoir from their pareut halos. then the iustauce of a wandering BIT passiug near the ceuter of the primary could cause instabilities iu its accretion disk. trigecring an accretion eveut of sufficient magnitude to power a ULX with a high luminosity as seen in HLX-1.," If we presume that the wandering BHs in our simulations retain a nuclear star cluster and gas reservoir from their parent halos, then the instance of a wandering BH passing near the center of the primary could cause instabilities in its accretion disk, triggering an accretion event of sufficient magnitude to power a ULX with a high luminosity as seen in HLX-1."
+ Our simulations do not have sufficient resolutiou to follow the tidally stripped cores of galaxies im detail. though other siuulations have shown that a tidally stripped chwart ealaxy can retain its core after a close passage with the primary (Alaver&Wadslev2001).," Our simulations do not have sufficient resolution to follow the tidally stripped cores of galaxies in detail, though other simulations have shown that a tidally stripped dwarf galaxy can retain its core after a close passage with the primary \citep{Mayer03}."
+ Thus. iu the likely instance that a wandering DIT retains the core of its host ealaxy. its passage near the galaxy disk cau explain the origin aud propertics of ITLA-1.," Thus, in the likely instance that a wandering BH retains the core of its host galaxy, its passage near the galaxy disk can explain the origin and properties of HLX-1."
+ Previous studies have estimated hunuinosities for massive DIIS wandering through the ISAL but prior to this Letter none have explored the issue in a cosinological context.," Previous studies have estimated luminosities for massive BHs wandering through the ISM, but prior to this Letter none have explored the issue in a cosmological context."
+ Ivolik(2001). showed that INMDBIIS with masses ranging from 107101 can produce Imninosities of ULXs if they pass through or near molecular clouds., \citet{Krolik04} showed that IMBHs with masses ranging from $10^2 - 10^4$ can produce luminosities of ULXs if they pass through or near molecular clouds.
+" SMDITI gravitational recoil eveuts passing through the disk iuav exhibit N-rav eissiou of L>10°"" cress ! (Fujita2009).", SMBH gravitational recoil events passing through the disk may exhibit X-ray emission of $L > 10^{39}$ ergs $^{-1}$ \citep{Fujita09}.
+. Mapollietal.(2008) performed an N-body|SPIT simulation of IAIBUs in a galaxy merger. and found that a few halo IMDIIS reside in orbits that pass through the disk. which iu our scenario may be observable as ULXs.," \citet{Mapelli08} performed an N-body+SPH simulation of IMBHs in a galaxy merger, and found that a few halo IMBHs reside in orbits that pass through the disk, which in our scenario may be observable as ULXs."
+ We provide compclling dynamical scenario for the presence of “wauderime” iassive DIIS in the halos of ealaxies., We provide compelling dynamical scenario for the presence of “wandering” massive BHs in the halos of galaxies.
+ A natural consequence of the hierarchical build up of ealaxies in a ACDAL scenario. the tidal stripping of galaxies containing seed DIIS cau populate the halo of a massive disk galaxv with wandering BUs.," A natural consequence of the hierarchical build up of galaxies in a $\Lambda$ CDM scenario, the tidal stripping of galaxies containing seed BHs can populate the halo of a massive disk galaxy with wandering BHs."
+" These objects often retain their original seed mass. are found hroughout the galaxy halo. aud may pass through the ealactic disk at απ average rate of 10.6 Gyr1,"," These objects often retain their original seed mass, are found throughout the galaxy halo, and may pass through the galactic disk at an average rate of 10.6 $^{-1}$."
+ We predict hat Local Croup dwarf galaxies such as the \aecllanic Clouds are likely to host IMDIIs., We predict that Local Group dwarf galaxies such as the Magellanic Clouds are likely to host IMBHs.
+ Detections of these wandering DIIS may sivo an upper lait to the initial nass of BIT seeds. and may allow us to differentiate )etwoeen the various proposed formation mechanisms of such seeds.," Detections of these wandering BHs may give an upper limit to the initial mass of BH seeds, and may allow us to differentiate between the various proposed formation mechanisms of such seeds."
+ Our scenario provides a plivsically motivated explanation for offunclear ULXs as IMDIIs which have ο. stripped from their host galaxies. if they retain a eas rescrvoirfaccretion disk that. when cynamically destabilized. uuelt be fimneled toward the black hole and form an accretion disk.," Our scenario provides a physically motivated explanation for off-nuclear ULXs as IMBHs which have been stripped from their host galaxies, if they retain a gas reservoir/accretion disk that, when dynamically destabilized, might be funneled toward the black hole and form an accretion disk."
+ JMD evatefully acknowledges the National Science Foundation Caaduate Research Fellowship Program., JMB gratefully acknowledges the National Science Foundation Graduate Research Fellowship Program.
+ FO acknowledges support from IST GO-1125. NASA ITP NNXNOSAGCSIC. and NSF ITR eraut. PITY-0205113 (also. supporting TRO).," FG acknowledges support from HST GO-1125, NASA ITP NNX08AG84G, and NSF ITR grant PHY-0205413 (also supporting TRQ)."
+ ATV acknowledges support from SAO TAI9-0006N and NASA ATP NNNLOAC'SLG awards., MV acknowledges support from SAO TM9-0006X and NASA ATP NNX10AC84G awards.
+ Simulations were runusing computer resourees and techuicalsupport from NAS., Simulations were runusing computer resources and technicalsupport from NAS.
+ The authors thank Charlotte Christensen. Advicnne Stilp. Amy IKinball aud Braut Robertson for assistance aud helpful coments. as," The authors thank Charlotte Christensen, Adrienne Stilp, Amy Kimball and Brant Robertson for assistance and helpful comments, as"
+We also plot b against the mean number of subhalos in a halo. finding a steeper relationship (in that b increases faster for à given increase inlog Nu).,"We also plot $b$ against the mean number of subhalos in a halo, finding a steeper relationship (in that $b$ increases faster for a given increase in$\log N_{\rm sub}$ )."
+" Because INS, increases more slowly than Mya, as My, is increased (see e.g. Figure 2)) this behaviour is expected."," Because $N_{\rm sub}$ increases more slowly than $M_{\rm halo}$ as $M_{\rm halo}$ is increased (see e.g., Figure \ref{hod}) ) this behaviour is expected."
+ We show b as a function of Now for two dillerent subhalo mass thresholds. finding that b increases faster lor a higher threshold mass.," We show $b$ as a function of $N_{\rm sub}$ for two different subhalo mass thresholds, finding that $b$ increases faster for a higher threshold mass."
+" “This again is expected: as for higher mass thresholds the HOD curves stay longer on the shallower initial part of the Now.Mya, relationship (see e.g.. Figure 3))."," This again is expected as for higher mass thresholds the HOD curves stay longer on the shallower initial part of the $N_{\rm sub}-M_{\rm halo}$ relationship (see e.g., Figure \ref{hod2}) )."
+ We take samples of halos which are chosen to cither have fixed halo mass or fixed number of subhalos (as defined in Section 2.3). and compute how the large scale bias byeypμ]] changes as we vary the other parameter.," We take samples of halos which are chosen to either have fixed halo mass or fixed number of subhalos (as defined in Section 2.3), and compute how the large scale bias $b_{\rm HOD/all}$ changes as we vary the other parameter."
+ Our results are shown in Ligure 9.. where on the αι axis we show the percentiles of the distribution.," Our results are shown in Figure \ref{bias1}, where on the $x-$ axis we show the percentiles of the distribution."
+ For example in the top panel a point abr=[99% shows byobeny for the top La of halos by occupation. at fixed halo mass.," For example in the top panel a point at $x=+99\%$ shows $b_{\rm HOD/all}$ for the top $1\%$ of halos by occupation, at fixed halo mass."
+ Likewise a point at ee99% shows byopea for the bottom I4 of halos by occupation. at .r—50% for the bottom 50% ancl so on.," Likewise a point at $x=-99\%$ shows $b_{\rm HOD/all}$ for the bottom $1\%$ of halos by occupation, at $x=-50\%$ for the bottom $50\%$ and so on."
+ Points at à=0 show results for all halos in that halo mass xn. and the bias μου is measured relative to this (as in equation 1).," Points at $x=0$ show results for all halos in that halo mass bin, and the bias $b_{\rm HOD/all}$ is measured relative to this (as in equation 1)."
+ This is cillerent from 6 computed with respect o the dark matter distribution. and ensures that all curves xs through byepeay=1 at wr=0.," This is different from $b$ computed with respect to the dark matter distribution, and ensures that all curves pass through $b_{\rm HOD/all}=1$ at $x=0$."
+ In the top panel of Figure 9.. the values of byopsan are shown for variations in the Now percentile. at. fixed. halo mass.," In the top panel of Figure \ref{bias1}, the values of $b_{\rm HOD/all}$ are shown for variations in the $N_{\rm sub}$ percentile, at fixed halo mass."
+ We show results for two cdillerent halo masses. and two cdillerent. cutolls in the mass of a subhalos.," We show results for two different halo masses, and two different cutoffs in the mass of a subhalos."
+ The mean halo mass {νο shown in the figure caption for the curves was computed by averaging over the masses of halos above a mass threshold., The mean halo mass $\langle M_{\rm halo}\rangle$ shown in the figure caption for the curves was computed by averaging over the masses of halos above a mass threshold.
+ As required. the (Aya) are approximately constant for the dillerent bins of Now percentile. ancl in all cases within 104 of the £A) value given.," As required, the $\langle M_{\rm halo}\rangle$ are approximately constant for the different bins of $N_{\rm sub}$ percentile, and in all cases within $10\%$ of the $\langle M_{\rm halo}\rangle$ value given."
+ We can see that byobsay does change significantlywith Nou. percentile. even though the mean halo mass is the same for all points along the curve.," We can see that $b_{\rm HOD/all}$ does change significantlywith $N_{\rm sub}$ percentile, even though the mean halo mass is the same for all points along the curve."
+ This is what we expect from looking at. Figures l and 6., This is what we expect from looking at Figures \ref{slice} and \ref{xibias}.
+ We can also see that the steepest change in b occurs for the larger mass halos., We can also see that the steepest change in $b$ occurs for the larger mass halos.
+ This analysis is equivalent to adding together the trend in b that one gets by moving up cilferent columns in Figure 7.., This analysis is equivalent to adding together the trend in $b$ that one gets by moving up different columns in Figure \ref{biasgrid}.
+ If we now turn to the results showing how 6 varies as Adyas ds changed for fixed Nou. (bottom panel of Figure 9)). we can see that the situation is somewhat different.," If we now turn to the results showing how $b$ varies as $M_{\rm halo}$ is changed for fixed $N_{\rm sub}$ (bottom panel of Figure \ref{bias1}) ), we can see that the situation is somewhat different."
+" In this plot. the (Nou) values shown are computed from the average number of subhalos above a threshold in Now. for each of the bins of Ajay, percentile."," In this plot, the $\langle N_{\rm sub}\rangle$ values shown are computed from the average number of subhalos above a threshold in $N_{\rm sub}$, for each of the bins of $M_{\rm halo}$ percentile."
+ Here we see that [or a low number of subhalos. {Niu}=1.2. there is a strong trend of b with halo mass. one which is strong for both values of subhalo mass cutoll.," Here we see that for a low number of subhalos, $\langle N_{\rm sub}\rangle=1.2$, there is a strong trend of $b$ with halo mass, one which is strong for both values of subhalo mass cutoff."
+ However when we increase the ING threshold to give a higher value of CN~ S. we find that there is no dependence of byopsda on Mya; percentile.," However when we increase the $N_{\rm sub}$ threshold to give a higher value of $\langle N_{\rm sub}\rangle\sim 8$ , we find that there is no dependence of $b_{\rm HOD/all}$ on $M_{\rm halo}$ percentile."
+ This independence of halo clustering and halo mass is interesting enough that we revisit it in Figure LO. , This independence of halo clustering and halo mass is interesting enough that we revisit it in Figure \ref{bias2}. .
+where we now turn to plotting b with respect to dark matter halo clustering., where we now turn to plotting $b$ with respect to dark matter halo clustering.
+ The scale of the elfects can then be judged, The scale of the effects can then be judged
+"We applied both tests to both the full data. sample and also restricting the source luminosities to the range 10°""erestoeLye107eres.","We applied both tests to both the full data sample and also restricting the source luminosities to the range $10^{36} {\rm \, erg \, s^{-1}} < L_X < 10^{38} {\rm \, erg \,
+s^{-1}}$."
+ In alb cases. the distributions of the spatial displacements for the X-ray sources in the three galaxies are consistent with being drawn from the same distribution.," In all cases, the distributions of the spatial displacements for the X-ray sources in the three galaxies are consistent with being drawn from the same distribution."
+ The median clisplacement for NGC 1569 appears somewhat smaller than for the other ealaxies. but the dillerence is not statistically significant.," The median displacement for NGC 1569 appears somewhat smaller than for the other galaxies, but the difference is not statistically significant."
+ The confidence level error intervals for the median overlap for all three galaxies., The confidence level error intervals for the median overlap for all three galaxies.
+ Fig., Fig.
+ 3. shows the luminosity of cach X-ray source plotted versus its spatial displacement. [rom the nearest star cluster., \ref{lumd} shows the luminosity of each X-ray source plotted versus its spatial displacement from the nearest star cluster.
+ “Phe striking feature of the plot is that there are no high luminosity sources at large displacements from star clusters., The striking feature of the plot is that there are no high luminosity sources at large displacements from star clusters.
+ Furthermore. there is an apparent trend. of decreasing luminosity with increasing displacement from the nearest star cluster.," Furthermore, there is an apparent trend of decreasing luminosity with increasing displacement from the nearest star cluster."
+ To evaluate if this trend is statistically significant. we performed. a boot-strap! analysis.," To evaluate if this trend is statistically significant, we performed a `boot-strap' analysis."
+ We used the set of displacements ancl luminosities shown in Fig., We used the set of displacements and luminosities shown in Fig.
+ 3. and then randomly re-arranged the pairings., \ref{lumd} and then randomly re-arranged the pairings.
+ This corresponds to the null hypothesis in which luminosity ancl displacement are unrelated., This corresponds to the null hypothesis in which luminosity and displacement are unrelated.
+" In. LO"" trials. we found 1920 cases in which no sources lav in the same region of the luninosity-displacement plot which is empty. in. Fig. 3..."," In $10^7$ trials, we found 1920 cases in which no sources lay in the same region of the luminosity-displacement plot which is empty in Fig. \ref{lumd}."
+ This corresponds to à chance probability of occurrence of 10 7.," This corresponds to a chance probability of occurrence of $1.9 \times
+10^{-4}$ ."
+ For MS2 there is also an absence of dim sources at small displacements., For M82 there is also an absence of dim sources at small displacements.
+ “This is likely due to the high level of diffuse X-ray emission in the central 200 pe of M82. (Cirillithsetal.2000). which precludes detection of dim sources in the central regions of MS2., This is likely due to the high level of diffuse X-ray emission in the central 200 pc of M82 \cite{griffiths00} which precludes detection of dim sources in the central regions of M82.
+ Dim sources are detected very near star clusters in NGC 1569., Dim sources are detected very near star clusters in NGC 1569.
+ We have shown that X-ray. sources in these three starburst galaxies are preferentially located near star clusters., We have shown that X-ray sources in these three starburst galaxies are preferentially located near star clusters.
+ Because the star clusters are very good tracers of current. star formation activity in the galaxies. this confirms that the X-rav binaries are voung objects associated with current star formation.," Because the star clusters are very good tracers of current star formation activity in the galaxies, this confirms that the X-ray binaries are young objects associated with current star formation."
+ We also found. significant. displacements of the X-ray sources [rom the clusters., We also found significant displacements of the X-ray sources from the clusters.
+ Because X-ray binaries. unlike other bright X-ray sources such as voung supernovae. can exhibit. high. velocities. this suggests that much of the X-rav source population consists of X-rav. binaries.," Because X-ray binaries, unlike other bright X-ray sources such as young supernovae, can exhibit high velocities, this suggests that much of the X-ray source population consists of X-ray binaries."
+ Several mechanisms can put à binary in motion., Several mechanisms can put a binary in motion.
+ For neutron star binaries. a kick’ due to an asymmetric explosion in the formation of the neutron star can lead to high. velocities (Lyne&Lorimer1994).," For neutron star binaries, a `kick' due to an asymmetric explosion in the formation of the neutron star can lead to high velocities \cite{lyne94}."
+. Even in the absence of ‘kicks’ [rom supernova explosions. momentum conservation following a symmetric ejection of matter in the formation of the neutron star or black hole in a binary with a high mass companion can produce a runaway speed of 50kms the ejected matter continues to move with the instantaneous. orbital velocity of the compact object at the moment of ejection and the binary must move in the opposite direction. to conserve momoentum (Nelemans.Tauris.&vandenHeuvel1999:vandenLleuvelοἱal. 2000).," Even in the absence of `kicks' from supernova explosions, momentum conservation following a symmetric ejection of matter in the formation of the neutron star or black hole in a binary with a high mass companion can produce a runaway speed of $\sim 50 \rm \, km \, s^{-1}$; the ejected matter continues to move with the instantaneous orbital velocity of the compact object at the moment of ejection and the binary must move in the opposite direction to conserve momentum \cite{nelemans99,vdh00}."
+. For binaries in clusters. interactions with other stars and binaries in the cluster can eject the binary from the cluster (Phinney&Sigurdssonquist. 1993).," For binaries in clusters, interactions with other stars and binaries in the cluster can eject the binary from the cluster \cite{phinney91,kulkarni93,sigurdsson93}."
+. In voung. dense star clusters such interactions can occur on time scales ofa few Myr (PortegiesZwartetal. 1999).," In young, dense star clusters such interactions can occur on time scales of a few Myr \cite{portegies99}."
+. Objects ejected via dynamical interactions tend to escape with close to the mininium energy needed to escape (Joshi.Nave.&Rasio2001)., Objects ejected via dynamical interactions tend to escape with close to the minimum energy needed to escape \cite{joshi01}.
+. The runaway velocities. (at infinity) should. be of the same magnitude as the stellar velocity. dispersions of the clusters. which are. typically 10I5kms in these galaxies (Smith&Gallagher2001:AMeCrady.Gilbert.&Craham. 2008).," The runaway velocities (at infinity) should be of the same magnitude as the stellar velocity dispersions of the clusters, which are typically $10-15 \, \rm km \, s^{-1}$ in these galaxies \cite{smith01,mccrady03}."
+. The displacements we observe in the starburst) galaxies are likely due to motion of the X-ray sources caused by one or more of these mechanisms., The displacements we observe in the starburst galaxies are likely due to motion of the X-ray sources caused by one or more of these mechanisms.
+ Furthermore. we found that there is an absence of bright X-ray sources with large clisplacements.," Furthermore, we found that there is an absence of bright X-ray sources with large displacements."
+ This suggests that there is some correlation. between the maximum possible brightness of an X-ray source ancl its motion., This suggests that there is some correlation between the maximum possible brightness of an X-ray source and its motion.
+ This correlation appears to hold. only [ο (isotropic equivalent)luminosities above 107eres," This correlation appears to hold only for (isotropic equivalent)luminosities above $10^{38} \rm \, erg \, s^{-1}$."
+" The excluced region appears to be bounded by a linear relation between (isotropic equivalent) source Luminosity Ly and source displacement. from the nearest star cluster dL,<(110)oresLyfe pe)."," The excluded region appears to be bounded by a linear relation between (isotropic equivalent) source luminosity $L_X$ and source displacement from the nearest star cluster $d$, $L_x < (1 \times 10^{41} \, {\rm erg
+\, s^{-1}}) / (d/{\rm pc})$ ."
+ lo discussing this correlation. we first consider the case where the X-ray sources emit isotropically.," In discussing this correlation, we first consider the case where the X-ray sources emit isotropically."
+ In this case. he systems producing such high luminosities likely contain lack holes accreting via Roche lobe overllow because such veh luminosities would. be cdillicult to achieve in a wind accretor due to the low ellicieney of wind. capture and slack holes ancl needed to not violate the Exdington limit (Bloncin.Stevens.&Ixallman1991:Petterson 1978)..," In this case, the systems producing such high luminosities likely contain black holes accreting via Roche lobe overflow because such high luminosities would be difficult to achieve in a wind accretor due to the low efficiency of wind capture and black holes and needed to not violate the Eddington limit \cite{blondin91,petterson78}. ."
+ If the X-ray sources are ejected from the star clusters with speeds which are roughly. independent. of mass. then he inverse correlation between maximum X-ray source," If the X-ray sources are ejected from the star clusters with speeds which are roughly independent of mass, then the inverse correlation between maximum X-ray source"
+the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology. funded by the National Aeronautics and Space Administration and the National Science Foundation.,"the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation."
+ This research made use of data products from the Midcourse Space Experiment. funded by the Ballistic Missile Defense Organization with additional support from NASA Oftice of Space Science.," This research made use of data products from the Midcourse Space Experiment, funded by the Ballistic Missile Defense Organization with additional support from NASA Office of Space Science."
+"Since we did not detect any !?CO line, we convert our 3o noise into an upper limit on the flux.","Since we did not detect any $^{12}$ CO line, we convert our $3\sigma$ noise into an upper limit on the flux."
+ A knowledge of the flux will enable us to estimate upper limits on gas disk mass., A knowledge of the flux will enable us to estimate upper limits on gas disk mass.
+" The observed flux is the double-integral of theobserved intensity /, over frequency v and solid angle €): Assuming the brightness temperature does not vary substantially over the telescope beam and that the lineis a Gaussian with line-width Av, then the upper limit on the 3c line flux F is The 3c upper limits on F are listed in Table B]."," The observed flux is the double-integral of theobserved intensity $I_\nu$ over frequency $\nu$ and solid angle $\Omega$: Assuming the brightness temperature does not vary substantially over the telescope beam and that the lineis a Gaussian with line-width $\Delta v$, then the upper limit on the $3\sigma$ line flux $F$ is The $3\sigma$ upper limits on $F$ are listed in Table \ref{obssummary}."
+" Similarly, we can derive the column density in the optically thin limit to be where Ay is the Einstein A coefficient (spontaneous emission) and has units of s!."," Similarly, we can derive the column density in the optically thin limit to be where $A_{ul}$ is the Einstein A coefficient (spontaneous emission) and has units of $^{-1}$."
+" Similar to the analysis of (2006),, we assume an excitation temperature Tex7:20 K. Then in our case for '?CO J=2—1, F(Tex,Eu,v)©28.00."," Similar to the analysis of , we assume an excitation temperature $T_\mathrm{ex}\approx20$ K. Then in our case for $^{12}$ CO $J=2-1$, $\mathcal{F}(T_\mathrm{ex}, E_u, \nu)\approx28.00$."
+ The Einstein Ay=6.91x1077 s~ and partition function Q(20K)=15.9 (CDM9?))., The Einstein $A_{ul}=6.91\times10^{-7}$ $^{-1}$ and partition function $Q(20\mathrm{~K})=15.9$ ).
+ We compute and tabulate in Table 3| the !?CO number densities and Hy gas masses in the optically thin limit., We compute and tabulate in Table \ref{obssummary} the $^{12}$ CO number densities and $_2$ gas masses in the optically thin limit.
+" Gas disk masses were derived from (1986)),, where [Hj/CO]z 103, ug=1.36 is the mean molecular weight, my, is themass of an Hy molecule, and d=140 pc is the distance to Taurus."," Gas disk masses were derived from , where $\left[\mathrm{H}_2/\mathrm{CO}\right]\approx10^4$ , $\mu_G=1.36$ is the mean molecular weight, $m_{\mathrm{H}_2}$ is themass of an $_2$ molecule, and $d\approx140$ pc is the distance to Taurus."
+stars.,stars.
+ Similar analyses for more stars is needed to confirm this in a statistical sense., Similar analyses for more stars is needed to confirm this in a statistical sense.
+ Only some stars show pulsation characteristics of both > Dor and 9 Scuti type (c.g.Henry&Fekel 2005)., Only some stars show pulsation characteristics of both $\gamma$ Dor and $\delta$ Scuti type \citep[e.g.][]{henry2005}.
+. Interestingly. most. of these. so-called. hybrid stars appear to be metallic line (Am) stars.," Interestingly, most of these, so-called, hybrid stars appear to be metallic line (Am) stars."
+ Furthermore. some verv interesting spectroscopic multiple svstems with of X / EF components are known. such as DG Leo (Erématetal.2005a) and 87 Tau (Lampensetal. 2007).," Furthermore, some very interesting spectroscopic multiple systems with of A / F components are known, such as DG Leo \citep{fremat2005b} and $\theta ^2$ Tau \citep{lampens2007}."
+. Phese systems have at least one pulsating component and a non-pulsating or stable one., These systems have at least one pulsating component and a non-pulsating or stable one.
+ Further investigation of these svstems in terms of chemical composition would allow for a direct study of the relation between multiplicity. »ulsations and chemical composition.," Further investigation of these systems in terms of chemical composition would allow for a direct study of the relation between multiplicity, pulsations and chemical composition."
+ Abundance determination is not only very interesting. as emphasised by the science cases described above. but also xwticularlv. challenging for A- and E-type stars. because of the occurrence of rapid rotation and/or blended lines in hese stars.," Abundance determination is not only very interesting, as emphasised by the science cases described above, but also particularly challenging for A- and F-type stars, because of the occurrence of rapid rotation and/or blended lines in these stars."
+ So far. only a few chemical abundance analysis ools for these stars have been developed by. e.g. Erspamer&North(2002b) and Drunttetal.(2002.2004.2008). which is based on Valenti&Piskunov(1996).," So far, only a few chemical abundance analysis tools for these stars have been developed by, e.g. \citet{erspamer2002} and \citet{bruntt2002,bruntt2004,bruntt2008} which is based on \citet{valenti1996}."
+. With the spectra we ave obtained for all science cases described above in mind. we independently developed: a semi-automatie procedure o determine fundamental stellar. parameters (Section 2) and chemical abundances (Section. 3). which we test. on simulated data (Section 4).," With the spectra we have obtained for all science cases described above in mind, we independently developed a semi-automatic procedure to determine fundamental stellar parameters (Section 2) and chemical abundances (Section 3), which we test on simulated data (Section 4)."
+ Phe main advantage of such tests is that the input stellar parameters are well known a priory., The main advantage of such tests is that the input stellar parameters are well known a priory.
+ These tests can provide an estimate of the error due to the procedure. and the uncertainties due to errors in the stellar parameters.," These tests can provide an estimate of the error due to the procedure, and the uncertainties due to errors in the stellar parameters."
+ Furthermore. gravitational darkening cLlects induced by fast rotation allect the abundance determination of stars seen nearly pole-on.," Furthermore, gravitational darkening effects induced by fast rotation affect the abundance determination of stars seen nearly pole-on."
+. We analysecl Vega (Section 5) to compare results obtained. with our method. with literature values. ancl subsequently investigated the effects on abundance determination. for. simulated: fast-rotating stars (Frématctal.2005b) seen at several low inclination angles (Section 6).," We analysed Vega (Section 5) to compare results obtained with our method with literature values, and subsequently investigated the effects on abundance determination for simulated fast-rotating stars \citep{fremat2005} seen at several low inclination angles (Section 6)."
+ The radial velocity (RV). projected rotational velocity (osin ). elfective temperature (Liar) and surface gravity (log ο) are determined. from. spectra using a procedure. partly owed. on the fitting of high-resolution spectroscopic data with svnthetic spectra interpolated in a flux: &rid.," The radial velocity (RV), projected rotational velocity $\upsilon \sin i$ ), effective temperature $_{\rm eff}$ ) and surface gravity $\log$ g) are determined from spectra using a procedure partly based on the fitting of high-resolution spectroscopic data with synthetic spectra interpolated in a flux grid."
+ This erid was prepared by applving the SYNSPISC programme (Ilubeny&Lanz1995.anclreferences.therein) and hy adding opacities sources due to Rayleigh scattering. and o the 1 ions.," This grid was prepared by applying the SYNSPEC programme \citep[][and references therein]{hubeny1995}, and by adding opacities sources due to Rayleigh scattering and to the $^{-}$ ions."
+ JE calculations were performed. with the ATLASS LTE atmosphere models computed by Castelli&Ixurucz (2003)., All calculations were performed with the ATLAS9 LTE atmosphere models computed by \citet{castelli2003}.
+. Phe microturbulent velocity was fixed at 2 kiss+. while a solar-type chemical composition (CirevesseSauval1998). was considered.," The microturbulent velocity was fixed at 2 $^{-1}$, while a solar-type chemical composition \citep{grevesse1998} was considered."
+ The procedure we follow for stellar parameter determination (RV. osin. Farlog ο) consists of four consecutive steps (seoalso.brématctal.2007).," The procedure we follow for stellar parameter determination (RV, $\upsilon \sin i$, $_{\rm eff}$ $\log$ g) consists of four consecutive steps \citep[see also][]{fremat2007}."
+.. First. RY is determined rom the cross-correlation function. (CCL) computed. with a svnthetic spectrum. corresponding to the stars. known μα»ectral type.," First, RV is determined from the cross-correlation function (CCF) computed with a synthetic spectrum corresponding to the star's known spectral type."
+ Then. esin is determined based on metallic ine fitting in the sensitive regions indicated in Table 1..," Then, $\upsilon \sin i$ is determined based on metallic line fitting in the sensitive regions indicated in Table \ref{vsiniregions}."
+ The fitting is performed. with the Girfit. computer code (lrématctal.2006).. where the Minuit. Fortran package eveloped and. applied. by CERN is used to perform least squares minimisation based on a simplex method (Press 2002).," The fitting is performed with the Girfit computer code \citep{fremat2006}, where the Minuit Fortran package developed and applied by CERN is used to perform least squares minimisation based on a simplex method \citep{press2002}."
+. All regions in Table 1. are fitted independently with 10 previously. determined. RN. while log &. Fr and. esin; we free parameters.," All regions in Table \ref{vsiniregions} are fitted independently with the previously determined RV while $\log$ g, $_{\rm eff}$ and $\upsilon \sin i$ are free parameters."
+ The stars we are concerned with are generally. cooler man SSOO WK. meaning that the hydrogen lines are sensitive o the cllective temperature. ancl insensitive to. Stark oadening and surface gravity variations.," The stars we are concerned with are generally cooler than 8500 K, meaning that the hydrogen lines are sensitive to the effective temperature, and insensitive to Stark broadening and surface gravity variations."
+ Therefore. we use Balmer lines to determine Tar.," Therefore, we use Balmer lines to determine $_{\rm eff}$."
+ For both A- and E-tvpe stars. we use here the LE2 and La lines. since around these lines we can obtain reliable continuum normalisation.," For both A- and F-type stars, we use here the $\beta$ and $\alpha$ lines, since around these lines we can obtain reliable continuum normalisation."
+ In early A-νρο stars. a reliable continuum could also be defined around he Le and Le lines. but from tests on synthetic spectra we found that this is not the case for cooler stars. which are relatively more alfected by metallic line blends.," In early A-type stars, a reliable continuum could also be defined around the $\gamma$ and $\delta$ lines, but from tests on synthetic spectra we found that this is not the case for cooler stars, which are relatively more affected by metallic line blends."
+ For this reason. and in order to conduct a homogeneous analysis of a sample of A ancl E stars. we do not take these lines into account.," For this reason, and in order to conduct a homogeneous analysis of a sample of A and F stars, we do not take these lines into account."
+ LE3 and Lla are fitted with the same packages as described above. with the previously obtained: values for eosin? and BN. while log & is set to 4.0 in the first iteration.," $\beta$ and $\alpha$ are fitted with the same packages as described above, with the previously obtained values for $\upsilon \sin i$ and RV, while $\log$ g is set to 4.0 in the first iteration."
+ Finally. log g has been derived by combining the llipparcos parallax. apparent magnitude (reddening is not taken into account at the moment as we currently only have stars of 2H magnitude. and brighter. at. our disposal). aud derived Tr to obtain the stars’ luminosity (erspamer& 2003)..," Finally, $\log$ g has been derived by combining the Hipparcos parallax, apparent magnitude (reddening is not taken into account at the moment as we currently only have stars of $^{th}$ magnitude and brighter at our disposal) and derived $_{\rm eff}$ to obtain the stars' luminosity \citep{erspamer2003}. ."
+ Tar ane luminosity values are then used to, $_{\rm eff}$ and luminosity values are then used to
+which can be explained with similar excitation schemes of C-type transitions in this energy range for the neutrals and the cations.,which can be explained with similar excitation schemes of C-type transitions in this energy range for the neutrals and the cations.
+ The spectra of ionized HBC in different inert gas matrices illustrate another important point., The spectra of ionized HBC in different inert gas matrices illustrate another important point.
+ The widths and shapes of the broad bands at wavelengths shorter than 600 nm (in the case of HBC) do not depend on the specific matrix environment., The widths and shapes of the broad bands at wavelengths shorter than 600 nm (in the case of HBC) do not depend on the specific matrix environment.
+ These bands are much broader than what is expected from typical matrix-induced broadening effects indicating very short lifetimes of the excited states which are not entirely caused by the interaction with the matrix atoms. (, These bands are much broader than what is expected from typical matrix-induced broadening effects indicating very short lifetimes of the excited states which are not entirely caused by the interaction with the matrix atoms. (
+"Alternatively, other broadening mechanisms, like Franck-Condon vibronic broadening, may be considered (see,e.g.,Sassaraetal. 2001)","Alternatively, other broadening mechanisms, like Franck-Condon vibronic broadening, may be considered \citep[see, e.g.,][]{sassara01}. .)"
+".. Therefore, absorption spectra of cold ionized HBC in the gas phase (below 600 nm) will most likely resemble the presented MIS spectra."," Therefore, absorption spectra of cold ionized HBC in the gas phase (below 600 nm) will most likely resemble the presented MIS spectra."
+" However, a small shift of the band positions to shorter wavelengths can be expected."," However, a small shift of the band positions to shorter wavelengths can be expected."
+" If the interstellar UV bump at 217.5 nm is indeed a collective feature of the interstellar PAHs that are observed through their ΤΗ, emission signatures, its observational constraints could be explained as follows."," If the interstellar UV bump at 217.5 nm is indeed a collective feature of the interstellar PAHs that are observed through their IR emission signatures, its observational constraints could be explained as follows."
+ The varying width for different lines of sight might be an indication for different degrees of ionization., The varying width for different lines of sight might be an indication for different degrees of ionization.
+" Alternatively, this could just hint at different size distributions of the molecules. ("," Alternatively, this could just hint at different size distributions of the molecules. ("
+"Also, both explanations could be possible.)","Also, both explanations could be possible.)"
+ Its seemingly fixed position can probably be explained by a more or less stable mean size of the PAHs in the mixture., Its seemingly fixed position can probably be explained by a more or less stable mean size of the PAHs in the mixture.
+" Within the observational and computational limitations, it was already shown that the interstellar abundance of PAHs could suffice to explain the observed bump strength (Cecchi-Pestellini2008;Steglichetal. |2010)."," Within the observational and computational limitations, it was already shown that the interstellar abundance of PAHs could suffice to explain the observed bump strength \citep{cecchi08, steglich10}."
+". As suggested by astronomical observations, the outflows of carbon-rich stars inject −−hydrogenated amorphous carbon (HAC) into the ISM "," As suggested by astronomical observations, the outflows of carbon-rich stars inject hydrogenated amorphous carbon (HAC) into the ISM \citep{mathis90}."
+Recent laboratory investigations of such materials indicate the formation of more aromatic compounds under long-term FUV light exposure and the appearance of a bump-like absorption feature around 217.5 nm ⋅⋅, Recent laboratory investigations of such materials indicate the formation of more aromatic compounds under long-term FUV light exposure and the appearance of a bump-like absorption feature around 217.5 nm \citep{gadallah11}.
+" However, this feature is yet too broad in comparison with the interstellar UV bump, probably because the electronic π resonances in the HACs are still affected by chemical and physical bonding between the aromatic subunits."," However, this feature is yet too broad in comparison with the interstellar UV bump, probably because the electronic $\pi$ resonances in the HACs are still affected by chemical and physical bonding between the aromatic subunits."
+" Furthermore, these laboratory materials seem to need slightly more carbon than available for interstellar dust in order to reproduce the observed strength of the interstellar feature which may be rationalized by a high content of aliphatic structures."," Furthermore, these laboratory materials seem to need slightly more carbon than available for interstellar dust in order to reproduce the observed strength of the interstellar feature \citep{gadallah11} which may be rationalized by a high content of aliphatic structures."
+" Further chemical alteration by, e.g., FUV irradiation or interstellar shocks in the ISM could lead to the destruction of such structures and the release of aromatic units, i.e., large PAHs, with sizes sufficient to explain the MIR emission as well as the UV bump and the smooth extinction curve below 400 nm."," Further chemical alteration by, e.g., FUV irradiation or interstellar shocks in the ISM could lead to the destruction of such structures and the release of aromatic units, i.e., large PAHs, with sizes sufficient to explain the MIR emission as well as the UV bump and the smooth extinction curve below 400 nm."
+" These large PAHs, either neutral or ionized, with variable hydrogen saturation, may then represent a more or less stable end product in the ISM."," These large PAHs, either neutral or ionized, with variable hydrogen saturation, may then represent a more or less stable end product in the ISM."
+" Concerning the PAH hypothesis, we found no evidence that neutral and singly ionized cationic PAHs studied so far by us or other laboratory groups could be responsible for any of the DIBs."," Concerning the PAH hypothesis, we found no evidence that neutral and singly ionized cationic PAHs studied so far by us or other laboratory groups could be responsible for any of the DIBs."
+" Especially, the case of HBC, which is rather large and by comparison very stable, is compelling."," Especially, the case of HBC, which is rather large and by comparison very stable, is compelling."
+" The apparent absence of electronic features of neutral and cationic HBC in the diffuse ISM suggests that, due to molecular diversity, the abundances of individual PAH species of this size are probably too low to allow detection in the UV-VIS spectral domain."," The apparent absence of electronic features of neutral and cationic HBC in the diffuse ISM suggests that, due to molecular diversity, the abundances of individual PAH species of this size are probably too low to allow detection in the UV-VIS spectral domain."
+" Therefore, we conclude that the DIBs (at least the prominent ones) are most likely not signatures of the interstellar PAH family, if we consider only neutral and singly ionized cationic species with normal hydrogen This work was supported by the Deutsche Forschungsgemeinschaft (DFG, project no."," Therefore, we conclude that the DIBs (at least the prominent ones) are most likely not signatures of the interstellar PAH family, if we consider only neutral and singly ionized cationic species with normal hydrogen This work was supported by the Deutsche Forschungsgemeinschaft (DFG, project no."
+ Hu 474/21-2)., Hu 474/21-2).
+" The authors are grateful to Harald Mutschke and Kamel A. Khalil Gadallah for measuring the UV flux of the H2 discharge lamp as well as to Hans-Joachim Knóllker Dresden) and Klaus Mülllen (MPI for Polymer Research(TU Mainz) for providing the DBR and HBC samples, respectively."," The authors are grateful to Harald Mutschke and Kamel A. Khalil Gadallah for measuring the UV flux of the $_2$ discharge lamp as well as to Hans-Joachim Knöllker (TU Dresden) and Klaus Mülllen (MPI for Polymer Research Mainz) for providing the DBR and HBC samples, respectively."
+ M.S. thanks Gaéll Rouillé for his constant support through helpful discussions., M.S. thanks Gaëll Rouillé for his constant support through helpful discussions.
+ We also thank the anonymous referee for suggestions that helped to improve the quality of the, We also thank the anonymous referee for suggestions that helped to improve the quality of the
+selected |classical pulsars. B1929+10 (Beckeretal.2006).. B0950408 (Zavlin&Pavlov2004:Zharikovetal. 2004). B1133+16 (Zharikovetal. 2008)... and J0108-1431 (Mignantetal.2003) only have upper limits on the thermal emission from the whole NS surface.,"selected 	classical pulsars, B1929+10 \citep{beck06}, B0950+08 \citep{zav04,zhar04}, , B1133+16 \citep{zhar08}, , 	and J0108-1431 \citep{mig03} only have upper limits on the thermal emission from the whole NS surface."
+ The upper panel of Fig., 	The upper panel of Fig.
+ 4. shows the thermal evolution of MSPs including vortex creep and rotochemical heating with modified Urca reactions., \ref{fig:teo_obs1} shows the thermal evolution of MSPs including vortex creep and rotochemical heating 	with modified Urca reactions.
+ For the vortex creep mechanism the parameter J is highly uncertain. so we adjust it to the lowest value that is consistent with the thermal emission from MSP ," For the vortex creep mechanism the parameter $J$ is highly uncertain, so we adjust it to the lowest value 	that is consistent with the thermal emission from MSP J0437-4715."
+This constrains it to J>5.5x10? erg s if vortex creep is J0437-4715.postulated as the main heating mechanism for this pulsar., 	This constrains it to $J\geq5.5\times10^{43}$ erg s if vortex creep is postulated as the main heating 	mechanism for this pulsar.
+ For rotochemical heating. the surface temperature of this MSP can be quite precisely predicted. considering its mass M=1.76M.. (Verbiestetal2008) and an interior model givenby the A18 + ou + UIX* EOS.," For rotochemical heating, the surface temperature of this MSP can be quite precisely predicted, 	considering its mass $M=1.76M_{\odot}$ \citep{ver08} and an interior model givenby the A18 + $\delta \upsilon$ + UIX* EOS."
+ This prediction is 1.7c below the observation of Kargaltsevetal.(2004)., 	This prediction is $1.7 \sigma$ below the observation of \citet{kar04}.
+ Similarly. the bottom panel shows the thermal evolution of the classical pulsars.," 	Similarly, the bottom panel shows the thermal evolution of the classical pulsars."
+ For the vortex creep mechanism. we assumed that the parameter J is universal.," For the vortex creep mechanism, we assumed 	that the parameter $J$ is universal."
+ Thus. to solve the thermal evolution in this regime. we used the previous value of J imposed by the MSP J0437-4715. resulting in a “phenomenological” temperature prediction slightly above the constraints for the pulsars B1929+10 and B0950-08.," Thus, to solve the thermal evolution in this regime, we used the previous value of $J$ imposed by the MSP J0437-4715, resulting in a “phenomenological” temperature prediction slightly above the constraints for the 	pulsars B1929+10 and B0950+08."
+ On the other hand. as we saw in Sect. ??..," On the other hand, as we saw in Sect. \ref{sec:rq},"
+ the evolution with rotochemical heating depends on the initial period. of rotation. Po. for NSs with relatively high magnetic fields.," the evolution with rotochemical heating 	depends on the initial period of rotation, $P_0$, for NSs with relatively high magnetic fields."
+ Because of this. rotochemical heating can easily accommodate substantially lower temperatures if more restrictive observations of classical pulsars are available.," Because of this, rotochemical 	heating can easily accommodate substantially lower temperatures if more restrictive observations of classical pulsars 	are available."
+ The evolution at late times (¢>10? yr) is remarkably similar for both mechanisms., 	The evolution at late times $t>10^8$ yr) is remarkably similar for both mechanisms.
+" We showed that the relation between surface temperature and age is T,e(7? for vortex creep (Sect. ?2))."," 	We showed that the relation between surface temperature and age is $T_s \propto t^{-3/8}$ for vortex creep 	(Sect. \ref{sec:vc}) ),"
+" very slightly steeper than the relation T,«17? for rotochemical heating with modified. Urca reactions (Sect. ??»", very slightly steeper than the relation $T_s \propto t^{-1/3}$ for rotochemical heating with modified 	Urca reactions (Sect. \ref{sec:rq}) )
+ in the latter stage of the thermal evolution., in the latter stage of the thermal evolution.
+ —Figure 5 shows the thermal evolution considering the fast direct Urca processes., 		 	Figure \ref{fig:teo_obs2} shows the thermal evolution considering the fast direct Urca processes.
+ For the vortex creep mechanism. we used the same excess of angular momentum J as in the previous evolutionary curves. fixed to be consistent with the observed temperature of MSP J0437-4715.," For the vortex creep mechanism, 	we used the same excess of angular momentum $J$ as in the previous evolutionary curves, fixed to be consistent with the 	observed temperature of MSP J0437-4715."
+ Because for ages 2 1O5yr this. mechanism does not depend on the Urca process type. the temperatures of the MSPs (upper panel) are the same as with modified Urea reactions.," Because for ages $\ga10^8$ yr this 	mechanism does not depend on the Urca process type, the temperatures of the MSPs (upper panel) are the same as with modified 	Urca reactions."
+ However. in classical pulsars (bottom panel) the thermal evolution is sensitive to the direct Urea processes and generates a predicted temperature slightly lower than the limit for PSR BO9SO+08.," However, in classical pulsars (bottom panel) the thermal evolution is sensitive to the direct Urca processes 	and generates a predicted temperature slightly lower than the limit for PSR B0950+08."
+ In this way. a more sensitive observation of this pulsar could discard vortex creep as the main source of the thermal emission detected from. the MSP J0437-4715.," In this way, a 	more sensitive observation of this pulsar could discard vortex creep as the main source of the thermal emission detected from 	the MSP J0437-4715."
+ On the contrary. the temperatures generated by rotochemical heating. both for classical pulsars and. MSPs. are strongly reduced if direct Urea processes are considered.," On the contrary, the temperatures generated by rotochemical heating, both for classical pulsars and MSPs, 	are strongly reduced if direct Urca processes are considered."
+ The temperature measured in J0437-4715 requires that the neutrino emission ts produced only by modified Urea reactions if rotochemical heating ts the main source of heat., The temperature measured in J0437-4715 requires that the neutrino emission is 		produced only by modified Urca reactions if rotochemical 	heating is the main source of heat.
+ Finally. we compare the excess angular momentum. J. that we estimated from the pinning energies of Tables | and 2 with the observations.," 	Finally, we compare the excess angular momentum, $J$, that we estimated from the pinning energies of Tables \ref{tab:uno} and \ref{tab:dos}
+ 	with the observations."
+ In order to do this. we numerically integrated Eq.," In order to do this, we numerically integrated Eq."
+ 9. over the inner crust of an NS with canonical mass M=Ι.Μ. and the Al8+o6v+UIX* EOS., \ref{eq:jota} over the inner crust of an NS with 	canonical mass $M=1.4M_{\odot}$ and the $\delta v$ +UIX* EOS.
+ In this way. with the pinning energy estimated by Donati&Pizzochero (2004)... we obtain J=3.8x107!eres for the Argonne interaction and J=5.9x107ergs for the Gogny interaction.," In this way, with the pinning energy estimated by \citet{don04}, 	we obtain $J=3.8\times10^{44}~\mathrm{erg~s}$ for the Argonne interaction and $J=5.9\times 10^{44}~\mathrm{erg~s}$ for the 	Gogny interaction."
+ Similarly. but with the pinning energy estimated by Avogadro (2008). J=1.2xI0?ergs with the SLy4 interaction. and J=6.8xI0? ere s with the Skm* interaction.," Similarly, but with the pinning energy estimated by \citet{avo08}, , $J=1.2\times10^{43}~\mathrm{erg~s}$ with 	the SLy4 interaction, and $J=6.8\times 10^{43}$ erg s with the Skm* interaction."
+ Comparing these results with the value J=5.5x10? erg s. which is marginally compatiblewith both the observed temperatures of MSP J0437-4715 and the upperlimits for the old classical," Comparing these results with the value $J = 5.5\times 10^{43}$ erg s, which is marginally compatiblewith both the observed temperatures of MSP J0437-4715 and the upperlimits 	for the old classical"
+Table , 
+The polarizabilities as functions of the wavelength for all different materials used in our computations are shown in Fig.[I].,The polarizabilities as functions of the wavelength for all different materials used in our computations are shown in Fig. \ref{fig:alphas}.
+ In Figs., In Figs.
+ we plot the total mass absorption coefficient for the various aggregates., \ref{fig:full spectra} we plot the total mass absorption coefficient for the various aggregates.
+ We clearly see the expected flattening of the 10m feature with increasing aggregate size.," We clearly see the expected flattening of the $\,\mu$ m feature with increasing aggregate size."
+" Also, for the more crystalline aggregates we see the spectral signature of the crystalline species appearing."," Also, for the more crystalline aggregates we see the spectral signature of the crystalline species appearing."
+ In Figs., In Figs.
+ B] and A] we plot the effective mass absorption coefficients of the amorphous silicate (Fig. B)), \ref{fig:ISM spectra} and \ref{fig:Forst spectra} we plot the effective mass absorption coefficients of the amorphous silicate (Fig. \ref{fig:ISM spectra}) )
+ and crystalline forsterite (Fig. HH) , and crystalline forsterite (Fig. \ref{fig:Forst spectra}) )
+components., components.
+ In these plots we also show the mass absorption coefficients for the cases in which the entire aggregate is composed of a single material (amorphous silicate and crystalline forsterite respectively)., In these plots we also show the mass absorption coefficients for the cases in which the entire aggregate is composed of a single material (amorphous silicate and crystalline forsterite respectively).
+ The abundance of amorphous silicate material in the aggregate is quite large., The abundance of amorphous silicate material in the aggregate is quite large.
+" Therefore, it is not very surprising that for the amorphous silicate the pure and effective mass absorption spectra are quite similar (see Fig. B))."," Therefore, it is not very surprising that for the amorphous silicate the pure and effective mass absorption spectra are quite similar (see Fig. \ref{fig:ISM spectra}) )."
+ For the forsterite component this is different., For the forsterite component this is different.
+" As can be seen from Fig. H],"," As can be seen from Fig. \ref{fig:Forst spectra},"
+ the effective mass absorption coefficients show a spectral signature typical for small grains even when the aggregate as a whole is large., the effective mass absorption coefficients show a spectral signature typical for small grains even when the aggregate as a whole is large.
+ From these results we conclude that small inclusions of materials with strong resonances can show the spectral signature of small grains even when they are embedded in a larger grain., From these results we conclude that small inclusions of materials with strong resonances can show the spectral signature of small grains even when they are embedded in a larger grain.
+ This is an important result to keep in mind when interpreting infrared spectra of environments where coagulated grains are expected to be present., This is an important result to keep in mind when interpreting infrared spectra of environments where coagulated grains are expected to be present.
+ To quantify the statement above we determined the for the various components making up the aggregate., To quantify the statement above we determined the for the various components making up the aggregate.
+ To do this we applied the following procedure., To do this we applied the following procedure.
+ The results of this procedure are shown in Fig. D]., The results of this procedure are shown in Fig. \ref{fig:sizes fit}.
+ It is clear from this figure that the forsterite and enstatite grains appear much smaller than the amorphous silicate grains., It is clear from this figure that the forsterite and enstatite grains appear much smaller than the amorphous silicate grains.
+ This is caused by the small volume fraction of the crystalline materials., This is caused by the small volume fraction of the crystalline materials.
+" Note that although the spectral of the low abundance materials is representative for small grains, the strength is reduced compared to the mass absorption coefficient of individual small grains."," Note that although the spectral of the low abundance materials is representative for small grains, the strength is reduced compared to the mass absorption coefficient of individual small grains."
+ This makes it very hard to get reliable abundance information when using single grains to analyze aggregate spectra., This makes it very hard to get reliable abundance information when using single grains to analyze aggregate spectra.
+ In the next section we will discuss an approximate method that can be used to correctly obtain abundance information from absorption or emission spectra caused by aggregate dust particles., In the next section we will discuss an approximate method that can be used to correctly obtain abundance information from absorption or emission spectra caused by aggregate dust particles.
+" In this section we outline a fast, empirical, effective medium approach to simulate the absorption spectra of the aggregates discussed above."," In this section we outline a fast, empirical, effective medium approach to simulate the absorption spectra of the aggregates discussed above."
+from that of our fiducial sample of 22 2dECGIUS mocks.,from that of our fiducial sample of 22 2dFGRS mocks.
+ We found that the error bars obtained. directly by averaging over a test ensemble of mocks are reproduced most. closely by using the scaled. fractional scatter derived from the fiducial ensemble of 22 mocks. rather than by taking the absolute error.," We found that the error bars obtained directly by averaging over a test ensemble of mocks are reproduced most closely by using the scaled fractional scatter derived from the fiducial ensemble of 22 mocks, rather than by taking the absolute error."
+ In this section we give a brief overview of the clustering of 245 galaxies in redshift space. for samples selected by luminosity and. spectral type.," In this section we give a brief overview of the clustering of 2dFGRS galaxies in redshift space, for samples selected by luminosity and spectral type."
+ First. in Section 4.1.. we give a qualitative impression of the clustering dillerences bv plotting the spatial distribution of galaxies in. volume-limited. samples.," First, in Section \ref{sec:cone}, we give a qualitative impression of the clustering differences by plotting the spatial distribution of galaxies in volume-limited samples."
+ “Phen we quantify these dillerences. by measuring the spherically averaged correlation function. e£(5).," Then we quantify these differences by measuring the spherically averaged correlation function, $\xi(s)$."
+ A more comprehensive analysis of the clustering of 2PdPGRS galaxies in redshift space will be presented. by Llawkins (2002)., A more comprehensive analysis of the clustering of 2dFGRS galaxies in redshift space will be presented by Hawkins (2002).
+ ]t is instructive to gain a visual impression of the spatial distribution. of 2dEGIUS galaxies before. interpreting the measuredcorrelation. functions., It is instructive to gain a visual impression of the spatial distribution of 2dFGRS galaxies before interpreting the measuredcorrelation functions.
+ In. Fig. 3..," In Fig. \ref{fig:coneplot},"
+" we show the spatial distribution of galaxies in two ranges of absolute magnitude: in panel (a) we show a sample of faint galaxies 180=AL,Slogh 19.0) and in panel (hb) a sample of bright galaxies 20.02A4,Slog,h3 21.0)."," we show the spatial distribution of galaxies in two ranges of absolute magnitude: in panel (a) we show a sample of faint galaxies $-18.0\,\ge\,M_{b_{\rm J}}-5\log_{10}\,h\,\ge\,-19.0$ ) and in panel (b) a sample of bright galaxies $-20.0\,\ge\,M_{b_{\rm J}}-5\log_{10}\,h\,\ge\,-21.0$ )."
+ Within each panel. carly ancl late type galaxies. as distinguished by their spectral types. are plotted: with different symbols: the positions of early-twpes are indicated by circles and the late-types are marked by stars.," Within each panel, early and late type galaxies, as distinguished by their spectral types, are plotted with different symbols; the positions of early-types are indicated by circles and the late-types are marked by stars."
+ For clarity. we show only a three degree slice in declination cut from the SGP region and we have sparsely sampled the galaxies. so that the space densities of the two spectral classes are equal.," For clarity, we show only a three degree slice in declination cut from the SGP region and we have sparsely sampled the galaxies, so that the space densities of the two spectral classes are equal."
+ In order to expand the scale of the plot. the range of redshifts shown is restricted. taking a subset of the full voblumce-Imited sample in cach case. (," In order to expand the scale of the plot, the range of redshifts shown is restricted, taking a subset of the full volume-limited sample in each case. ("
+Note also that the redshift ranges diller between the two panels.),Note also that the redshift ranges differ between the two panels.)
+ A hierarchy. of structures is readily apparent in. these plots. ranging from isolated objects. to groups of a handful of galaxies and on through to rich clusters containing over a hundred members.," A hierarchy of structures is readily apparent in these plots, ranging from isolated objects, to groups of a handful of galaxies and on through to rich clusters containing over a hundred members."
+ Lt is interesting to see how structures are traced by galaxies in the different Luminosity bins by comparing common structures between the two panels., It is interesting to see how structures are traced by galaxies in the different luminosity bins by comparing common structures between the two panels.
+ For example. the prominent. structure (possibly a supercluster of galaxies) seen atacOP and zc0.061 is clearly visible in both panels.," For example, the prominent structure (possibly a supercluster of galaxies) seen at $\alpha\,\simeq\,0^{\rm h}$ and $z\,\simeq\,0.061$ is clearly visible in both panels."
+ Phe same is true for the overdensity seen at α50315' at zz0.068.," The same is true for the overdensity seen at $\alpha\,\simeq\,03^{\rm h}15'$ at $z\,\simeq\,0.068$."
+ ‘This is the first time that a large enough survey has been available. both in terms of the volume spannecl and the number of measured. redshifts. to allow a comparison of the clustering of galaxies of dillerent. spectral types ancl luminosities in. representative volunmie-limitecl samples. without the complication of the strong radial gradient in number density seen in [lux-limited samples.," This is the first time that a large enough survey has been available, both in terms of the volume spanned and the number of measured redshifts, to allow a comparison of the clustering of galaxies of different spectral types and luminosities in representative volume-limited samples, without the complication of the strong radial gradient in number density seen in flux-limited samples."
+ ]t is apparent from a comparison of the distribution of the different spectral tvpes in Fig. 3((, It is apparent from a comparison of the distribution of the different spectral types in Fig. \ref{fig:coneplot}( (
+a). that the fain earlv-tvpe galaxies tend to be grouped. into structures on small scales whereas the faint late-tvpes are more spread out.,"a), that the faint early-type galaxies tend to be grouped into structures on small scales whereas the faint late-types are more spread out."
+ One would therefore anticipate that the earlv-tvpes shoulc have a stronger clustering amplitude than the Iate-types. an expectation that is borne out in Section 4.2..," One would therefore anticipate that the early-types should have a stronger clustering amplitude than the late-types, an expectation that is borne out in Section \ref{sec:xi.s}."
+ In Fig. 3((, In Fig. \ref{fig:coneplot}( (
+b). the distinction between the distribution of the spectral types is less apparent.,"b), the distinction between the distribution of the spectral types is less apparent."
+ This is partly due to the greater importance of projection ellects in the declination direction. as the cone extends to a greater redshift) than in Fig. 3((," This is partly due to the greater importance of projection effects in the declination direction, as the cone extends to a greater redshift than in Fig. \ref{fig:coneplot}( ("
+a).,a).
+ However. close examination of the largest structures sugeests that carly-types are more abundant. in them than fate-types. again implving a stronger clustering amplitude.," However, close examination of the largest structures suggests that early-types are more abundant in them than late-types, again implying a stronger clustering amplitude."
+ In Fie. 4..," In Fig. \ref{fig:xis},"
+ we show the spherically averaged redshift. space correlation function. €(s). as a function of luminosity and spectral type.," we show the spherically averaged redshift space correlation function, $\xi(s)$, as a function of luminosity and spectral type."
+ Results are shown for samples selected. in bins of width one magnitude. as indicated by the Legend in each panel.," Results are shown for samples selected in bins of width one magnitude, as indicated by the legend in each panel."
+ Phe top panel shows the correlation functions of all galaxies that have been assigned a spectral type. the micelle panel shows results for galaxies classified as carly-types (9<< 148) and the bottom. panel shows results for late-tvpes {η 1.4).," The top panel shows the correlation functions of all galaxies that have been assigned a spectral type, the middle panel shows results for galaxies classified as early-types $\eta<-1.4$ ) and the bottom panel shows results for late-types $\eta>-1.4$ )."
+" Note that. at. present. there are insullicient numbers of late-type galaxies to permit a reliable measurement. of the correlation function for the brightest magnitude bin 21.02A, 5log,,h22.0."," Note that, at present, there are insufficient numbers of late-type galaxies to permit a reliable measurement of the correlation function for the brightest magnitude bin $-21.0\,\ge\,M_{b_{\rm J}}-5\log_{10}\,h\,\ge\,-22.0$."
+ Several deductions can be mace immediately from Fig. 4.., Several deductions can be made immediately from Fig. \ref{fig:xis}.
+ In all cases. the redshift space correlation function is well deseribed by a power-law over a fairly limited range of scales.," In all cases, the redshift space correlation function is well described by a power-law over a fairly limited range of scales."
+ The correlation functions of earlv-tvpe. galaxies are somewhat steeper than those of late-types., The correlation functions of early-type galaxies are somewhat steeper than those of late-types.
+ However. the main dillerence is that the earlv-tvpe galaxies have a stronger clustering amplitude than the late-twpe galaxics.," However, the main difference is that the early-type galaxies have a stronger clustering amplitude than the late-type galaxies."
+ The correlation length. defined here as the pair separation [or which €(sy)= 1. varies for earlv-tvpes from [lor galaxies with absolute magnitudes: around 19.510 sy=&9+0.7h ," The correlation length, defined here as the pair separation for which $\xi(s_{0})\,=\,1$ , varies for early-types from $s_{0}\,=\,7.1\,\pm\,0.7\,$ for galaxies with absolute magnitudes around $M_{b_{\rm J}}-5\log_{10}\,h\,\sim\,-19.5$ to $s_{0}\,=\,8.9\,\pm\,0.7\,$ "
+neglected.,neglected.
+" This requires T""=F (ie. 2vy and p/p’=p/p. where unprimed quantities are the usual ones for Z=a7|. which implies that 77=a. i.e. that the length and time units scale by the same factor."," This requires $\Gamma'=\Gamma$ (i.e. $v'=v$ ) and $p'/\rho' = p/\rho$, where unprimed quantities are the usual ones for $\zeta = \alpha = \eta = 1$ , which implies that $\eta = \alpha$, i.e. that the length and time units scale by the same factor."
+ A more elegant way of deriving this is that special relativity introduces c (the speed of light in vacuum). a UDC with units of 7/7. and thus requires a/77=| in order to keep its value the same in all our family of physical systems.," A more elegant way of deriving this is that special relativity introduces $c$ (the speed of light in vacuum), a UDC with units of $l/t$, and thus requires $\alpha/\eta = 1$ in order to keep its value the same in all our family of physical systems."
+" This reduces the number of free scaling parameters to two (V,=2 since Ny=1)."," This reduces the number of free scaling parameters to two $N_f = 2$ since $N_{\rm udc}
+= 1$ )."
+" In particular. vf)=at)maubax. Pip=pip£a. and the family of physical systems corresponding to a particular simulation is given by In particular, the total energy E. and mass AM either in a particular computational cell or in the whole computational box tor volume V) scale as E'/E=M'/M c."," In particular, $x'_\mu = (t',\vec{r'}) = \alpha(t,\vec{r}) =
+\alpha x_\mu$, $\rho'/\rho = p'/p =\zeta\alpha^{-3}$, and the family of physical systems corresponding to a particular simulation is given by In particular, the total energy $E$, and mass $M$ either in a particular computational cell or in the whole computational box (or volume $V$ ) scale as $E'/E = M'/M = \zeta$ ."
+ When gravity is included. this introduces the gravitational constant G — a UDC with units of n!P47.," When gravity is included, this introduces the gravitational constant $G$ – a UDC with units of $m^{-1}l^3t^{-2}$."
+" Thus. in order for it to keep the same value in all our family of systems requires that When there is only weak or Newtonian gravity (and general relativistic effects can be neglected). and Newtonian (bulk or thermal) motions. then G is the only UDC CV. IN,= 2) and When there effects of general relativity cannot be ignored. or or Newtonian gravity with relativistic velocities (either bulk or hermal). then there are two relevant UDCs. G and ο ελule= 2). which imply Eq. (6))"," Thus, in order for it to keep the same value in all our family of systems requires that When there is only weak or Newtonian gravity (and general relativistic effects can be neglected), and Newtonian (bulk or thermal) motions, then $G$ is the only UDC $N_{\rm udc} = 1$, $N_f =
+2$ ) and When there effects of general relativity cannot be ignored, or for Newtonian gravity with relativistic velocities (either bulk or thermal), then there are two relevant UDCs, $G$ and $c$ $N_{\rm
+udc} = 2$ ), which imply Eq. \ref{eq:gravity}) )"
+ and a=η. respectively.," and $\alpha = \eta$, respectively."
+ Together this implies hat €=ajg. ie. that all three scaling coethcients are equal. the MHD equations are also based on Maxwell's equations. and thus include c as a UDC. so they require that a=η.," Together this implies that $\zeta = \alpha =
+\eta$, i.e. that all three scaling coefficients are equal, the MHD equations are also based on Maxwell's equations, and thus include $c$ as a UDC, so they require that $\alpha = \eta$."
+" This holds even in the Newtonian case. where there are two free parameters (V,= 2) describing the relevant family of physical systems corresponding to a particular simulation. according to Eq. (900)."," This holds even in the Newtonian case, where there are two free parameters $N_f
+= 2$ ) describing the relevant family of physical systems corresponding to a particular simulation, according to Eq. \ref{eq:c}) )."
+" If gravity is included. even if weak or Newtonian gravity. then this introduces a second UDC. C. resulting in only one free parameter describing the relevant family of physical systems (V,=| since Nodeude= 2). according to Eq. (9))."," If gravity is included, even if weak or Newtonian gravity, then this introduces a second UDC, $G$, resulting in only one free parameter describing the relevant family of physical systems $N_f = 1$ since $N_{\rm udc} =
+2$ ), according to Eq. \ref{eq:G,c}) )."
+ there are various types of simulations that aim to describe the motions of discrete point-like particles. under the influence of the mutual forces that they exert on each other. rather than a continuous medium that is described by hydrodynamic or MHD equations.," there are various types of simulations that aim to describe the motions of discrete point-like particles, under the influence of the mutual forces that they exert on each other, rather than a continuous medium that is described by hydrodynamic or MHD equations."
+ Here I briefly go over two important types of such simulations., Here I briefly go over two important types of such simulations.
+ The first type is particle in cell (PIC) simulations of the motions of charged particles of either positive or negative electric charge under the mutual electromagnetic forces that they exert on each other., The first type is particle in cell (PIC) simulations of the motions of charged particles of either positive or negative electric charge under the mutual electromagnetic forces that they exert on each other.
+ In this case Maxwell's equations introduce c as a UDC (implying a@=jp., In this case Maxwell's equations introduce $c$ as a UDC (implying $\alpha=\eta$ ).
+" If we do not mind that the sealing woule change the rest mass and/or electric charge of particles. then this would be the only constraint (N44=IX. implying IN,=2 anc Eq. (90)."," If we do not mind that the scaling would change the rest mass and/or electric charge of particles, then this would be the only constraint $N_{\rm udc} = 1$ ), implying $N_f = 2$ and Eq. \ref{eq:c}) )."
+" If. however. it is important for us to accurately mode a specific particle species (such as electrons/positrons) of a given universal rest mass and electric charge. then this would add two more constraints (and. altogether [IN3). thus removing al the remaining freedom in the scaling parameters (V,;=0) and implying £2«e—37|."," If, however, it is important for us to accurately model a specific particle species (such as electrons/positrons) of a given universal rest mass and electric charge, then this would add two more constraints (and altogether $N_{\rm udc} = 3$ ), thus removing all the remaining freedom in the scaling parameters $N_f = 0$ ) and implying $\zeta=\alpha=\eta=1$."
+ The second type is N-body simulations. that are often used in cosmology and stellar or planetary dynamics. where point- like masses move under their mutual gravitational forces.," The second type is $N$ -body simulations, that are often used in cosmology and stellar or planetary dynamics, where $N$ point-like masses move under their mutual gravitational forces."
+ Since gravity is the only force involved. G must obviously remain constant. implying Eq. (6)).," Since gravity is the only force involved, $G$ must obviously remain constant, implying Eq. \ref{eq:gravity}) )."
+" If there are only Newtonian gravity and velocities then there are two free parameters CV,=2. N44.=by and Eq. (7))"," If there are only Newtonian gravity and velocities then there are two free parameters $N_f = 2$ , $N_{\rm udc}
+= 1$ ) and Eq. \ref{eq:G}) )"
+ holds., holds.
+" Otherwise. if relativistic etfects cannot be neglected. then c also enters the relevant equations as a second UDC 2) resulting in only one free parameter (NV,=1). and implying Eq. (8)."," Otherwise, if relativistic effects cannot be neglected, then $c$ also enters the relevant equations as a second UDC $N_{\rm udc} = 2$ ) resulting in only one free parameter $N_f = 1$ ), and implying Eq. \ref{eq:G,c}) )."
+ In cosmological N-body simulations with Newtonian gravity and velocities. the only bona tide UDC is ο. implying Eq. ¢7))," In cosmological $N$ -body simulations with Newtonian gravity and velocities, the only bona fide UDC is $G$, implying Eq. \ref{eq:G}) )"
+" with N,=2.", with $N_f = 2$.
+ The situation is more complicated. however. since we usually want the simulations to agree with the cosmological model of our observed universe. whose parameters are reasonably well known.," The situation is more complicated, however, since we usually want the simulations to agree with the cosmological model of our observed universe, whose parameters are reasonably well known."
+ Thus. some of these cosmological parameters might be treated as UDCs. depending on the purpose of the simulations.," Thus, some of these cosmological parameters might be treated as UDCs, depending on the purpose of the simulations."
+ For example. if the Hubble constant Hy is treated as a UDC. then it would imply 7=lI. which together with G (that implies c-a Eq. [6]]).," For example, if the Hubble constant $H_0$ is treated as a UDC, then it would imply $\eta = 1$, which together with $G$ (that implies $\zeta
+=\alpha^3$; Eq. \ref{eq:gravity}] ]),"
+" results in Nj.= 2. N,=Land In particular this would leave the mass density unchanged. Puy=pypu. so that the effective Q0) that is implied by the average value of 69) over the computational box would still follow the same original assumed cosmology."," results in $N_{\rm udc} = 2$ , $N_f = 1$ and In particular this would leave the mass density unchanged, $\rho'(t')
+=\rho'(t) = \rho(t)$, so that the effective $\Omega_M(t)$ that is implied by the average value of $\langle\rho\rangle$ over the computational box would still follow the same original assumed cosmology."
+ In cosmological N-body simulations. the initial conditions are considered to be scale-invariant. since the amplitude of the initial fluctuations in the gravitational potential are (at least nearly) independent of the wavenumber Κ.," In cosmological $N$ -body simulations, the initial conditions are considered to be scale-invariant, since the amplitude of the initial fluctuations in the gravitational potential are (at least nearly) independent of the wavenumber $k$."
+ However. the corresponding fluctuations in density scale as op/pytky «K.," However, the corresponding fluctuations in density scale as $\langle\delta\rho/\rho\rangle(k)\propto k^2$ ."
+ This introduces a time dependent scale. ο)=Ὁπκι). at which the densiwe fluctuations become of order unity. (6e/p)[A;CO]=1. and thus enter the strongly non-linear stage of their evolution.," This introduces a time dependent scale, $l_1(t) =2\pi/k_1(t)$, at which the density fluctuations become of order unity, $\langle\delta\rho/\rho\rangle[k_1(t)]\equiv 1$, and thus enter the strongly non-linear stage of their evolution."
+ This scale changes under a rescaling of the length units. and so does wy or the normalization of the initial power spectrum (i.e. the length scale that c represents would no longer be 82' Mpc upon rescaling of the length. but instead ax8/! Mpc).," This scale changes under a rescaling of the length units, and so does $\sigma_8$ or the normalization of the initial power spectrum (i.e., the length scale that $\sigma_8$ represents would no longer be $8h^{-1}\;$ Mpc upon rescaling of the length, but instead $\alpha\times 8h^{-1}\;$ Mpc)."
+ Sticking to the observed value of c. and treating both Hp and oy as UDCs. would etfectively remove the last degree of freedom in our rescaling(i.e.," Sticking to the observed value of $\sigma_8$, and treating both $H_0$ and $\sigma_8$ as UDCs, would effectively remove the last degree of freedom in our rescaling(i.e."
+" result in NV,= 0).", result in $N_f = 0$ ).
+ If. however. the cosmological parameters such as Ayor ox are not treated as UDCs. and are allowed to vary (even if only over a limited range that is consistent with current observational constraints). then a rescaling ofthe units given by Eq. (7))," If, however, the cosmological parameters such as $H_0$or $\sigma_8$ are not treated as UDCs, and are allowed to vary (even if only over a limited range that is consistent with current observational constraints), then a rescaling ofthe units given by Eq. \ref{eq:G}) )"
+ could help to reduce thenumber of simulations required in order to numericallystudy a large parameter space with ditterent cosmologies., could help to reduce thenumber of simulations required in order to numericallystudy a large parameter space with different cosmologies.
+ Depending on the physics that are included in a simulation. further," Depending on the physics that are included in a simulation, further"
+ , 
+where the circular velocity is 20kms+.,where the circular velocity is $20~\kms$.
+" Lf we again accept the orbit is inclined so that all angular momentum is in the plane of the sky. the relationship for the projected position of such an orbit is simply Ape ⋅(6) whereHepes the Leyangularap fro,frequencyWhewIGI1 VD, ponet."," If we again accept the orbit is inclined so that all angular momentum is in the plane of the sky, the relationship for the projected position of such an orbit is simply ) [kpc] where the angular frequency${1 \over 100Myr}$ is $V_c D_{p,max}^{-1}$."
+" lt can"" casily be seen that for £277.5Mgr. DycLAkpe which is the maximum projected distance of the + largest GCs."," It can easily be seen that for $t>77.5~Myr$, $D_p>1.4~kpc$ which is the maximum projected distance of the 4 largest GCs."
+ What we want to know is whether observing the 4 GC's within 1.4πρὸ is à contrivance if they have a circular orbit beyond 2Ape., What we want to know is whether observing the 4 GCs within $1.4~kpc$ is a contrivance if they have a circular orbit beyond $2~kpc$.
+ Since of the time a GC on a circular orbit at 2Ape is within 1.4Ape. the probability of observing it within 14kpe is (42).," Since of the time a GC on a circular orbit at $2~kpc$ is within $1.4~kpc$, the probability of observing it within $1.4~kpc$ is $\left({77.5 \over 100}\right)$."
+ Therefore. the probability of observing all four⋅ within⋠⋠ 1.4Apc is ⋠∣⋉⋉↽⊥(557).=⋅⋅0.36. which. one could argue makes this a perfectly plausible scenario.," Therefore, the probability of observing all four within $1.4~kpc$ is $\left({77.5 \over 100}\right)^4=0.36$, which one could argue makes this a perfectly plausible scenario."
+ One final point worth mentioning is that the timing problem is aggravated. greatly. by the existence of GC 3 (the most massive one by a factor of wo)., One final point worth mentioning is that the timing problem is aggravated greatly by the existence of GC 3 (the most massive one by a factor of two).
+ It is entirely possible that GC 3 orbits at à cowafe disance from Fornax. whereas the other four are not on orbis bevond the tical radius but instead. are gradually spiralling in.," It is entirely possible that GC 3 orbits at a “safe” distance from Fornax, whereas the other four are not on orbits beyond the tidal radius but instead are gradually spiralling in."
+ Lf this is the case then from Fig 9 we know that tjo next. two largest. GC's (2 and 5) will take at least S€/grs to decay., If this is the case then from Fig \ref{fig:dfm} we know that the next two largest GCs (2 and 5) will take at least $8~Gyrs$ to decay.
+ Comparing that to the first panel of Fig 10.. to survive à Hubble time and plunge to within 1.2Ape o| Fornax the racial orbit. must be relatively extreme.," Comparing that to the first panel of Fig \ref{fig:gcs}, to survive a Hubble time and plunge to within $1.2~kpc$ of Fornax the radial orbit must be relatively extreme."
+ Recal that Fie LO was made to show the orbit which could. plunge deepest into Fornax aud still survive a Hubble time., Recall that Fig \ref{fig:gcs} was made to show the orbit which could plunge deepest into Fornax and still survive a Hubble time.
+ Lf instead the orbit only skims the edge of Fornax. the radial orbit need. not be as extreme.," If instead the orbit only skims the edge of Fornax, the radial orbit need not be as extreme."
+ Thus. different orbits exist. both circular and exotic. which can survive in orbit of Fornax at the correct distance for a IIubble time.," Thus, different orbits exist, both circular and exotic, which can survive in orbit of Fornax at the correct distance for a Hubble time."
+ The fact that 5 GCs currently. orbit. the Fornax chwarl. although it has no stellar nuceus. would seem to suggest that dynamical friction is à weak means of draining their orbital energy.," The fact that 5 GCs currently orbit the Fornax dwarf, although it has no stellar nucleus, would seem to suggest that dynamical friction is a weak means of draining their orbital energy."
+ Significantlv. t10 DM profile necessary to match the losVD profile of 116 Fornax dwarf can easily accommocdate the slow decay of the GC's to orbital clistances resembling their current oncs (Fig 8))," Significantly, the DM profile necessary to match the losVD profile of the Fornax dwarf can easily accommodate the slow decay of the GCs to orbital distances resembling their current ones (Fig \ref{fig:dfn}) )."
+ Dynamical friction insile the tidal raclius is far stronger in MOND (see Fieὃν 5))., Dynamical friction inside the tidal radius is far stronger in MOND (see Fig \ref{fig:df}) ).
+ GCs beginningoὃν on circular orbits narrowly inside the tidal racius lose their orbital energy to DF and sink to the centre of Fornax within a Hubble time., GCs beginning on circular orbits narrowly inside the tidal radius lose their orbital energy to DF and sink to the centre of Fornax within a Hubble time.
+ In the case of the largest GC. he orbit cannot be sustained for more than 5Cyrs.," In the case of the largest GC, the orbit cannot be sustained for more than $5~Gyrs$."
+ On the other hand. highly racial orbits that enter well inside the uminous part of the dwarf can survive for a Hubble time. with the final 2 or 3λCyrs spent almost entirely within the uminous part before eventually spiralling in.," On the other hand, highly radial orbits that enter well inside the luminous part of the dwarf can survive for a Hubble time, with the final 2 or $3~Gyrs$ spent almost entirely within the luminous part before eventually spiralling in."
+ Furthermore. if these GCs were captured. in the last 4 Clyrs. there is also not enough time to drain the angular momentum of even he most massive one.," Furthermore, if these GCs were captured in the last $4~Gyrs$ , there is also not enough time to drain the angular momentum of even the most massive one."
+ It must be, It must be
+Because we are relving on relative fluxes al different. Irequencies to identify Ls. ib is critical that we understand the impact of beam sizes on these flux measurements.,"Because we are relying on relative fluxes at different frequencies to identify s, it is critical that we understand the impact of beam sizes on these flux measurements."
+ In some cases (as for NGC 6946). the multilrequency observations were done in different array configurations to match beam sizes al different frequencies.," In some cases (as for NGC 6946), the multifrequency observations were done in different array configurations to match beam sizes at different frequencies."
+ However. in some observations presented here the beam sizes are not well matched (as for M33).," However, in some observations presented here the beam sizes are not well matched (as for M33)."
+ In these cases. tvpicallv a single configuration is used [or all frequencies causing tlie beam size al the hieher frequency to be smaller than the beam size at lower frequency.," In these cases, typically a single configuration is used for all frequencies causing the beam size at the higher frequency to be smaller than the beam size at lower frequency."
+ If the object (UDILLII) is point-like wilh respect to the beams. (he larger (lower Irequenev) beam will contain more non-thermal backeround contribution (as discussed above). which can only cause these objects to appearless (hermal.," If the object ) is point-like with respect to the beams, the larger (lower frequency) beam will contain more non-thermal background contribution (as discussed above), which can only cause these objects to appear thermal."
+ If the object is extended. (he lower frequeney beam will contain more flux from (he same physical environment than (he higher lrequeney beam again. artificially boosting the relative Εικ at the lower frequency. and also causing the object to appearfess thermal.," If the object is extended, the lower frequency beam will contain more flux from the same physical environment than the higher frequency beam — again, artificially boosting the relative flux at the lower frequency, and also causing the object to appear thermal."
+ In other words. if the beam size at lower frequencies is larger than the beam size at higher frequencies (as is (he case for the mis-mnatcelied beams in this paper). it will only disguise the oplically thick signature of UDILIIs and. artificially initate an optically thick thermal spectral index.," In other words, if the beam size at lower frequencies is larger than the beam size at higher frequencies (as is the case for the mis-matched beams in this paper), it will only disguise the optically thick signature of s and artificially imitate an optically thick thermal spectral index."
+ The condition that candidates must have positive spectral indices has resulted in (he detection of fourteen sources in M32. [ive sources in NGC 253. and sixteen sources in NGC 6946.," The condition that candidates must have positive spectral indices has resulted in the detection of fourteen sources in M33, five sources in NGC 253, and sixteen sources in NGC 6946."
+" These sources are listed in Tables | 3. along with their spectral indices (a. where 5S,x »)."," These sources are listed in Tables \ref{tbl-1} – \ref{tbl-3} along with their spectral indices $\alpha$, where $S_{\nu} \propto \nu^{\alpha}$ )."
+ Within the uncertainties. eight of these objects are not inconsistent with having aE0. and we cannot rule out an optically thin rregion as the source.," Within the uncertainties, eight of these objects are not inconsistent with having $\alpha \lesssim 0$, and we cannot rule out an optically thin region as the source."
+ The remaining twenty-seven objects are only consistent with having an oplically thick origin. and we consider (hese sources to be strong candidates for having an UCILIIL-like origin.," The remaining twenty-seven objects are only consistent with having an optically thick origin, and we consider these sources to be strong candidates for having an -like origin."
+ Ulvestad&Antonueci(1997). have previously discussed the brightest of (hese five sources in NGC 253 (source #33 in this paper)., \citet{ulvestad97} have previously discussed the brightest of these five sources in NGC 253 (source 3 in this paper).
+" Their derived properties to size =2.4x1.2 pe. EM—2x10 ? pe. and n,=1.3x10! 5) are in excellent agreement with the values we derive lor this object in this section."," Their derived properties $Q_{Lyc} = 5.2 \times 10^{51}$ $^{-1}$, size $=2.4 \times
+1.2$ pc, $EM = 2 \times 10^8$ $^{-6}$ pc, and $n_e = 1.3 \times
+10^4$ $^{-3}$ ) are in excellent agreement with the values we derive for this object in this section."
+ However. the properties of the other four sources were not discussed. and we wish to add these sources to (he current sample of UDILIIs in the literature.," However, the properties of the other four sources were not discussed, and we wish to add these sources to the current sample of s in the literature."
+ As the gas and cust associated. with Ls clissipates and (he extinction lessens. these objects will become more easily visible in optical light.," As the gas and dust associated with s dissipates and the extinction lessens, these objects will become more easily visible in optical light."
+ On the other hand. UDILIIs will become less obvious in radio observations as their ambient densities decrease and (heir thermal bremsstrahilung emission decreases as à result.," On the other hand, s will become less obvious in radio observations as their ambient densities decrease and their thermal bremsstrahlung emission decreases as a result."
+the data or biases that mimic Keplerian periodicity.,the data or biases that mimic Keplerian periodicity.
+" The MAP estimate and the corresponding Do99 set of the period of this fourth signal is 3120 [2560, 3940] days for Lick1 data and 3860 [3180, 5160] for Lick2 data."," The MAP estimate and the corresponding $\mathcal{D}_{0.99}$ set of the period of this fourth signal is 3120 [2560, 3940] days for Lick1 data and 3860 [3180, 5160] for Lick2 data."
+ The latter of these estimates appears to be very close to the estimate of Curieletal.(2011) of 3848.86-:0.74 days., The latter of these estimates appears to be very close to the estimate of \citet{curiel2011} of $\pm$ 0.74 days.
+" However, because of the difference of more than 700 days between the MAP estimates of the periods from Lick! and Lick2, we cannot conclude, based on the Lick data alone, that there are indeed four Keplerian signals in the data."," However, because of the difference of more than 700 days between the MAP estimates of the periods from Lick1 and Lick2, we cannot conclude, based on the Lick data alone, that there are indeed four Keplerian signals in the data."
+ This inconsistency is seen the most clearly when looking at the equiprobability contours of the parameter posterior densities given each data set., This inconsistency is seen the most clearly when looking at the equiprobability contours of the parameter posterior densities given each data set.
+" The contours containing50%,,95%,, and of the density are shown in Fig."," The contours containing, and of the density are shown in Fig."
+ 2 for the period and amplitude parameters of v And d (top) and the proposed v And e (bottom)., \ref{upsAnd_contours} for the period and amplitude parameters of $\upsilon$ And d (top) and the proposed $\upsilon$ And e (bottom).
+ The Lick1 contours are shown in red and Lick2 contours in blue., The Lick1 contours are shown in red and Lick2 contours in blue.
+" As seen in this Fig.,"," As seen in this Fig.,"
+ the estimated period and amplitude of the v And d differ also between the two Lick data sets., the estimated period and amplitude of the $\upsilon$ And d differ also between the two Lick data sets.
+" Because of the inconsistency of the Lick data sets published in Fischeretal.(2003) and Wrightetal.(2009), we use the model inadequacy criterion to find out if either of these two data sets is also inconsistent with the combined ELODIE, AFOE, HET, and HJS data."," Because of the inconsistency of the Lick data sets published in \citet{fischer2003} and \citet{wright2009}, we use the model inadequacy criterion to find out if either of these two data sets is also inconsistent with the combined ELODIE, AFOE, HET, and HJS data."
+" We denote this combined data as m and use m, and m» todenote the Lickl and Lick2 data, respectively, and calculate the Bayes factors B(m,m) and B(m,mz) for the model M4."," We denote this combined data as $m$ and use $m_{1}$ and $m_{2}$ todenote the Lick1 and Lick2 data, respectively, and calculate the Bayes factors $B(m,m_{1})$ and $B(m,m_{2})$ for the model $\mathcal{M}_{I,4}$."
+" The logarithms of these factors are 4.01 and -10.20, respectively (Table 6))."," The logarithms of these factors are 4.01 and -10.20, respectively (Table \ref{dataset_inconsistency}) )."
+" This implies that the Lick2 data set is inconsistent with the rest of the data and the 4-companion model is an inadequate description with a probability of more than 0.999, whereas the Lickl data cannot be shown inconsistent with the rest of the data with a probability exceeding5%."," This implies that the Lick2 data set is inconsistent with the rest of the data and the 4-companion model is an inadequate description with a probability of more than 0.999, whereas the Lick1 data cannot be shown inconsistent with the rest of the data with a probability exceeding."
+". Therefore, it appears that Lick1 data (Fischeretal.,2003) is consistent with the other four data sets but the Lick2 data (Wrightetal.,2009) is not."," Therefore, it appears that Lick1 data \citep{fischer2003} is consistent with the other four data sets but the Lick2 data \citep{wright2009} is not."
+" We also investigated whether some of the ELODIE, AFOE, HET, HJS, and Lick data sets were inconsistent with the rest of the data by calculating the Bayes factors B(m;,m), where 1,..., 5, refers to each of these sets, respectively, and m contains all the data except the set m;."," We also investigated whether some of the ELODIE, AFOE, HET, HJS, and Lick data sets were inconsistent with the rest of the data by calculating the Bayes factors $B(m_{i}, m)$, where $m_{i}, i=1, ..., 5$ , refers to each of these sets, respectively, and $m$ contains all the data except the set $m_{i}$."
+ We performed these calculations using both Lick! data and Lick2 data., We performed these calculations using both Lick1 data and Lick2 data.
+ The probabilities of the model M4 being inadequate in describing each of these sets with respect the the rest of the data are shown in Table 6..," The probabilities of the model $\mathcal{M}_{I,4}$ being inadequate in describing each of these sets with respect the the rest of the data are shown in Table \ref{dataset_inconsistency}."
+" The results in Table 6 show that while Lick2 data is inconsistent with the rest of the measurements with respect to the model M4, the AFOE data is also inconsistent with the rest of the measurements regardless of using the Lick1 or"," The results in Table \ref{dataset_inconsistency} show that while Lick2 data is inconsistent with the rest of the measurements with respect to the model $\mathcal{M}_{I,4}$ , the AFOE data is also inconsistent with the rest of the measurements regardless of using the Lick1 or"
+developed by Fregeauetal.(2004).,developed by \citet{FCZR2004}.
+ The stellar and binary internal evolution was done according to BSE code (Hurley.Pols.&Tout2000:Hurley.Tout.&Pols 2002).," The stellar and binary internal evolution was done according to BSE code \citep[]{Hu2000,Hu2002}."
+. The description of the escape process was based on the theory presented by Fokushigezie (2000)., The description of the escape process was based on the theory presented by \citet{FH2000}.
+". So. the MOCCA code is able to follow all channels of interaction up to binary-binary hierarchy including merging of stars. and the escape is not immediate any more. stars need time to find their way to escape from the £L, and £L» Lagrangian points."," So, the MOCCA code is able to follow all channels of interaction up to binary-binary hierarchy including merging of stars, and the escape is not immediate any more, stars need time to find their way to escape from the $L_1$ and $L_2$ Lagrangian points."
+ The probability of escape and probability for interaction are characterized by some free parameters which were adjusted by comparison of MOCCA and N-body simulation results for large iN systems. up to WV200000.," The probability of escape and probability for interaction are characterized by some free parameters which were adjusted by comparison of MOCCA and N-body simulation results for large $N$ systems, up to $N=200000$."
+ It was shown that the free parameters of the MOCCA code can be successfully calibrated against N-body simulations and that the free parameters do not depend on ΑΝ., It was shown that the free parameters of the MOCCA code can be successfully calibrated against N-body simulations and that the free parameters do not depend on $N$.
+ MOCCA code is not only able to follow evolution of the cluster total mass. Lagrangian radii and the core radius. but also is able to reproduce in an reasonably accuracy distributions of binary parameters and number of BH-BH binaries and BSS.," MOCCA code is not only able to follow evolution of the cluster total mass, Lagrangian radii and the core radius, but also is able to reproduce in an reasonably accuracy distributions of binary parameters and number of BH-BH binaries and BSS."
+ It also reproduces well the results obtained by Baumgardt(2001). for single mass tidally limited systems for the half-mass time and evolution of potential escapers., It also reproduces well the results obtained by \citet{Ba2001} for single mass tidally limited systems for the half-mass time and evolution of potential escapers.
+ The code is able to cope with very diverse systems. from single mass. isolated system without primordial binaries to multi mass. idally limited system with large fraction of binaries.," The code is able to cope with very diverse systems, from single mass, isolated system without primordial binaries to multi mass, tidally limited system with large fraction of binaries."
+ It was shown hat simplified and faster version of the MOCCA code (without direct FewBody integrator - MOCCA-NOoFB) is a method of choice or projects which aim is to investigate the evolution. of star clusters from the point of view their global properties., It was shown that simplified and faster version of the MOCCA code (without direct FewBody integrator - MOCCA-NoFB) is a method of choice for projects which aim is to investigate the evolution of star clusters from the point of view their global properties.
+" For other »urposes. particularly when properties of ""peculiar"" objects and heir distributions are in area of interests one should use slower TOCCA code."," For other purposes, particularly when properties of ""peculiar"" objects and their distributions are in area of interests one should use slower MOCCA code."
+" It is worth to note that MOCCA and MOCCA-oFB simulations presented in this paper needs only about three and two hours, respectively. to be completed on a standard Opetron 2.4Ghz CPU."," It is worth to note that MOCCA and MOCCA-NoFB simulations presented in this paper needs only about three and two hours, respectively, to be completed on a standard Opetron 2.4Ghz CPU."
+ Despite these successes the MOCA code has still some known shortcomings. which we summarise here.," Despite these successes the MOCA code has still some known shortcomings, which we summarise here."
+ Despite these limitations. some of which are difficult to cure. the MOCCA code presented in this paper shows its potential power in simulations of real star clusters. from open clusters to rich globular clusters.," Despite these limitations, some of which are difficult to cure, the MOCCA code presented in this paper shows its potential power in simulations of real star clusters, from open clusters to rich globular clusters."
+ Monte Carlo models are feasible in a reasonable time (a few days) for globular clusters. which are. still for some time. too large for direct N-body simulations.," Monte Carlo models are feasible in a reasonable time (a few days) for globular clusters, which are, still for some time, too large for direct N-body simulations."
+ The MOCA code is able to provided data as detailed as N-body code can do., The MOCA code is able to provided data as detailed as N-body code can do.
+ Only those two codes can provides such comprehensive information., Only those two codes can provides such comprehensive information.
+ Even when ;N-body simulations eventually become possible. Monte Carlo models will remain as a quicker way of exploring the parameter space for the large scale A’-body simulations.," Even when $N$ -body simulations eventually become possible, Monte Carlo models will remain as a quicker way of exploring the parameter space for the large scale $N$ -body simulations."
+ The authors are grateful to John Fregeau for making his FewBody code accessible to the public and for very helpful suggestions related to efficient use of the code., The authors are grateful to John Fregeau for making his FewBody code accessible to the public and for very helpful suggestions related to efficient use of the code.
+ MG was supported by the Polish Ministry of Sciences and Higher Education through the grants. 92/N-ASTROSIM/2008/0 and N N203 38036., MG was supported by the Polish Ministry of Sciences and Higher Education through the grants 92/N-ASTROSIM/2008/0 and N N203 38036.
+ He warmly thanks Douglas Heggie for his hospitality during a visit to Edinburgh University., He warmly thanks Douglas Heggie for his hospitality during a visit to Edinburgh University.
+ DCH thanks MG for his kind hospitality during recent visit to CAMK., DCH thanks MG for his kind hospitality during recent visit to CAMK.
+"oon linear scales, but on translinear scales, to exploit the reduced nonlinearity in the shape and covariance of","on linear scales, but on translinear scales, to exploit the reduced nonlinearity in the shape and covariance of."
+" In principle, this could be done by setting the variancePes).. of the Gaussian in sso that the large-scale amplitudeP; of mmatches that ofP5,, or equivalently P;multiplying bby a factor to line it up with oon linear scales"," In principle, this could be done by setting the variance of the Gaussian in so that the large-scale amplitude of matches that of, or equivalently multiplying by a factor to line it up with on linear scales."
+" We experimented with this, estimating this factor by averaging buttheuncertaintyinthis f actohpedeeclsnedngob lrianceonsmallscales, comparabletothato P.."," We experimented with this, estimating this factor by averaging $(k)/$ $(k)$ over a range of $k$, but the uncertainty in this factor produced strong covariance on small scales, comparable to that of."
+" In fact, P5(k)/Pcys)these experiments (k)overarangeofk,partially motivated the form of the ccovariance matrix found in ?.."," In fact, these experiments partially motivated the form of the covariance matrix found in \citet{n11a}."
+" reffig:ellipses.ombinedshowstheresults f romanotherpossibility aand ttogether, Py1+6)and aand ttogether."," \\ref{fig:ellipses_combined} shows the results from another possibility, in which we analyze and together, and and together."
+" PasGenerally, the constraints from analyzing two power spectra together are simply the minimum at each kmax of the results from analyzing each individually."," Generally, the constraints from analyzing two power spectra together are simply the minimum at each $\kmax$ of the results from analyzing each individually."
+" This is unsurprising given the high, but not total, degeneracy between the power spectra."," This is unsurprising given the high, but not total, degeneracy between the power spectra."
+" However, for the combination of aandPg(5), there are significant gains over analyzing each individually for some parameter combinations, but the constraints are never better than for(1,5)."," However, for the combination of and, there are significant gains over analyzing each individually for some parameter combinations, but the constraints are never better than for."
+" We have explored the sensitivity to cosmological parameters of the power spectra of various transformations of the overdensity field.Ps5,"," We have explored the sensitivity to cosmological parameters of the power spectra of various transformations of the overdensity field.,"
+", the conventional power spectrum, benefits from being exactly the linear power spectrum on linear scales."," the conventional power spectrum, benefits from being exactly the linear power spectrum on linear scales."
+ iis in the simple effects of smoothing on it., Another benefit of is in the simple effects of smoothing on it.
+" However, on translinear scales it suffers strong nonlinearities, both in the mean shape and in the covariance, degrading parameter constraints."," However, on translinear scales it suffers strong nonlinearities, both in the mean shape and in the covariance, degrading parameter constraints.,"
+" spectrum of the log-density, has the most RnuhejhudimpbysePcosmology-constraining power of any power spectrum considered here."," the power spectrum of the log-density, has the most cosmology-constraining power of any power spectrum considered here."
+" Typically pushing to the smallest scales analyzed here, constraints in marginalized and unmarginalized error bars are a factor of 2-3 smaller than forP;."," Typically, pushing to the smallest scales analyzed here, constraints in marginalized and unmarginalized error bars are a factor of 2-3 smaller than for."
+. The generality of this result suggests that it would hold for other cosmological parameters as well., The generality of this result suggests that it would hold for other cosmological parameters as well.
+" This improvement over ccomes from the high diagonality of Pj,(1.,5)'ss covariance matrix, and from the small departures from the shape of the linear power spectrum."," This improvement over comes from the high diagonality of s covariance matrix, and from the small departures from the shape of the linear power spectrum."
+" In particular, the tilt n, in the linear power spectrum is dramatically better-preserved in"," In particular, the tilt $n_s$ in the linear power spectrum is dramatically better-preserved in"
+to be applied to determine the flux in regions with signal (?).,to be applied to determine the flux in regions with signal .
+. This is described in for the THINGS galaxies., This is described in for the THINGS galaxies.
+" This scaling of the residuals, and therefore the noise, yields correct fluxes in areas with signal, however the noise properties are no longer representative."," This scaling of the residuals, and therefore the noise, yields correct fluxes in areas with signal, however the noise properties are no longer representative."
+ Residual-scaled cubes should thus be only used to measure fluxes in areas with significant signal., Residual-scaled cubes should thus be only used to measure fluxes in areas with significant signal.
+ Any other applications which critically depend on noise properties (such as profile fitting) should use the unscaled cubes discussion)., Any other applications which critically depend on noise properties (such as profile fitting) should use the unscaled cubes .
+". In order to illustrate the limitations of aand potential advantages that ccould provide, the remainder of this paper is devoted to a practical comparison between aand ((Section 3))."," In order to illustrate the limitations of and potential advantages that could provide, the remainder of this paper is devoted to a practical comparison between and (Section \ref{sec:clean-msclean}) )."
+ We also compare the of iin the CASA and AIPS software packagesimplementation to classical wwith the use of wwindows (Section 4))., We also compare the implementation of in the CASA and AIPS software packages to classical with the use of windows (Section \ref{sec:ms-mr-w-c}) ).
+ We use a small sample of galaxies from THINGS to compare the results., We use a small sample of galaxies from THINGS to compare the results.
+ We chose three galaxies from the THINGS survey to use for our comparison of aandCLEAN., We chose three galaxies from the THINGS survey to use for our comparison of and.
+". These galaxies are2403,, and and they cover a range of mmasses and morphologies(?)."," These galaxies are, and and they cover a range of masses and morphologies."
+". Each galaxy was observed with the VLA in B, C and D configurations."," Each galaxy was observed with the VLA in B, C and D configurations."
+ This study uses the calibrated and combined data-set from all arrays for each galaxy., This study uses the calibrated and combined data-set from all arrays for each galaxy.
+ More details on the observations and generation of the data-sets can be found in?., More details on the observations and generation of the data-sets can be found in.
+. The iimplementation in the CASA software package and the iimplementation in the AIPS software package was used on the data-set for this comparison., The implementation in the CASA software package and the implementation in the AIPS software package was used on the data-set for this comparison.
+" For each galaxy, two data cubes were generated with two different weightings; a ‘natural’ weighting and a ‘robust’ weighting with a robustness parameter of 0.5(?)."," For each galaxy, two data cubes were generated with two different weightings; a `natural' weighting and a `robust' weighting with a robustness parameter of $0.5$."
+. These weightings were the same as used in the creation of the original THINGS data cubes., These weightings were the same as used in the creation of the original THINGS data cubes.
+ The data cubes were created with the same spatial and velocity pixel sizes as the original THINGS cubes?)., The data cubes were created with the same spatial and velocity pixel sizes as the original THINGS cubes.
+. The cubes were then *MSCLEAN-ed' down to 2.50 which was also the flux threshold used in THINGS., The cubes were then -ed' down to $2.5\sigma$ which was also the flux threshold used in THINGS.
+" As mentioned in Section 2.3,, ccan usually use much higher gain factors than classicalCLEAN."," As mentioned in Section \ref{sec:clean-msclean-issues}, can usually use much higher gain factors than classical."
+". In this study, the ggain factor was set at 0.7 for all three galaxiesgain)."," In this study, the gain factor was set at $0.7$ for all three galaxies."
+. No wwindows were used., No windows were used.
+" A maximum number of literations of 1200 (NGC 2403), 700 (Holmberg II) and 1000 (IC 2574) were chosen for each galaxy per channel."," A maximum number of iterations of $1200$ (NGC 2403), $700$ (Holmberg II) and $1000$ (IC 2574) were chosen for each galaxy per channel."
+ In all cases the flux limit was reached well before the iteration limit., In all cases the flux limit was reached well before the iteration limit.
+" Unless mentioned otherwise below, the cubes were also corrected for primary beam attenuation using the task of the software package."," Unless mentioned otherwise below, the cubes were also corrected for primary beam attenuation using the task of the software package."
+ Residual cubes were also created for the aand ddata., Residual cubes were also created for the and data.
+" As discussed briefly in Section 2.2,, CASA (and AIPS++) use paraboloids for the shape of the scale components."," As discussed briefly in Section \ref{sec:msclean-proc}, CASA (and AIPS++) use paraboloids for the shape of the scale components."
+ Extensive experimentation shows that the results did not depend significantly on the number of scales chosen nor their distribution., Extensive experimentation shows that the results did not depend significantly on the number of scales chosen nor their distribution.
+ The most important choice is that of the largest size in the distribution., The most important choice is that of the largest size in the distribution.
+ We found that optimum results were obtained when this was chosen to correspond roughly to the size of the coherent structures visible in individual channel maps., We found that optimum results were obtained when this was chosen to correspond roughly to the size of the largest coherent structures visible in individual channel maps.
+ This choice does not have to be exact., This choice does not have to be exact.
+ We chose total of six scales largestdistributed where each scale is three times larger than the preceding scale., We chose a total of six scales distributed where each scale is three times larger than the preceding scale.
+" Again, this choice was not critical,a but we found that this distribution was more efficient (in terms of number of iterations required) than a linear distribution."," Again, this choice was not critical, but we found that this distribution was more efficient (in terms of number of iterations required) than a linear distribution."
+" The largest scale for each galaxy was 130"" (NGC 2403), 270"" (Holmberg II) and 160"" (IC 2574)."," The largest scale for each galaxy was $130\arcsec$ (NGC 2403), $270\arcsec$ (Holmberg II) and $160\arcsec$ (IC 2574)."
+" For comparison, Figure 1. shows how the range of scales chosen and structure size for a single channel map are related in the galaxy NGC 2403."," For comparison, Figure \ref{fig:ngc2403-scale2struct-comp} shows how the range of scales chosen and structure size for a single channel map are related in the galaxy NGC 2403."
+" It can be seen that for the largest scale size, the largest coherent structure visible in the channel map fits within its diameter."," It can be seen that for the largest scale size, the largest coherent structure visible in the channel map fits within its diameter."
+" Choosing an even larger scale would have no effect, as the largest structure is already optimally contained within the scale distribution shown, and wwould simply not choose these even larger scales."," Choosing an even larger scale would have no effect, as the largest structure is already optimally contained within the scale distribution shown, and would simply not choose these even larger scales."
+" Choosing a smaller, largest scale would simply increase the number of iterations."," Choosing a smaller, largest scale would simply increase the number of iterations."
+ Figure 2 shows the choice of scales the algorithm made with iteration number for a single velocity channel of Holmberg II., Figure \ref{fig:msclean-scalechoice} shows the choice of scales the algorithm made with iteration number for a single velocity channel of Holmberg II.
+" In other words, the figure shows the sscale component used as the algorithm progressed to lower flux levels."," In other words, the figure shows the scale component used as the algorithm progressed to lower flux levels."
+ A similar trend is observed across all channels of all three galaxies., A similar trend is observed across all channels of all three galaxies.
+ It shows that iindeed removes emission at larger scales before smaller scales (as explained in Section 2.3))., It shows that indeed removes emission at larger scales before smaller scales (as explained in Section \ref{sec:clean-msclean-issues}) ).
+ In Table | we compare the sizes of the restoring beams of the AIPS aand the CASA/AIPS++MSCLEAN., In Table \ref{tab:beam-size-comp} we compare the sizes of the restoring beams of the AIPS and the CASA/AIPS++.
+. These are both determined by a Gaussian function to the central component of the respective dirty beams., These are both determined by fitting a Gaussian function to the central component of the respective dirty beams.
+" As their dirty beams are identical, the small fittingdifferences seen in Table | are due to the different fitting procedures for a Gaussian function used in the restore processes of both software packages."," As their dirty beams are identical, the small differences seen in Table \ref{tab:beam-size-comp} are due to the different fitting procedures for a Gaussian function used in the restore processes of both software packages."
+Thus. with these parameters now identified. a solution of Eq. (1)),"Thus, with these parameters now identified, a solution of Eq. \ref{eq:1}) )"
+ is obtained that is continuous through the critical points with a tolerance for testing continuity of about 10.7. about an order of magnitude greater than the integration step size.," is obtained that is continuous through the critical points with a tolerance for testing continuity of about $10^{-8}$, about an order of magnitude greater than the integration step size."
+ Phe polvtropic exponent is varied until a suitable set of values for the parameters [rootsrgnp.tof. are obtained so that the solution would be first continuous through all the eritical points ancl second. result in a temperature of about TOO Ix. in the vicinity of 3074.," The polytropic exponent is varied until a suitable set of values for the parameters $\{r_s,u_s,r_f,u_f,u_0\}$, are obtained so that the solution would be first continuous through all the critical points and second, result in a temperature of about $700~$ K in the vicinity of $30R_0$."
+ Observations of the circumstellar atmosphere of Detelgeuse indicate that dust. primarily exists in the form ofa shell at a distance of about ~BORG (c.g.2).., Observations of the circumstellar atmosphere of Betelgeuse indicate that dust primarily exists in the form of a shell at a distance of about $\sim 30 R_0$ \citep[e.g.][]{Bloemhof1984}.
+ However. recent observations as cliscussed above. seem to indicate that here may be small amounts of alumina (ALLO) present rather close to the photosphere (see??).. within 1.5/4.," However, recent observations as discussed above, seem to indicate that there may be small amounts of alumina $_2$ $_3$ ) present rather close to the photosphere \citep[see][]{Perrin2007,Verhoelst2006}, within $1.5 R_0$."
+" l]lowever. it is nought that this alumina may be either ransient. perhaps even destroved at around. 1542, or even ""urther out in the chromosphere (see2).. or is present in such small quantities as to be unable to support a dust-driven wind at these distances."," However, it is thought that this alumina may be either transient, perhaps even destroyed at around $1.5R_0$ or even further out in the chromosphere \citep[see][]{Verhoelst2006}, or is present in such small quantities as to be unable to support a dust-driven wind at these distances."
+ There is a third possibility. rowever remote. that the small amount of alumina. present does result in a mild dust-driven. wind. but the alumina is ransparent until it accumulates silicates on its surface 7).. which occurs in the dust shell at around 302%).," There is a third possibility, however remote, that the small amount of alumina present does result in a mild dust-driven wind, but the alumina is transparent until it accumulates silicates on its surface \citep[see][]{Onaka1989}, which occurs in the dust shell at around $30 R_0$."
+" Regardless. observations do not indicate a significant oesence of dust in the region 1.5/2,2orz30/2."," Regardless, observations do not indicate a significant presence of dust in the region $1.5 R_0 \stackrel{_<}{_\sim} r \stackrel{_<}{_\sim} 30 R_0$."
+ In this paper. we address each of these possibilities. within the ramework of the hybrid-dust-driven wind model.," In this paper, we address each of these possibilities, within the framework of the hybrid-dust-driven wind model."
+ Physically. here are four distinct scenarios that emerge and these are isted in Table 2.," Physically, there are four distinct scenarios that emerge and these are listed in Table 2."
+ Each of these scenarios are explored. in the following discussion., Each of these scenarios are explored in the following discussion.
+ This is the simplest anc perhaps the most likely scenario to fit the observations of Detelgeuse's atmosphere., This is the simplest and perhaps the most likely scenario to fit the observations of Betelgeuse's atmosphere.
+ Llercin. there is no dust. formation at close distances ancl silicate dust condenses at a distance of about 3025 where the temperature drops to about TOO Wk. Stellar material is transported [rom the photosphere to this distance of 2011 by means of a Weber-Davis magneto-rotaional wind.," Herein, there is no dust formation at close distances and silicate dust condenses at a distance of about $30 R_0$ where the temperature drops to about $700~$ K. Stellar material is transported from the photosphere to this distance of $30 R_0$ by means of a Weber-Davis magneto-rotational wind."
+" This is shown in Figure 1.. with two cases corresponcding to silicate cust formation in reasonable ancl large amounts at a distance of about 30/2,."," This is shown in Figure \ref{fig:figure1}, with two cases corresponding to silicate dust formation in reasonable and large amounts at a distance of about $30R_0$."
+ The eas velocities of the critical wind solutions of Eq. (1)), The gas velocities of the critical hybrid-wind solutions of Eq. \ref{eq:1}) )
+ are shown in Figure 1. as the red and green lone-dashed lines., are shown in Figure \ref{fig:figure1} as the red and green long-dashed lines.
+ These solutions start at the base of the wind subsonic and accelerate first through the sonic point at about 5.27Ry: after this the wind. is supersonic., These solutions start at the base of the wind subsonic and accelerate first through the sonic point at about $5.27R_0$; after this the wind is supersonic.
+ It can be readily seen that the Mach numbers are small for the critical solution close to the photosphere (ri 51)., It can be readily seen that the Mach numbers are small for the critical solution close to the photosphere $r\stackrel{_<}{_\sim}5R_0$ ).
+ Bevond the sonic point. the wind mildly. accelerates through the racial and fast Alfvénn points. that are nearly coincident upon one another. and. emerges super-Alfvénnic ab large distances (r 25/4).," Beyond the sonic point, the wind mildly accelerates through the radial and fast Alfvénn points, that are nearly coincident upon one another, and emerges super-Alfvénnic at large distances $r > 25 R_0$ )."
+ X little further out. dust condenses out from the gas at a radial distance of ry=304. shown bv the vertical dashed. line marked. 77/7.," A little further out, dust condenses out from the gas at a radial distance of $r_d = 30 R_0$, shown by the vertical dashed line marked $r_d$ ""."
+ t this distance from the surface of Detelgeuse. the temperature has dropped: below the cust condensation temperature for silicates of about zTOO Ix: the temperature is given by the blue solid line ancl should. be interpreted. using the right-hand axis.," At this distance from the surface of Betelgeuse, the temperature has dropped below the dust condensation temperature for silicates of about $\approx 700$ K; the temperature is given by the blue solid line and should be interpreted using the right-hand axis."
+ The temperature profile shown in Figure 1.. corresponds to Scenario la. wherein silicate dust. [orms in reasonable amounts at around. 30725.," The temperature profile shown in Figure \ref{fig:figure1}, corresponds to Scenario 1a, wherein silicate dust forms in reasonable amounts at around $30R_0$."
+ hereafter the wine is a coupled MlIID-dust-driven wind and the gas outllow rapidly approaches the terminal velocity., Thereafter the wind is a coupled MHD-dust-driven wind and the gas outflow rapidly approaches the terminal velocity.
+ Fhus. the solution is purcly WD from the phototsphere (à= Ry) to the dus condensation radius (r— ry) and thereafter it dis hybrid.," Thus, the solution is purely WD from the phototsphere $r=R_0$ ) to the dust condensation radius $r=r_d$ ) and thereafter it is hybrid."
+ This is the primary salient feature of the hybrid wind mode developed for Detelgeuse., This is the primary salient feature of the hybrid wind model developed for Betelgeuse.
+ Also shown in Figure 1. are the corresponding dus velocities for these two models in this scenario. shown by the solid lines that lie above the long-dashed ones.," Also shown in Figure \ref{fig:figure1} are the corresponding dust velocities for these two models in this scenario, shown by the solid lines that lie above the long-dashed ones."
+ It can be seen that the dust. grains are moving racially faster thanthe gas and dragging the eas along with them., It can be seen that the dust grains are moving radially faster thanthe gas and dragging the gas along with them.
+ The solid red line represents the cust velocity profile for a model with parameters Py=0.5 and (0j= 1/2000: Scenario la., The solid red line represents the dust velocity profile for a model with parameters $\Gamma_d=0.5$ and $\langle \delta \rangle = 1/2000$ ; Scenario 1a.
+ In order to investigate the ellect of changing the average, In order to investigate the effect of changing the average
+"With reasonable assumptions for the average photospheric velocities of SNe Ib. le and Ie-BL. we may break the degeneracy between M,; and Εκ implied by our mean 7, estimates derived from our Άρης measurements.","With reasonable assumptions for the average photospheric velocities of SNe Ib, Ic and Ic-BL, we may break the degeneracy between $M_{\rm ej}$ and $E_K$ implied by our mean $\tau_c$ estimates derived from our $\Delta m_{15}$ measurements."
+" Within the framework of the analytic light curve models. these two physical parameters depend on vy, as follows: For SNe Ib anc Ie we adopt v,7:10.000kms! at maximum light. in line with spectroscopic observations (Tathesonetal.. 2001)."," Within the framework of the analytic light curve models, these two physical parameters depend on $v_{\rm
+ph}$ as follows: For SNe Ib and Ic we adopt $v_{\rm ph}\approx 10,000~\rm km~s^{-1}$ at maximum light, in line with spectroscopic observations \citep{mfl+01}."
+". For SNe Ic-BL and ergine-driven Se. the observed range of vy, values is broader: we adopt vy,=20.000kms! for the SNe Ie-BL at maximum light (e.g.. Pianetαἱ. 2006))."," For SNe Ic-BL and engine-driven SNe, the observed range of $v_{\rm
+ph}$ values is broader; we adopt $v_{\rm ph}=20,000~\rm km~s^{-1}$ for the SNe Ic-BL at maximum light (e.g., \citealt{pmm+06}) )."
+" Combining these assumptions with our average 7, estimates for the combined sample (Table 6)). we find typical values of Mj2M. and Exs;~| erg for SNe Ib and Ic while for SNe Ic-BL we find higher values. M;z5M.. and Eys,~10 erg (Figure 24))."," Combining these assumptions with our average $\tau_c$ estimates for the combined sample (Table \ref{tab:params}) ), we find typical values of $M_{\rm
+ej}\approx 2~M_{\odot}$ and $E_{K,51}\approx 1$ erg for SNe Ib and Ic while for SNe Ic-BL we find higher values, $M_{\rm ej}\approx
+5~M_{\odot}$ and $E_{K,51}\approx 10$ erg (Figure \ref{fig:me}) )."
+" For the three local engine-driven explosions we estimate M;~4M... and Ey,~9 erg (Table 85)."," For the three local engine-driven explosions we estimate $M_{\rm
+ej}\approx 4~M_{\odot}$ and $E_{K,51}\approx 9$ erg (Table \ref{tab:peak}) )."
+ Thus. based on these reasonable estimates of vy. we find no evidence for different explosion parameters among SNe Ib and Ic. while those of SNe Ic-BL are distinct and more closely resemble the values inferred for engine-driven explosions.," Thus, based on these reasonable estimates of $v_{\rm ph}$, we find no evidence for different explosion parameters among SNe Ib and Ic, while those of SNe Ic-BL are distinct and more closely resemble the values inferred for engine-driven explosions."
+ Here we consider the implications for the progenitors of SNe Ibe based on their derived explosion parameters together with host galaxy and/or explosion site diagnostics., Here we consider the implications for the progenitors of SNe Ibc based on their derived explosion parameters together with host galaxy and/or explosion site diagnostics.
+ While we derive explosion properties for SNe Ib and le. their explosion site properties imply statistically significantdissinilarities.," While we derive explosion properties for SNe Ib and Ic, their explosion site properties imply statistically significant."
+ SNe Ic favor the more luminous and metal-rich regions of their host galaxies than SNe Ib (Kelly.Kirshner&Pahre2008:Mod-Jazetal.2010 but see Andersonetαἱ. 2010)).," SNe Ic favor the more luminous and metal-rich regions of their host galaxies than SNe Ib \citealt{kkp08,mbf+10} but see \citealt{acj+10}) )."
+ These explosion site diagnostics suggest that SNe Ic progenitors are more massive and/or younger and characterized by a slightly higher metal content than the progenitors of SNe Ib., These explosion site diagnostics suggest that SNe Ic progenitors are more massive and/or younger and characterized by a slightly higher metal content than the progenitors of SNe Ib.
+ More massive progenitors may produce explosions with larger values unless they are able to lose mass more efficiently throughout., More massive progenitors may produce explosions with larger $M_{\rm ej}$ values unless they are able to lose mass more efficiently throughout.
+Μα Since the line-driven of Wolf-Rayet stars are enhanced at higher metallicity (Mwinds Z*: Vink&deKoter 2005)). we speculate that SNe le progenitors are initially more massive than those of SNe [b but lose mass more efficiently. resulting in explosion parameters that appear indistinguishable from their helium-rich cousins.," Since the line-driven winds of Wolf-Rayet stars are enhanced at higher metallicity $\dot{M}\propto Z^{0.8}$; \citealt{vd05}) ), we speculate that SNe Ic progenitors are initially more massive than those of SNe Ib but lose mass more efficiently, resulting in explosion parameters that appear indistinguishable from their helium-rich cousins."
+ Next we consider the more extreme explosion parameters we infer for SNe Ie-BL in comparison to those of SNe Ibe., Next we consider the more extreme explosion parameters we infer for SNe Ic-BL in comparison to those of SNe Ibc.
+ Explosion site metallicities for SNe Ic-BL do not reveal strong dissimilarities from those of SNe Ib and Ie (Andersoneral.2010:Modjazetaf.2010 but see Arcavietal. 2010).," Explosion site metallicities for SNe Ic-BL do not reveal strong dissimilarities from those of SNe Ib and Ic \citealt{acj+10,mbf+10} but see \citealt{agk+10}) )."
+ However. a study of the explosion site luminosities indicate that SNe Ic-BL trace the star-forming light of their host galaxies at least as tightly as SNe Ic and perhaps as tightly as GRB- (Kelly.Kirshner&Pahre.2008).," However, a study of the explosion site luminosities indicate that SNe Ic-BL trace the star-forming light of their host galaxies at least as tightly as SNe Ic and perhaps as tightly as GRB-SNe \citep{kkp08}."
+. At the same time. the light curves of SNe Ic-BL show evidence for more powerful explosions. with somewhat larger values than ordinary SNe Ibe (Table 8)).," At the same time, the light curves of SNe Ic-BL show evidence for more powerful explosions, with somewhat larger $M_{\rm ej}$ values than ordinary SNe Ibc (Table \ref{tab:peak}) )."
+" We therefore speculateM, that SNe Ic-BL may represent the explosion of more massive progenitor stars than other SNe [be across a broad range of metallicities.", We therefore speculate that SNe Ic-BL may represent the explosion of more massive progenitor stars than other SNe Ibc across a broad range of metallicities.
+ This may also explain the existence of broad-lined SNe of Types [b and IIb that stem from metal-poor progenitors (e.g.. 22003be: Hamuyetal...2009)).," This may also explain the existence of broad-lined SNe of Types Ib and IIb that stem from metal-poor progenitors (e.g., 2003bg; \citealt{hdm+09}) )."
+ Finally we consider the relation of engine-driven SN progenitors to those of SNe Ic-BL., Finally we consider the relation of engine-driven SN progenitors to those of SNe Ic-BL.
+ Based on their optical properties (ight curves and spectra). SNe Ic-BL and engine- SNe appear statistically indistinguishable. however radio observations reveal relativistic outflows in only a small fraction (Soderbergeraf..2010).," Based on their optical properties (light curves and spectra), SNe Ic-BL and engine-driven SNe appear statistically indistinguishable, however radio observations reveal relativistic outflows in only a small fraction \citep{scp+10}."
+. Meanwhile. explosion site metallicity measurements indicate that engine-driven explosions tend to populate the low end of the metallicity distribution for SNe Ibe (Modjazeta£..2010).. although there are some notable exceptions (see Levesqueefαἰ.2010).," Meanwhile, explosion site metallicity measurements indicate that engine-driven explosions tend to populate the low end of the metallicity distribution for SNe Ibc \citep{mbf+10}, although there are some notable exceptions (see \citealt{lsf+10}) )."
+ In these cases. an additional key parameter may be at play and possibly related to the rotation. binarity. or the nature of the compact remnant of the progenitor system (black hole. neutron star or magnetar: Yoon.Woosley&Langer 2010)).," In these cases, an additional key parameter may be at play and possibly related to the rotation, binarity, or the nature of the compact remnant of the progenitor system (black hole, neutron star or magnetar; \citealt{ywl10}) )."
+ We present the first uniform and statistical sample of SNe Ibe multi-band light curves available to date., We present the first uniform and statistical sample of SNe Ibc multi-band light curves available to date.
+ We find a significant dispersion among the light curves both in. peak luminosity and decay rate., We find a significant dispersion among the light curves both in peak luminosity and decay rate.
+" Through à comparison with the existing small sample of well-observed SNe Ibe light curves from the literature. we find that a significant fraction of SNe Ibe are heavily extinguished with E(BV),z0.4 mag."," Through a comparison with the existing small sample of well-observed SNe Ibc light curves from the literature, we find that a significant fraction of SNe Ibc are heavily extinguished with $E(B-V)_{\rm host}\approx 0.4$ mag."
+" After correcting our light curves for host— galaxy extinction, we compare differential and cumulative distributions for the peak absolute magnitudes of SNe Ib. Ic. and Ic-BL."," After correcting our light curves for host galaxy extinction, we compare differential and cumulative distributions for the peak absolute magnitudes of SNe Ib, Ic, and Ic-BL."
+ We find that the peak luminosity distributions for SNe Ib and Ie are statistically indistinguishable. with 7—18 mag and are therefore not inconsistent with a (Mj;single progenitor channel for helium-poor and helium-rich events.," We find that the peak luminosity distributions for SNe Ib and Ic are statistically indistinguishable, with $\langle M_{R,\rm
+peak}\rangle \approx -18$ mag and are therefore not inconsistent with a single progenitor channel for helium-poor and helium-rich events."
+ A comparison of their early decline rates supports this hypothesis and we note there is no evidence for a correlation with peak absolute magnitude as Is seen for SNe Ia. However. we find that SNe Ic-BL are typically more luminous with (Mj)~—19 mag.," A comparison of their early decline rates supports this hypothesis and we note there is no evidence for a correlation with peak absolute magnitude as is seen for SNe Ia. However, we find that SNe Ic-BL are typically more luminous with $\langle M_R\rangle \approx -19$ mag."
+ The probability that they belong to the same set of progenitors and/or explosions as ordinary SNe Ib and Ie ts Just2%., The probability that they belong to the same set of progenitors and/or explosions as ordinary SNe Ib and Ic is just.
+. We compare these observed light curve properties with those of the three nearby engine-driven SNe discovered within the same volume and find a significant overlap with the SNe Ie-BL population., We compare these observed light curve properties with those of the three nearby engine-driven SNe discovered within the same volume and find a significant overlap with the SNe Ic-BL population.
+ The probability that engine-driven SNe are drawn from the SNe Ic-BL population is25%.. based on these optical diagnostics alone.," The probability that engine-driven SNe are drawn from the SNe Ic-BL population is, based on these optical diagnostics alone."
+ This result underscores the point that neither high luminosity nor fast photospheric velocity can be used as a robust indicator for an engine-driven explosion. and thus radio searches for a relativistic outflow are required to distinguish (e.g.. Soderbergefal. 2010)).," This result underscores the point that neither high luminosity nor fast photospheric velocity can be used as a robust indicator for an engine-driven explosion, and thus radio searches for a relativistic outflow are required to distinguish (e.g., \citealt{scp+10}) )."
+ We model the optical light curves in a systematic fashion to extract estimates of the three physical parameters. νι. and Ex and find that the results reiterate those reviewed M.above.," We model the optical light curves in a systematic fashion to extract estimates of the three physical parameters, $M_{\rm Ni}$, $M_{\rm
+ej}$, and $E_K$ and find that the results reiterate those reviewed above."
+ Namely. thereis no significant difference between SNe Ib and le and we estimate their physical parameters to be Mwi720.2M ..Mj~2M... and Exσι~|erg.," Namely, thereis no significant difference between SNe Ib and Ic and we estimate their physical parameters to be $M_{\rm Ni}\approx
+0.2~M_{\odot}$, $M_{\rm ej}\approx 2~M_{\odot}$, and $E_{K,51}\approx
+1$erg."
+ Meanwhile. SNe Ic-BL and engine-driven events are similar to one another and distinct from ordinary SNe Ib and Ic: for them we estimate more extreme physical parameters of My;zz0.5 M... Myz SM... and Ex.s;~10 erg.," Meanwhile, SNe Ic-BL and engine-driven events are similar to one another and distinct from ordinary SNe Ib and Ic; for them we estimate more extreme physical parameters of $M_{\rm Ni}\approx
+0.5~M_{\odot}$ , $M_{\rm ej}\approx 5~M_{\odot}$ , and $E_{K,51}\approx
+10$ erg."
+ The statistically significant difference between ordinary SNe Ibe and SNe Ic-BL (including engine- events) suggests that the latter share a distinct Ni, The statistically significant difference between ordinary SNe Ibc and SNe Ic-BL (including engine-driven events) suggests that the latter share a distinct $^{56}$ Ni
+Since the discovery of large sanples of close compact Dinaries (754. the formation of these systems has been intensively discussed.,"Since the discovery of large samples of close compact binaries \citep[]{kraft64-1}, the formation of these systems has been intensively discussed."
+ Based on the first rough sketches provided bv 7.. the picture that most close-conrpact-biuaryv stars form through comunon-cuvelope (CE) evolution has been established.," Based on the first rough sketches provided by \citet{paczynski76-1}, the picture that most close-compact-binary stars form through common-envelope (CE) evolution has been established."
+ Once the more lnassive star in the initial binary svsteni expauds ou the first giaut. branch (FCB) or ou the asviuptotic eiat branch (ACB). it eventually fills its Roche lobe. aud dynamically unstable mass transfer to the less massive conrponeut leads to the formation of a gaseous cuvelope around the core of the elant and the companion.," Once the more massive star in the initial binary system expands on the first giant branch (FGB) or on the asymptotic giant branch (AGB), it eventually fills its Roche lobe, and dynamically unstable mass transfer to the less massive component leads to the formation of a gaseous envelope around the core of the giant and the companion."
+ This euvelope is expelled by the dissipation of orbital energy (seee.g.2.forareceutreview)..., This envelope is expelled by the dissipation of orbital energy \citep[see e.g.][ for a recent review]{webbink08-1}.
+ CE evolution teruinates the mass growth of the core of the giaut aud therefore niass fransfer initiated when the primary was on the FOB produces post-coummon-cuvelope binaries (PCEBs) containing low-mass (Mag50.5 Afsun)) Ue-core white dwarfs (WDs). while mass trausfer starting with the primary on the AGB produces PCEBs containing more massive (Maa20.5 Afsun)) C/O-core WDs.," CE evolution terminates the mass growth of the core of the giant and therefore mass transfer initiated when the primary was on the FGB produces post-common-envelope binaries (PCEBs) containing low-mass $\Mwd\,\lappr\,0.5$ ) He-core white dwarfs (WDs), while mass transfer starting with the primary on the AGB produces PCEBs containing more massive $\Mwd\,\gappr\,0.5$ ) C/O-core WDs."
+ Despite some recent progress. sumnulatious of CE evolution still fail to follow the complete spiralin aud euvelope ejection processes (e.g.7.—andreferences therein).," Despite some recent progress, simulations of CE evolution still fail to follow the complete spiral-in and envelope ejection processes \citep[e.g.][ and references therein]{taam+ricker06-1}."
+ However. a relation between the final aud. the initial orbital parameters can be obtained from a simple paramctrized euergy equation (c.g. 2: 23).," However, a relation between the final and the initial orbital parameters can be obtained from a simple parametrized energy equation (e.g. \citealt{webbink84-1}; \citealt{iben+livio93-1}) )."
+ A laree sample of PCEBs with known stellay masses aud orbital periods covering the eutire parameter space. combined with aux relation between the binary and/or stellar parameters of PCEBs that can be identified (c.g. a relation between the orbital period aud oue of the masses). iuight therefore xovide crucial iuformation for the energv budect of CE evolution.," A large sample of PCEBs with known stellar masses and orbital periods covering the entire parameter space, combined with any relation between the binary and/or stellar parameters of PCEBs that can be identified (e.g. a relation between the orbital period and one of the masses), might therefore provide crucial information for the energy budget of CE evolution."
+ The potential of PCEDs to coustiün current nodels of CE evolution was realized early by ?7 απ iu more detail later bw. e.g. 7. ?.. aud ?..," The potential of PCEBs to constrain current models of CE evolution was realized early by \citet{ritter86-1} and in more detail later by, e.g., \citet{schreiber+gaensicke03-1}, \citet{nelemans+tout05-1}, and \citet{davisetal10-1}."
+ However. he PCEB samples analyzed by these authors were not rolmogencous and also heavily affected by observational vases. Which made amy investigation of possible relations vetween the stellar masses aud the orbital period a futile exercise.," However, the PCEB samples analyzed by these authors were not homogeneous and also heavily affected by observational biases, which made any investigation of possible relations between the stellar masses and the orbital period a futile exercise."
+ Over the past five vers; we performed the first laree-scale survey of PCEBs among white-dwarf/manu-sequeuce (WDMS) binaries ideutified by the Sloan Digital Sky Survey (SDSS.?)..," Over the past five years, we performed the first large-scale survey of PCEBs among white-dwarf/main-sequence (WDMS) binaries identified by the Sloan Digital Sky Survey \citep[SDSS,][]{abazajianetal09-1}."
+ Given that the orbital period aud both stellar masses can be measured relatively easily. the large and more homogeneous sample of SDSS PCEDs now available allows us for the first time to test for possible dependencies of the orbital period of PCEBs ou the stellar masses.," Given that the orbital period and both stellar masses can be measured relatively easily, the large and more homogeneous sample of SDSS PCEBs now available allows us for the first time to test for possible dependencies of the orbital period of PCEBs on the stellar masses."
+ In particular. we can separate the sample according to the evolutionary state of the WD progenitor at the ouset of common cuvelope evolution. which provides additional observational constraints for binary population models.," In particular, we can separate the sample according to the evolutionary state of the WD progenitor at the onset of common envelope evolution, which provides additional observational constraints for binary population models."
+coupling of the two shells by waves in the nonconvective region between them.,coupling of the two shells by waves in the nonconvective region between them.
+" This behavior seems robust; we expect it to persist, so that at core collapse this part of the star (at least) will have significant nonspherical distortion."," This behavior seems robust; we expect it to persist, so that at core collapse this part of the star (at least) will have significant nonspherical distortion."
+" Another asymmetry, density perturbations, was found by Bazannonspherical&Arnett(1998);Asida&Arnett (2000)."," Another asymmetry, nonspherical density perturbations, was found by \citet{bazan1998,asida2000}."
+". The fluctuations in density and temperature, presented in the top panel of Figure 3 as root mean square deviations from an angular mean, reach values as large as ~10% and are localized at the nonconvective region just beyond the convective boundary panel)."," The fluctuations in density and temperature, presented in the top panel of Figure \ref{f3} as root mean square deviations from an angular mean, reach values as large as $\sim$ and are localized at the nonconvective region just beyond the convective boundary (top panel)."
+" The fluctuations are coincident with regions (topwhere the buoyancy frequency, vp, is large, which can be seen in the bottom panel of Figure 3.."," The fluctuations are coincident with regions where the buoyancy frequency, $\nu_B$, is large, which can be seen in the bottom panel of Figure \ref{f3}."
+" Here, is a measure of the “stiffness” of the stratification v2(Turner1973), and is proportional to the restoring buoyancy force on perturbed stellar matter, where g is the gravity, and the term in parentheses is the difference between the fractional density gradient of the stellar structure and the fractional density changedue to a radial (adiabatic) Lagrangian displacement."," Here, $\nu_B^2$ is a measure of the “stiffness” of the stratification \citep{turner73}, and is proportional to the restoring buoyancy force on perturbed stellar matter, where $g$ is the gravity, and the term in parentheses is the difference between the fractional density gradient of the stellar structure and the fractional density changedue to a radial (adiabatic) Lagrangian displacement."
+ Regions where vg is zero are unstable to convective motions., Regions where $\nu_B$ is zero are unstable to convective motions.
+" The spikes in vg in our model are due to steep, stabilizing composition gradients which separate fuel from ash and lead to sharp gradients in density."," The spikes in $\nu_B$ in our model are due to steep, stabilizing composition gradients which separate fuel from ash and lead to sharp gradients in density."
+ Convection excites wave motions in the adjacent stable layers which give rise to the density perturbations., Convection excites wave motions in the adjacent stable layers which give rise to the density perturbations.
+" Similar internal wave phenomena can be observed in laboratory ice-water convection experiments where the largest temperature fluctuations are measured immediately above the convecting layer where the buoyancy frequency is large (Townsend1966),, highlighting the generality of this phenomenon."," Similar internal wave phenomena can be observed in laboratory ice-water convection experiments where the largest temperature fluctuations are measured immediately above the convecting layer where the buoyancy frequency is large \citep{townsend1966a}, highlighting the generality of this phenomenon."
+ The underlying stellar structure determines the set of discrete resonant modes that can be excited., The underlying stellar structure determines the set of discrete resonant modes that can be excited.
+" The narrow stable layers which bound the convective shells in our simulation, including the (truncated) core layer, are isolated enough from other wave propagation regions to act as resonating cavities."," The narrow stable layers which bound the convective shells in our simulation, including the (truncated) core layer, are isolated enough from other wave propagation regions to act as resonating cavities."
+" These modes are the deeper, interior counterparts to the modes observed in helio and asteroseismology studies of milder evolutionary stages."," These modes are the deeper, interior counterparts to the modes observed in helio and asteroseismology studies of milder evolutionary stages."
+" Each mode can be uniquely identified by its horizontal wavenumber index, |, and its oscillation frequency, w."," Each mode can be uniquely identified by its horizontal wavenumber index, $l$, and its oscillation frequency, $\omega$ ."
+ We identify excited modes in our simulation by isolating, We identify excited modes in our simulation by isolating
+for LkCa 15).,for LkCa 15).
+" In comparison. the line Iuminosities of the S(1) and 5(2) emission Chat are detected from TW Ilva in SII2 are 0.8xLO""L. and 0.5xLOOL... respectively."," In comparison, the line luminosities of the S(1) and S(2) emission that are detected from TW Hya in SH2 are $0.8 \times 10^{-6}\Lsun$ and $0.5 \times 10^{-6}\Lsun$, respectively."
+ Since the 5(2) line of TW Iva is blended with an OL! feature. we assumed an ΟΠΗ contribution that is similar to that of neighboring OIL features in estimating the flux of the S(2) line.," Since the S(2) line of TW Hya is blended with an OH feature, we assumed an OH contribution that is similar to that of neighboring OH features in estimating the flux of the S(2) line."
+ The observed SCL) luminosity is similar to the value of 0.2xLO°L. that is predicted for TW Ilva by Gorti Hollenbach (2008).," The observed S(1) luminosity is similar to the value of $0.2\times 10^{-6}\,\Lsun$ that is predicted for TW Hya by Gorti Hollenbach (2008)."
+ Interestinglv. the 1I» emission from TW Ilva is stronger than (he upper limits on the ll» flux that Bitner et ((2008) measured for TW Ilva.," Interestingly, the $\Htwo$ emission from TW Hya is stronger than the upper limits on the $\Htwo$ flux that Bitner et (2008) measured for TW Hya."
+" Whereas we measure line f[Iuxes of 13x10therestem? and 0.6x10theresbem? for the SC) and S(2) lines (5-0 and 3-0 detections. respectively, in SII2). Ditier et ((2008) provide 3-0 upper limits for the 1) and $(2) lines of 0.7xLOterestem? and 0.6x10terestem7. respectively,"," Whereas we measure line fluxes of $1.3\times 10^{-14}\,\ergperssqcm$ and $0.6\times 10^{-14}\,\ergperssqcm$ for the S(1) and S(2) lines $\sigma$ and $\sigma$ detections, respectively, in SH2), Bitner et (2008) provide $\sigma$ upper limits for the S(1) and S(2) lines of $0.7 \times 10^{-14}\,\ergperssqcm$ and $0.6 \times 10^{-14}\,\ergperssqcm$, respectively."
+" Perhaps the narrower slit (0.81"" and 0.54"" for the S1 and S2 lines. respectively) of the TEXES Ineastvements plavs a role here as well."," Perhaps the narrower slit $0.81\arcsec$ and $0.54\arcsec$ for the S1 and S2 lines, respectively) of the TEXES measurements plays a role here as well."
+ The line widths that are measured for 1» emission from voung stus. when it is detected by TEXES in more distant sources (140 ppc). are ~JOkms+. which is consistent with much of the emission arising [rom the AAU region of the disk.," The line widths that are measured for $\Htwo$ emission from young stars, when it is detected by TEXES in more distant sources $\sim 140$ pc), are $\sim 10\kms$, which is consistent with much of the emission arising from the AU region of the disk."
+ As discussed in the context of the [Nel] emission from TW Ίνα (section 3.4). much of this region of the disk is exclucled by the narrow slit in a nearby source such as TW Ilva.," As discussed in the context of the [NeII] emission from TW Hya (section 3.4), much of this region of the disk is excluded by the narrow slit in a nearby source such as TW Hya."
+ As in the case of [Nell]. near-sinnutaneous narrow and wide slit ground-based observations of the II» emission from TW Iva. taken in good seeing. would be useful in constraining the spatial extent of the emission.," As in the case of [NeII], near-simultaneous narrow and wide slit ground-based observations of the $\Htwo$ emission from TW Hya, taken in good seeing, would be useful in constraining the spatial extent of the emission."
+ We also detect molecules other (han OLI and Is in the TW Ilva spectrum (COs. ICO and possibly Cll).," We also detect molecules other than OH and $\Htwo$ in the TW Hya spectrum $\COtwo$, $\HCOp$, and possibly $\CHthree$ )."
+ The properties of these emission features will be discussed in greater detail in a future publication., The properties of these emission features will be discussed in greater detail in a future publication.
+ As awell-stuclied. nearby transition object. TW IHva has been a touchstone for understanding the origin and evolutionary state of such svstenms.," As a well-studied, nearby transition object, TW Hya has been a touchstone for understanding the origin and evolutionary state of such systems."
+ Transition objects have unusual spectral energy distributions eliaracterized bv a deficit of infrared excess at short wavelengths. a trait that is interpreted. as indicating that the disk is optically thin in the continuum. within a eiven radius yj. (Strom et 11989).," Transition objects have unusual spectral energy distributions characterized by a deficit of infrared excess at short wavelengths, a trait that is interpreted as indicating that the disk is optically thin in the continuum within a given radius $R_{\rm hole}$ (Strom et 1989)."
+" Analvses of spectral energy distributions of transition objects find fj, values of a few AU to ~50 ΛΑΔΙ: in some cases (Espaillat et 22007).", Analyses of spectral energy distributions of transition objects find $R_{\rm hole}$ values of a few AU to $\sim 50$ AU in some cases (Espaillat et 2007).
+ In comparison. the spectral energy distribution of TW Iva indicates Chat its disk is optically thin within e4 AU of the star (Calvet et 22002: Uchida et 22004).," In comparison, the spectral energy distribution of TW Hya indicates that its disk is optically thin within $\sim 4$ AU of the star (Calvet et 2002; Uchida et 2004)."
+ An inner hole in, An inner hole in
+"with a Galactic rest-frame (GRF) velocity of~700 aat its current position, making it the fastest halo object known.","with a Galactic rest-frame (GRF) velocity of$\sim$ at its current position, making it the fastest halo object known."
+ This allows a significant lower limit for the total halo mass of the Galaxy to be set., This allows a significant lower limit for the total halo mass of the Galaxy to be set.
+" The present work provides a glimpse to the detailed kinematic feasible once the European Space Agency’s investigationsGaia satellite mission (e.g.Turonetal.2005) becomes operational, at much higher precision."," The present work provides a glimpse to the detailed kinematic investigations feasible once the European Space Agency's Gaia satellite mission \citep[e.g.][]{2005ESASP.576.....T}
+ becomes operational, at much higher precision."
+ In a labelsec:tar--pmprevious paper (Tillichetal.2009) we already investigated the high-velocity tail of the Xueetal.(2008) sample and found a hyper-velocity candidate of spectral type A among the stars with highest radial velocities in the GRF., In a previous paper \citep{2009A&A...507L..37T} we already investigated the high-velocity tail of the \cite{2008ApJ...684.1143X} sample and found a hyper-velocity candidate of spectral type A among the stars with highest radial velocities in the GRF.
+" Here we study the most extreme stars approaching us, applying the same techniques."," Here we study the most extreme stars approaching us, applying the same techniques."
+ We selected all stars with GRF velocities vage«—350 ffrom the RV-based sample of Xueetal.(2008) and obtained 5 targets for which we attempted to measure proper motions., We selected all stars with GRF velocities $v_{\rm GRF}<-350$ from the RV-based sample of \cite{2008ApJ...684.1143X} and obtained 5 targets for which we attempted to measure proper motions.
+ All available independent position measurements on Schmidt plates (APM - McMahonetal.(2000); SSS - Hamblyetal. (2001))) were collected and combined with the and other available positions (CMC14 Carlsberg-Meridian-Catalog (2006); 2MASS - Cutrietal.(2003);; UKIDSS - Lawrenceetal. (2007))) for a first linear proper motion fit., All available independent position measurements on Schmidt plates (APM - \cite{2000yCat.1267....0M}; SSS - \cite{2001MNRAS.326.1279H}) ) were collected and combined with the and other available positions (CMC14 \cite{2006yCat.1304....0C}; 2MASS - \cite{2003tmc..book.....C}; UKIDSS - \cite{2007MNRAS.379.1599L}) ) for a first linear proper motion fit.
+" However, there were even more measurements of Schmidt plates, from up to 14 different epochs in case of overlapping plates of the Digitised SkySurveys?."," However, there were even more measurements of Schmidt plates, from up to 14 different epochs in case of overlapping plates of the Digitised Sky."
+. FITS images of 15 15 arcmin size were extracted from all available plates byand ESO MIDAS tools were used to measure positions., FITS images of 15 by 15 arcmin size were extracted from all available plates and ESO MIDAS tools were used to measure positions.
+" For this purpose, we selected compact background galaxies around each target, identified fromSDSS,, to transform the target positions on all the Schmidt plates to the system."," For this purpose, we selected compact background galaxies around each target, identified from, to transform the target positions on all the Schmidt plates to the system."
+ The small fields allowed us to apply a simple model (shift+rotation) and to achieve a higher accuracy in our final proper motion fit for all our targets (see e.g. Fig. 1))., The small fields allowed us to apply a simple model (shift+rotation) and to achieve a higher accuracy in our final proper motion fit for all our targets (see e.g. Fig. \ref{fig_PMfit}) ).
+ We consider a positive detection if the proper motion errors o of both components are below aand if at least one of the components is (above 3o)., We consider a positive detection if the proper motion errors $\sigma$ of both components are below and if at least one of the components is significant (above $\sigma$ ).
+" For J1539+0239, the brightest of our five significanttargets, we found a highly significant proper motion of u,cosó= —10.6x1.6 aand ws= —10.0+2.3yr, whereas for the other four fainter stars (with V magnitudes between 17 and 20) the proper motion was consistent with zero."," For J1539+0239, the brightest of our five targets, we found a highly significant proper motion of $\mu_\alpha\cos\,\delta=-$ $\pm$ and $\mu_\delta=-$ $\pm$, whereas for the other four fainter stars (with $V$ magnitudes between 17 and 20) the proper motion was consistent with zero."
+ The proper motions as well as their transformations from equatorial to Galactic coordinates are summarized in Table 1.., The proper motions as well as their transformations from equatorial to Galactic coordinates are summarized in Table \ref{tab_HVS}.
+" In order to exclude RV variability, we re-observed J1539+0239 with the TWIN spectrograph at the 3.5m telescope on Calar Alto in May 2009."," In order to exclude RV variability, we re-observed J1539+0239 with the TWIN spectrograph at the 3.5m telescope on Calar Alto in May 2009."
+" Radial velocities were derived by y7-fitting of adequate synthetic spectra over the full spectral range, yielding a heliocentric radial velocity of Vrad=—366.6+4.0 ffor the TWIN spectrum, which is consistent with the vq=—372.6x5.8 mmeasured from the data within the mutual uncertainties."," Radial velocities were derived by $\chi^2$ -fitting of adequate synthetic spectra over the full spectral range, yielding a heliocentric radial velocity of $v_{\rm
+rad}=-366.6\pm4.0$ for the TWIN spectrum, which is consistent with the $v_{\rm
+rad}=-372.6\pm5.8$ measured from the data within the mutual uncertainties."
+ We use the latter value for the kinematic study., We use the latter value for the kinematic study.
+ A quantitative analysis of J1539+0239 was carried out following the hybrid NLTE approach discussed by Przybillaetal. (2006).., A quantitative analysis of J1539+0239 was carried out following the hybrid NLTE approach discussed by \cite{2006A&A...445.1099P}.
+" In brief, line-blanketed LTE model atmospheres were computed with ATLAS9 (Kurucz1993) and NLTE (and LTE) line-formation calculations were performed using updated versions ofDretam and (Giddings1981;Butler&Giddings 1985)."," In brief, line-blanketed LTE model atmospheres were computed with ATLAS9 \citep{1993KurCD..13.....K} and NLTE (and LTE) line-formation calculations were performed using updated versions of and \citep{1981PhDT.......113G,bugi85}."
+". Many astrophysically important chemical species were treated in NLTE, using state-of-the-art model atoms (H: Przybilla Butler 2004a;C1:: Przybilla et al."," Many astrophysically important chemical species were treated in NLTE, using state-of-the-art model atoms (H: Przybilla Butler 2004a;: Przybilla et al."
+ 2001b;Nr: Przybilla Butler 2001;Or: Przybilla et al., 2001b;: Przybilla Butler 2001;: Przybilla et al.
+ 2000;Mgi/m:: Przybilla et al., 2000;: Przybilla et al.
+ 2001a; andFem: Becker 1998)., 2001a; and: Becker 1998).
+" The effective temperatureand the surface gravity logg were determined by fits to the Stark-broadened Balmer and Paschen lines and an ionization equilibrium, here of MgΠΠ, in analogy to previous work on hyper-velocity stars at similar temperatures (Przybillaetal.2008;Tillich 2009).. "," The effective temperatureand the surface gravity $\log g$ were determined by fits to the Stark-broadened Balmer and Paschen lines and an ionization equilibrium, here of Mg, in analogy to previous work on hyper-velocity stars at similar temperatures \citep{2008A&A...488L..51P,2009A&A...507L..37T}. ."
+The stellar was derived by model fits to the observed metal line metallicityspectra., The stellar metallicity was derived by model fits to the observed metal line spectra.
+ Results are listed in Table 1 and a comparison of the resulting final synthetic spectrum with observation is shown in Fig. 2.., Results are listed in Table \ref{tab_HVS} and a comparison of the resulting final synthetic spectrum with observation is shown in Fig. \ref{linefits}. .
+" Overall,"," Overall,"
+that we can look for in the currently available observational data.,that we can look for in the currently available observational data.
+" In the left panel of Fig. 13,,"," In the left panel of Fig. \ref{fig:sSFR-Z_BH},"
+ specific star formation rate is plotted against gas phase metallicity for all our high mass model galaxies., specific star formation rate is plotted against gas phase metallicity for all our high mass model galaxies.
+" Galaxies are binned by sSFR, and Ζει, and the bins are coloured by the mean black hole mass of the galaxies in each bin."," Galaxies are binned by $_{e}$ and $Z_{\textnormal{cold}}$, and the bins are coloured by the mean black hole mass of the galaxies in each bin."
+ The number of galaxies in each bin is also indicated on the plot in white text., The number of galaxies in each bin is also indicated on the plot in white text.
+" It is clear that low-sSFR, low-Z galaxies have the largest central black holes in the model."," It is clear that low-sSFR, $Z$ galaxies have the largest central black holes in the model."
+ The right panel of Fig., The right panel of Fig.
+ 13 shows the same plot for an adapted version of our observational Sample T2 (see Appendix B for details)., \ref{fig:sSFR-Z_BH} shows the same plot for an adapted version of our observational Sample T2 (see Appendix B for details).
+" Black hole masses were estimated via the measured stellar velocity dispersion using the Mpy-o relation provided by Tremaineetal. (2002):: where a=813, 6=4.02, σ,= 200km/s and Mgy is in units of Ma."," Black hole masses were estimated via the measured stellar velocity dispersion using the $M_{\textnormal{BH}}$ $\sigma$ relation provided by \citet{T02}: where $\alpha = 8.13$, $\beta = 4.02$, $\sigma_{0} = 200$ km/s and $M_{\textnormal{BH}}$ is in units of $\textnormal{M}_{\textnormal{\astrosun}}$."
+" We can see that the observational data also contains a Z, high-Mpg population, though the distinction between this population and the majority of the sample is less clear than in the model."," We can see that the observational data also contains a $Z$, $M_{\textnormal{BH}}$ population, though the distinction between this population and the majority of the sample is less clear than in the model."
+" The dichotomy in the observations is more clearly seen in Fig. 14,,"," The dichotomy in the observations is more clearly seen in Fig. \ref{fig:BHDists},"
+ where the distribution of black hole masses is shown for galaxies contained within the two red boxes marked in Fig. 19.., where the distribution of black hole masses is shown for galaxies contained within the two red boxes marked in Fig. \ref{fig:sSFR-Z_BH}. .
+redshift are assunied couservativelv to have coustmedo4 he available gas in their host halo.,redshift are assumed conservatively to have consumed the available gas in their host halo.
+" Figure 3) also shows he αμα allowed DII mass aud the correspoudine iininimn apparent magnitude. in variauts of our fiducial uodel: examples in which yj, is either increased to 50Ίαν+ or decreased to Llaus* (dotted and shortdashed curves). and a model in which the typical seed nass is increased to 107ML, (longdashed. curves)."," Figure \ref{fig:future} also shows the maximum allowed BH mass and the corresponding minimum apparent magnitude, in variants of our fiducial model: examples in which $\sigma_{\rm min}$ is either increased to $50~{\rm km~s^{-1}}$ or decreased to $1~{\rm
+km~s^{-1}}$ (dotted and short--dashed curves), and a model in which the typical seed mass is increased to $10^3~{\rm M_\odot}$ (long–dashed curves)."
+ The dotlongdashed curves show the results when we allow DII to consue only of the available gas in its host wo., The dot–long–dashed curves show the results when we allow BHs to consume only of the available gas in its host halo.
+ Finally. the dotshortdashed curves that bracket our fiducial model assunie survey areas of 0.1. deg? uix 1000 dee?.," Finally, the dot–short–dashed curves that bracket our fiducial model assume survey areas of 0.1 $^2$ and 1000 $^2$."
+ We fud that the survey area euters oulv logarithmically iuto our constraimts., We find that the survey area enters only logarithmically into our constraints.
+ However. the allowed BIT mass aux the resulting apparent magnitude. are strong fuuctious of redshift.," However, the allowed BH mass and the resulting apparent magnitude, are strong functions of redshift."
+ At its plauned sensitivity. the will have a detection thresholc of ~32 mag in he 1 5;nü range (with <3 hows of integration and S/N=10).," At its planned sensitivity, the will have a detection threshold of $\sim$ 32 mag in the $1-5\mu$ m range (with $\lsim 3$ hours of integration and S/N=10)."
+ The assumed survey size of 0.1 dee?can be coveredby 23 images takenbyNCST ina total observation time of2.3 davs., The assumed survey size of 0.1 $^2$ can be covered by 23 images taken by in a total observation time of 2.3 days.
+ The proposed plans to map au area of ~10 dee? to APRIME=, The proposed plans to map an area of $\sim 10$ $^2$ to $K=27$.
+" At a Πας maenitude of A~20(30). future surveys can vield strong constraiuts on the erowth of supermassive black holes. as they can map out the ""brightest quas euvelope shown iu Figure 3. out to x—10(15)."," At a limiting magnitude of $K\sim 20~(30)$, future surveys can yield strong constraints on the growth of supermassive black holes, as they can map out the “brightest quasar” envelope shown in Figure \ref{fig:future}
+ out to $z\sim 10~(15)$."
+ Existing data on highl-redslüft quasars πράος that ΟΙ» as lmassive as ~3&107AL. were asscmbled when the universe was ouly a teuth of its present age.," Existing data on high-redshift quasars implies that BHs as massive as $\sim
+3\times10^9~{\rm M_\odot}$ were assembled when the universe was only a tenth of its present age."
+ Figure 2 shows that the massive DIT inferred for SDSS 1011-0125 at 2=DS can grow in hierarchical galaxy formation models with plausible parameter values for the initial seed mass (~LO ADL). the miunuu velocity dispersion of collapsed objects that harbor such a seed C=50Iaus 1). the radiative efficienev. (~ 6%). aud the hunuinositv in Eddington units (~ 1).," Figure 2 shows that the massive BH inferred for SDSS 1044-0125 at $z=5.8$ can grow in hierarchical galaxy formation models with plausible parameter values for the initial seed mass $\sim 10~{\rm M_\odot}$ ), the minimum velocity dispersion of collapsed objects that harbor such a seed $\ga 50~{\rm km~s^{-1}}$ ), the radiative efficiency $\sim 6\%$ ), and the luminosity in Eddington units $\sim 1$ )."
+ Figure 3 illustrates the upper euvelope of plausible Iuuinositv values for high-redshift quasars that Ισ be found in future survevs., Figure 3 illustrates the upper envelope of plausible luminosity values for high-redshift quasars that might be found in future surveys.
+ There are several caveats to the constraints we derived., There are several caveats to the constraints we derived.
+ First. beaming or lensing may affect the apparcut magnitude of some of the brightest aud highest redshift quasars (Barkana Loeh 2000b).," First, beaming or lensing may affect the apparent magnitude of some of the brightest and highest redshift quasars (Barkana Loeb 2000b)."
+ Second. we have asstuned that j aud € maintain the same values during the huninous quasar phase and the main erowth phase of the DII mass.," Second, we have assumed that $\eta$ and $\epsilon$ maintain the same values during the luminous quasar phase and the main growth phase of the BH mass."
+ It is possible. however. that € is negligiblv small caving the carly growth phase of the DIT. aud high curing the luminous quasar phase (e.g. Tachuelt Rees 1993).," It is possible, however, that $\epsilon$ is negligibly small during the early growth phase of the BH, and high during the luminous quasar phase (e.g. Haehnelt Rees 1993)."
+ This would be equivalent to a corresponding increase in the seed BIT mass in our model., This would be equivalent to a corresponding increase in the seed BH mass in our model.
+ Multiple seeds per halo which eventually coalesce are also equivalent to a single massive seed with the stm of thei masses., Multiple seeds per halo which eventually coalesce are also equivalent to a single massive seed with the sum of their masses.
+ Finally. although a recent cosmological simulation fuds good agreement with he PressSchechter aass function for the lowest halo nasses of interest hore at 2=10 (Jaug-Condell Ieruquist 2000): the accuracy of this ansatz still remains to be tested over a wider ranec of redshifts and iilo masses. where there nmüght be svstematic deviations (Joulàus et al.," Finally, although a recent cosmological simulation finds good agreement with the Press–Schechter mass function for the lowest halo masses of interest here at $z=10$ (Jang-Condell Hernquist 2000); the accuracy of this ansatz still remains to be tested over a wider range of redshifts and halo masses, where there might be systematic deviations (Jenkins et al."
+ 2000: Sheth Tormen 2000)., 2000; Sheth Tormen 2000).
+ Beceutlv. various modifications to the standard CDM nodel have been proposed because of a potential conflict )etween theory aud observations ou small spatial scales (e... Namioukowski Liddle 2000: Bode. Ostriker Turok 2000: Barkana. Παππά Ostriker 2000. aux references therein).," Recently, various modifications to the standard $\Lambda$ CDM model have been proposed because of a potential conflict between theory and observations on small spatial scales (e.g., Kamionkowski Liddle 2000; Bode, Ostriker Turok 2000; Barkana, Haiman Ostriker 2000, and references therein)."
+ Our coustraints are expected to tighten sienificantly in such models. since they ecuerically reduce sinall scale power aud thus eliminate poteutial hosts of seed BUs at ligh redshifts.," Our constraints are expected to tighten significantly in such models, since they generically reduce small scale power and thus eliminate potential hosts of seed BHs at high redshifts."
+ Future surveys of high-redshift quasars will test these models., Future surveys of high-redshift quasars will test these models.
+ We thank Xiaohui Fan. Michael Strauss. and Pa Tall for useful discussious.," We thank Xiaohui Fan, Michael Strauss, and Pat Hall for useful discussions."
+" ZIT was supported bv NASA through the Hubble Fellowship erant IIE-01119.01-09À. awarded by the Space Telescope Scicuce Iustitute. which is operated by the Association of Universities for Research iu Astronomy. Iuc.. for NASA under contract NAS 5-26555,"," ZH was supported by NASA through the Hubble Fellowship grant HF-01119.01-99A, awarded by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA under contract NAS 5-26555."
+NAGS This work was supported iu part bv NASA erauts 5- 5-7768. and NSF evauts AST-0000877. AST-0071019 for AL.," This work was supported in part by NASA grants NAG 5-7039, 5-7768, and NSF grants AST-9900877, AST-0071019 for AL."
+ We find nuuerical solutious with an unuplemenutation of the generic relaxation code describe in(1998).,The diffusion equation \ref{diffusion}) ) with a source term reads: We find numerical solutions with an implementation of the generic relaxation code described in.
+. The finite-clilference scheme is fully implicit in time., The finite-difference scheme is fully implicit in time.
+ Civen a clistribution {νο}at time /. we only need to solve the one-dimensional boundary-value problem in / to fine the solution at time Fddl.," Given a distribution $f(t,J)$at time $t$, we only need to solve the one-dimensional boundary-value problem in $J$ to find the solution at time $t+dt$."
+ We thus relax successively (rom one solution to another over a variable timestep. df.," We thus relax successively from one solution to another over a variable timestep, $dt$."
+ Contijs over the size of the timestep ensure that. at each grid point. the change ii each variable over d/ is within au acceptable tolerance.," Controls over the size of the timestep ensure that, at each grid point, the change in each variable over $dt$ is within an acceptable tolerance."
+ The system is described by two first-order di(ferential equatious equivalent to Eq. (36)).," The system is described by two first-order differential equations equivalent to Eq. \ref{FPsource}) ),"
+ supplemeuted with a third dillerential equation tliat defines the nunevical domain., supplemented with a third differential equation that defines the numerical domain.
+ For simplicity. we do uot implemeut au adaptive mesh but use iusteacl a fixed ericl ou the numerical domain. chosen to be uniformly spaced in logr.," For simplicity, we do not implement an adaptive mesh but use instead a fixed grid on the numerical domain, chosen to be uniformly spaced in $\log r$."
+ We take as variables the set CF.O(Df/ü.JJ).," We take as variables the set $(f,\partial(Df)/\partial J,J)$."
+ Numerical stability of the solutions is achieved with centered-difllereuciug SChemes for f and J. aud au upwind-differencing[n] scheme for Q(Df)/04.," Numerical stability of the solutions is achieved with centered-differencing schemes for $f$ and $J$, and an upwind-differencing scheme for $\partial(Df)/\partial J$."
+ Three boundary coucitious arerequired., Three boundary conditions arerequired.
+ The ununerical domain exteuds from the inner edge of the disk. at ~LO2? Av. out to 1000AU.," The numerical domain extends from the inner edge of the disk, at $ \simeq 10^{-2}\au$ , out to $1000\au$."
+ At the inner boundary we trivially set the value of J ancl are [ree to fix f=0., At the inner boundary we trivially set the value of $J$ and are free to fix $f=0$.
+ At the outer boundary we demaucd the flux be zero., At the outer boundary we demand the flux be zero.
+ In steady-state. as discussed in re[lsubsec:steady.. this is a boundary coucdition that must be satisfied if diffusion dominates at large radii.," In steady-state, as discussed in \\ref{subsec:steady}, this is a boundary condition that must be satisfied if diffusion dominates at large radii."
+ To avoid potential complications near the disk edge. we exteud our uumerical domaiu to sulliciently large racii that it does not influeuce our results.," To avoid potential complications near the disk edge, we extend our numerical domain to sufficiently large radii that it does not influence our results."
+ We have testec our numerical results for power-law models against the analytic steady-state utions described in & subsecisteady.., We have tested our numerical results for power-law models against the analytic steady-state solutions described in \\ref{subsec:steady}.
+ Tje numerical solutions converge satisfactorily for a modest grid spacing (typically 1000 eric points ove five decades in radius)., The numerical solutions converge satisfactorily for a modest grid spacing (typically 1000 grid points over five decades in radius).
+ Additional tests at ten times this resolution show clear merical convergetrce of the steady-state results., Additional tests at ten times this resolution show clear numerical convergence of the steady-state results.
+ Agreement with the late-time behavior expected 'om the time-deperdent self-similar solutious described iu yestbsec:timedep will also be demonstrated below (see Fig. 1)., Agreement with the late-time behavior expected from the time-dependent self-similar solutions described in \\ref{subsec:timedep} will also be demonstrated below (see Fig. \ref{fig:MMSNsv}) ).
+ We study the enbiued effects of migration aud cdiffusiou ou the orbital evolution of low-uiass proto-planets for two separate classes of disk inodels: the Minimum Mass Solar Nebula (MNSN) and T Tauri alpha disks., We study the combined effects of migration and diffusion on the orbital evolution of low-mass proto-planets for two separate classes of disk models: the Minimum Mass Solar Nebula (MMSN) and T Tauri alpha disks.
+ The MMSN model adopted here has a mean surface density Xxer.7 and an aspect ratio Afrx JO which are normalized to 200gcm7 and 0.1 at εξ LAU. respectively.," The MMSN model adopted here has a mean surface density $\bar{\Sigma}\propto r^{-3/2}$ and an aspect ratio $h/r\propto
+r^{1/4}$ , which are normalized to $4200\,{\rm g\,cm^{-2}}$ and 0.1 at $r\!=\!1\au$ , respectively."
+ TheT Tauri alpha disk model adopted is identical to the irradiated disk described by, TheT Tauri alpha disk model adopted is identical to the irradiated disk described by
+ , 
+LRGs.,LRGs.
+ The full cluster will be comprised of many. galaxy most of which will not be LRGs and therefore the true shape could well be more spherical., The full cluster will be comprised of many galaxy most of which will not be LRGs and therefore the true shape could well be more spherical.
+ Other techniques that have prerequisites on the shape of clusters would not be able to detect this particular object using only the LRGs., Other techniques that have prerequisites on the shape of clusters would not be able to detect this particular object using only the LRGs.
+ The spatial cistribution of galaxies and clusters of galaxies contains a wealth of information regarding the underlying cosmological model., The spatial distribution of galaxies and clusters of galaxies contains a wealth of information regarding the underlying cosmological model.
+ The most wielely used method in the literature for condensing this information is to measure the autocorrelation function of the positional data., The most widely used method in the literature for condensing this information is to measure the autocorrelation function of the positional data.
+ For our purposes we measure the correlation function simply to compare the statistical spatial distribution of our cluster sample to other derived cluster samples., For our purposes we measure the correlation function simply to compare the statistical spatial distribution of our cluster sample to other derived cluster samples.
+ We calculate. the two-point correlation function. for the 318 2SLAQ eroups and clusters obtained: previously., We calculate the two-point correlation function for the 313 2SLAQ groups and clusters obtained previously.
+ Using the PSLAQ spectroscopic galaxy catalogue we create a ranclom catalogue. which replicates the angular completeness on the sky. this process is handled straightforwardly in the package of software.," Using the 2SLAQ spectroscopic galaxy catalogue we create a random catalogue, which replicates the angular completeness on the sky, this process is handled straightforwardly in the package of software."
+ “Phe racial distribution of the random catalogue is obtained [ron a smooth spline fit to the galaxy recdshift distribution. n(z).," The radial distribution of the random catalogue is obtained from a smooth spline fit to the galaxy redshift distribution, $n(z)$."
+ The smoothing ensures the exclusion of large scale structure voids and filaments., The smoothing ensures the exclusion of large scale structure voids and filaments.
+ The correlation function is calculated with the Landy and Szalay estimator (?)::where DD. OR and RR ave the number of pairs of points in the cata. D. and random. £2. catalogues.," The correlation function is calculated with the Landy and Szalay estimator \citep{LandySzalay:93}: :where $DD$, $DR$ and $RR$ are the number of pairs of points in the data, $D$, and random, $R$, catalogues."
+ Phe number of pairs are calculated in 14 bins. equally. separated in log space. from r=590 Mpe ht.," The number of pairs are calculated in 14 bins, equally separated in log space, from $r=5-90$ Mpc $h^{-1}$."
+ The correlation function is usually represented as a power law. £(r)=(πα). thus making a comparison with other works quite straight forward.," The correlation function is usually represented as a power law, $\xi(r)=(r/r_0)^{\gamma}$, thus making a comparison with other works quite straight forward."
+ In the left panel of fig.14 we plot the two-point correlation function of the 25LAQ clusters (blue squares with error bars) and a best-fit power-law slope (red solid line)., In the left panel of \ref{fig:cluster_2pcf} we plot the two-point correlation function of the 2SLAQ clusters (blue squares with error bars) and a best-fit power-law slope (red solid line).
+ Phe error bars are estimated using the jackknife method. which involves dividing the survey into No sections with equal area or volume.," The error bars are estimated using the jackknife method, which involves dividing the survey into $N$ sections with equal area or volume."
+ The variance and mean are estimated from No measures of our statistic., The variance and mean are estimated from $N$ measures of our statistic.
+ Each measurement is performed on the survey with region i removed. where /2 1... ," Each measurement is performed on the survey with region $i$ removed, where $i=1,...,N$ ."
+"The jackknife estimate of the variance is (2):: where Nj, is the number of Jackknife samples used and (à; represents a single bin in our statistic. £."," The jackknife estimate of the variance is \citep{Lupton:93}: where $N_{jk}$ is the number of Jackknife samples used and $r_i$ represents a single bin in our statistic, $\xi$."
+" In this analysis we set Nj,=20 sample.", In this analysis we set $N_{jk}=20$ sample.
+" The power law model is best fit. with the parameters ro=244+4 Alpe h Lands=2.1202. and the value of the reduced chi-squared is 47,=0.94."," The power law model is best fit with the parameters $r_0=24\pm4$ Mpc $h^{-1}$ and $\gamma=-2.1 \pm 0.2$, and the value of the reduced chi-squared is $\chi^2_{red} = 0.94$."
+ The right panel in f[ig.14 shows the two-point correlation function. of the mock haloes (black squares with error bars) and. a best-fit power-law slope (red. solid lino)., The right panel in \ref{fig:cluster_2pcf} shows the two-point correlation function of the mock haloes (black squares with error bars) and a best-fit power-law slope (red solid line).
+ The correlation of the mock haloes is best [it with parameters ry=23.05+40.72 Alpe fhlands =1.9540.05. which are perfectly consistent with the parameters found for the 28LAO eroups and elusters.," The correlation of the mock haloes is best fit with parameters $r_0=23.05\pm0.72$ Mpc $h^{-1}$ and $\gamma=-1.95\pm0.05$, which are perfectly consistent with the parameters found for the 2SLAQ groups and clusters."
+ Ehe error bars are measured using the jackknife method as before., The error bars are measured using the jackknife method as before.
+ Although the errors are much Larger for the real 28LAC data. this analysis indicates that the catalogue of groups ancl clusters found within 2S5LAQ using the Dok code shows the correct. level of clustering with respect. to the mock halo catalogue.," Although the errors are much larger for the real 2SLAQ data, this analysis indicates that the catalogue of groups and clusters found within 2SLAQ using the DFoF code shows the correct level of clustering with respect to the mock halo catalogue."
+ In addition. significant amounts of impurity and/or incompleteness in the cluster catalogue would suppress the correlation length as impurity adds randomly placed: structure and. incompleteness removes correlatec structure.," In addition, significant amounts of impurity and/or incompleteness in the cluster catalogue would suppress the correlation length as impurity adds randomly placed structure and incompleteness removes correlated structure."
+ Thus. as both the real ancl mock PSLAQ data have consistent correlation lengths. this implies that the amount of incompleteness ancl impurity in the PSLAQ cluster catalogue is not significant.," Thus, as both the real and mock 2SLAQ data have consistent correlation lengths, this implies that the amount of incompleteness and impurity in the 2SLAQ cluster catalogue is not significant."
+ The correlation length of the 25LAQ LRG saniple was calculated by ?.. 7.45d:0.35 Alpe fi1 and by ?.. 7.540.04 Mpe fh +. however this value is expected to be larger for clusters of galaxies.," The correlation length of the 2SLAQ LRG sample was calculated by \citet{Ross:07}, $7.45\pm0.35$ Mpc $h^{-1}$ and by \citet{Sawangwit:09}, , $7.5\pm0.04$ Mpc $h^{-1}$ , however this value is expected to be larger for clusters of galaxies."
+ Comparing the ry value obtained with those found using other low redshift cluster samples shows good aereement. ry=26.044.5 Mpe fhο). 194 Mpe fh1 (2). ISS40.9 Mpe h+ (2).," Comparing the $r_0$ value obtained with those found using other low redshift cluster samples shows good agreement, $r_0 = 26.0 \pm 4.5$ Mpc $h^{-1}$ \citep{Borgani:99}, $19.4\leq r_0 \leq23.3$ Mpc $h^{-1}$ \citep{Miller:99}, $18.8 \pm 0.9$ Mpc $h^{-1}$ \citep{Collins:00}."
+ See ? [or a good review of ry values for various cluster samples., See \citet{Nichol:01} for a good review of $r_0$ values for various cluster samples.
+ ‘The clustering length found here is also consistent with that expected for ACDAL. 22<ryx27 Mpeh + (7)," The clustering length found here is also consistent with that expected for $\Lambda$ CDM, $22\leq r_0 \leq27$ Mpc $h^{-1}$ \citep{Colberg:00}."
+ We have written an optical cluster finding algorithm based on that of ?.., We have written an optical cluster finding algorithm based on that of \citet{HuchraGeller:82}.
+ Our dvnamic friends-ol-Friends (Dor) code uses a linking length that compensates for selection ellects by changing size depending on the surface number density of galaxies at à given recdshift., Our dynamic friends-of-friends (DFoF) code uses a linking length that compensates for selection effects by changing size depending on the surface number density of galaxies at a given redshift.
+ We produced a mock catalogue. which is representative of the 2PSLAQ catalogue. using the Llorizon 47 simulation and the HOD prescription of ?. in order to determine the linking parameters. Lynas) and epaosud(z). in the DEoE code.," We produced a mock catalogue, which is representative of the 2SLAQ catalogue, using the Horizon $4\pi$ simulation and the HOD prescription of \citet{blake:08} in order to determine the linking parameters, $R_{friend}(z)$ and $v_{friend}(z)$, in the DFoF code."
+ “The code was run with various combinations of the two parameters producing a set of distinct catalogues., The code was run with various combinations of the two parameters producing a set of distinct catalogues.
+ We wrote a membership matching code in order to determine the completeness ancl purity of the resulting catalogues relative to the original mock haloes., We wrote a membership matching code in order to determine the completeness and purity of the resulting catalogues relative to the original mock haloes.
+ Dased on this analysis we chose values of Lyrical.=0.5) 0.87 Alpe fh+ ancl Cpeeed(2=0.5) 900 which correspond to a catalogue that is 98 complete anel A pure with the largest number of groups and clusters possible.," Based on this analysis we chose values of $R_{friend}(z=0.5) = $ 0.87 Mpc $h^{-1}$ and $v_{friend}(z=0.5) = $ 900 $^{-1}$, which correspond to a catalogue that is $\%$ complete and $\%$ pure with the largest number of groups and clusters possible."
+ Running the Dok code with these values. we produced a catalogue5 of 313 eroups ancl clusters containing 1152 member galaxies.," Running the DFoF code with these values, we produced a catalogue of 313 groups and clusters containing 1152 member galaxies."
+" The 5galaxy ABgroups and clusters have an average5 velocity dispersion of m,=467.97 kms and an average size of Roy=0.78 Alpe h1", The galaxy groups and clusters have an average velocity dispersion of $\overline{\sigma}_{v}=467.97$ $^{-1}$ and an average size of $\overline{R}_{clt}=0.78$ Mpc $h^{-1}$.
+ We tested the validity of our catalogue by obtaining SDSS galaxies ina 1deg? region around each cluster centre., We tested the validity of our catalogue by obtaining SDSS galaxies ina $^2$ region around each cluster centre.
+ We then subtracted the background signal from all the SDSS galaxies within 1 Mpe ο of the cluster centre and stacked the results in bins according to cluster redshift., We then subtracted the background signal from all the SDSS galaxies within 1 Mpc $h^{-1}$ of the cluster centre and stacked the results in bins according to cluster redshift.
+ Each of these bins was examined in SDSS photometricredshift spaceand colour-magnitude space., Each of these bins was examined in SDSS photometricredshift spaceand colour-magnitude space.
+ This analysis shows that our groups, This analysis shows that our groups
+"""science demonstration"" program.","""science demonstration"" program."
+" The UT1-UT3 baseline was used, resulting in a projected baseline length of 65mm at a position angle of 81° east of north."," The UT1-UT3 baseline was used, resulting in a projected baseline length of m at a position angle of $^{\circ}$ east of north."
+ The grism was used to spectrally disperse the signal at a resolution of R 2 A/A4 = 230., The grism was used to spectrally disperse the signal at a resolution of R = $\lambda$ $\Delta\lambda$ = 230.
+" This ensures that relatively narrow ""dusty"" emission features such as those of crystalline silicates and PAHs are well resolved spectrally."," This ensures that relatively narrow ""dusty"" emission features such as those of crystalline silicates and PAHs are well resolved spectrally."
+" The seeing was constant at a value of 0.6"" and the atmospheric transparency was excellent.", The seeing was constant at a value of $\arcsec$ and the atmospheric transparency was excellent.
+" We performed observations in ""High-Sens"" mode, i.e. we took an interferometric measurement combining the light from both telescopes, and subsequent photometric measurements the signal from one telescope at a time."," We performed observations in ""High-Sens"" mode, i.e. we took an interferometric measurement combining the light from both telescopes, and subsequent photometric measurements the signal from one telescope at a time."
+" Stars of known brightness and angular diameter were observed, using the same procedure, for photometric calibration and monitoring system coherence losses (interferometric ""transfer function"")."," Stars of known brightness and angular diameter were observed, using the same procedure, for photometric calibration and monitoring system coherence losses (interferometric ""transfer function"")."
+" We found our interferometric measurement of 995881 to be of significantly higher quality than the corresponding photometric measurements, and chose to directly calibrate the correlated flux rather than converting to interferometric visibility by division over the photometry."," We found our interferometric measurement of 95881 to be of significantly higher quality than the corresponding photometric measurements, and chose to directly calibrate the correlated flux rather than converting to interferometric visibility by division over the photometry."
+ A low resolution Spitzer-IRS spectrum is used to compare and flux-calibrate the VISIR spectrum., A low resolution Spitzer-IRS spectrum is used to compare and flux-calibrate the VISIR spectrum.
+ For the acquisition and reduction of the Spitzer data we refer to?., For the acquisition and reduction of the Spitzer data we refer to.
+. Photometric data points were taken from ?.., Photometric data points were taken from \cite{2004A&A...426..151A}.
+" Together with the Spitzer spectrum they make up the SED, which allows us to constrain the disk geometry (see refsec:SED))."," Together with the Spitzer spectrum they make up the SED, which allows us to constrain the disk geometry (see \\ref{sec:SED}) )."
+" Finally, II] data were taken from an earlier study, see ?.."," Finally, I] data were taken from an earlier study, see \cite{2005A&A...436..209A}."
+ A good fit ?/v ==1.8) to the AMBER data could be obtained with an inclined uniformly emitting ring surrounding a like source representing the star , A good fit $\chi^2$ $\nu$ =1.8) to the AMBER data could be obtained with an inclined uniformly emitting ring surrounding a point-like source representing the star \\ref{fig:amber}) ).
+The inner and outer radius of this ring are 2.2+0.2 mmas AAU) and 2.70.3 mmas AAU)., The inner and outer radius of this ring are $\pm$ mas AU) and $\pm$ mas AU).
+ The disk inclination is 60? 10? and its position angle is 102°+2°., The disk inclination is $^\circ\pm$ $^\circ$ and its position angle is $^\circ\pm$ $^\circ$.
+ The mean radius of the ring is mmas (~0.4 AAU)., The mean radius of the ring is mas $\sim$ AU).
+ This is similar to the values found by ? for a group of Herbig Ae/Be stars., This is similar to the values found by \cite{2004ApJ...613.1049E} for a group of Herbig Ae/Be stars.
+ The fraction of the K-band flux that is ascribed to the star by the fit of a Kurucz model to the optical photometry (~20%)) is not exactly equal to the AMBER estimate of 32+4%., The fraction of the K-band flux that is ascribed to the star by the fit of a Kurucz model to the optical photometry $\sim$ ) is not exactly equal to the AMBER estimate of $\pm$.
+". It may be that this discrepancy is due to the presence of an unresolved disk component, e.g. hot gas the dust sublimation radius, that contributes to the K-band emission. (23; 25; 25; 25;"," It may be that this discrepancy is due to the presence of an unresolved disk component, e.g. hot gas the dust sublimation radius, that contributes to the K-band emission. \citealt{2007ApJ...657..347E}; \citealt{2008ApJ...676..490K}; \citealt{2008A&A...485..209A}; \citealt{2008ApJ...677L..51T};"
+ and ?))., and \citealt{2008A&A...483L..13I}) ).
+ ? have investigated the II] ΤΑ emission line in a large sample of Herbig stars., \cite{2005A&A...436..209A} have investigated the I] \\r{A} emission line in a large sample of Herbig stars.
+ The authors argue that the emission is non-thermal and originates from the disk surface of a flared circumstellar disk., The authors argue that the emission is non-thermal and originates from the disk surface of a flared circumstellar disk.
+ This is consistent with the large detection rate of II] emitters among the group I sources in the sample., This is consistent with the large detection rate of I] emitters among the group I sources in the sample.
+" Roughly half of the group II sources,"," Roughly half of the group II sources,"
+We analysed the time series of MACRO iuunous using two complementary approaches: search for periodicities aud search for bursts of eveuts.,We analysed the time series of MACRO muons using two complementary approaches: search for periodicities and search for bursts of events.
+ The Lomb-Scarele method was used in the first case. scan statistics in the secoud.," The Lomb-Scargle method was used in the first case, scan statistics in the second."
+" The two techuiques complete carly analyses performed with ""folding methods iu searching for periodicities aud time differcuces for burst eveuts.", The two techniques complete early analyses performed with “folding” methods in searching for periodicities and time differences for burst events.
+ No deviations from the expected distributions were found., No deviations from the expected distributions were found.
+ We thank the members of the MACRO Collaboration aud the personnel of the LNCS for their cooperation., We thank the members of the MACRO Collaboration and the personnel of the LNGS for their cooperation.
+ ILD. and S.N. thank the Abdus Salam ICTP TRIL Proeramuune for providing fellowships., H.D. and S.M. thank the Abdus Salam ICTP TRIL Programme for providing fellowships.
+ IL.D.. J.D.. €.CG.. F.M. aud A. Moussa thank the collaboration between the Universities of Bologna and Mohamed D of Oujda.," H.D., J.D., G.G., F.M. and A. Moussa thank the collaboration between the Universities of Bologna and Mohamed $1^{st}$ of Oujda."
+We thank O.B. Slee. A.L. Rov. M. Murgia. HH. Xndernach. AL. Ehle for providing us with their observational cala prior to publication. ancl allowing us to displav their 1.4 Gllz map of the relic in Abell 85.,"We thank O.B. Slee, A.L. Roy, M. Murgia, H. Andernach, M. Ehle for providing us with their observational data prior to publication, and allowing us to display their 1.4 GHz map of the relic in Abell 85."
+ We also thank €. €jiovannini and L. Feretti for access to their 330. MlIz daa ol the same relic., We also thank G. Giovannini and L. Feretti for access to their 330 MHz data of the same relic.
+ We acknowledge A. Ixerceks contributjons to à very early stage of this simulation project and L. Mülller's comments on the manuscript., We acknowledge A. Kercek's contributions to a very early stage of this simulation project and E. Mülller's comments on the manuscript.
+ Some of the com.»utations reported. here were performed. using the Ulx. Astrophysical Fluids Facility (UISΑΟ)., Some of the computations reported here were performed using the UK Astrophysical Fluids Facility (UKAFF).
+ This work was supported by the European Community Research and Training Network “Lhe Physics of the Intergalactic Medium., This work was supported by the European Community Research and Training Network `The Physics of the Intergalactic Medium'.
+"volume where the smoothed density exceeds ὃς,",volume where the smoothed density exceeds $\delta_c$.
+ Assuming a eaussian density field. then where & is the root-mean-square fluctuations. within the top-hat filter ancl erfe is the complementary error function.," Assuming a gaussian density field, then where $\sigma$ is the root-mean-square fluctuations within the top-hat filter and ${\rm erfc}$ is the complementary error function."
+ The Κον step was to realize that Uuetuations on dillerent massescales are not. independent., The key step was to realize that fluctuations on different mass-scales are not independent.
+ In fact. to a first approximation PS assumec that high-mass halos were entirely made up of lower-mass ones with πο underdense matter mixed in.," In fact, to a first approximation PS assumed that high-mass halos were entirely made up of lower-mass ones with no underdense matter mixed in."
+ Then £ must be regarded as a cumulative mass fraction and it can be clilferentiatecl to obtain the dillerential one. The main drawback of this approach is that. because of the above assumption of crowding together of low-mass ialos into larger ones. it scems to undercount the number of objects.," Then $F$ must be regarded as a cumulative mass fraction and it can be differentiated to obtain the differential one, The main drawback of this approach is that, because of the above assumption of crowding together of low-mass halos into larger ones, it seems to undercount the number of objects."
+ As Adu>0 (and therefore oi17» o) the fraction of the Universe which exceeds the density threshold. tends o one half., As $M\mapsto0$ (and therefore $\sigma\mapsto\infty$ ) the fraction of the Universe which exceeds the density threshold tends to one half.
+ For this reason it is usual to multiply f by wo to rellect the fact that most of the Universe today is contained in collapsed structures., For this reason it is usual to multiply $f$ by two to reflect the fact that most of the Universe today is contained in collapsed structures.
+ We call this the correctec S prediction., We call this the corrected PS prediction.
+" Extensions of the PS prescription. to caleulate explicitLy he integrated. merger history of all halos. were [irs developed by Bower (1991) and then rederived. using a more mathematically motivated theory (called the Excursion Se πουν, hereafter IL. Bond 1992). and tested against N-body experiments. by Lacey Cole (1993. 1994)."," Extensions of the PS prescription, to calculate explicitly the integrated merger history of all halos, were first developed by Bower (1991) and then rederived, using a more mathematically motivated theory (called the Excursion Set Theory, hereafter EST, Bond 1992), and tested against N-body experiments, by Lacey Cole (1993, 1994)."
+" In this formalism the top-ha smoothing radius about a given. point is first set to a very arge value and then gradually. reduced. until the enclosec overdensity exceeds 2, (in hierarchical cosmologies this wil always occur before the radius shrinks to zero).", In this formalism the top-hat smoothing radius about a given point is first set to a very large value and then gradually reduced until the enclosed overdensity exceeds $\delta_c$ (in hierarchical cosmologies this will always occur before the radius shrinks to zero).
+ This gives he largest. region. which will have collapsed around tha xont., This gives the largest region which will have collapsed around that point.
+ There may be smaller regions which have a larger overdensity. but. these merely represent. smaller structures which have been subsumed into the larger one., There may be smaller regions which have a larger overdensity but these merely represent smaller structures which have been subsumed into the larger one.
+ By varving he density. threshold. one can build up a picture of the collapse and. merger-history of the halos: in essence this xper describes a numerical representation of this process., By varying the density threshold one can build up a picture of the collapse and merger-history of the halos: in essence this paper describes a numerical representation of this process.
+ Surprisingly perhaps. the IST predicts the same distribution. of halo masses as oes the PS theory (but without the need. for the extra factor of two in normalization).," Surprisingly perhaps, the EST predicts the same distribution of halo masses as does the PS theory (but without the need for the extra factor of two in normalization)."
+ Despite being very idealizecl in nature. ignoring both the internal structure and. tidal forces. the derived. formulae provide a surprisingly good fit to the N-body results (Efstathiou 1988. Lacey Cole 1994. Gelb Bertschinger 1994).," Despite being very idealized in nature, ignoring both the internal structure and tidal forces, the derived formulae provide a surprisingly good fit to the N-body results (Efstathiou 1988, Lacey Cole 1994, Gelb Bertschinger 1994)."
+ Llowever we have to regard these sucesses with some scepticism. since the basic hypothesis of the E works very poorly on a object-by-object basis (White 1995). the numerical simulations are still plagued by resolution elfects and limited dywnamical range. ancl the halo statistics are sensitive to the scheme chosen for identifving halos.," However we have to regard these sucesses with some scepticism, since the basic hypothesis of the EST works very poorly on a object-by-object basis (White 1995), the numerical simulations are still plagued by resolution effects and limited dynamical range, and the halo statistics are sensitive to the scheme chosen for identifying halos."
+ Moreover one should always bear in mind. that the PS treatment is a linear approach to a problem. which is fundamentally non-linear in nature., Moreover one should always bear in mind that the PS treatment is a linear approach to a problem which is fundamentally non-linear in nature.
+ The full non-linear. evolution of structure ijs. best described. by an N-body simulation., The full non-linear evolution of structure is best described by an N-body simulation.
+ Moreover. with the introduction of techniques such as smoothed particle ivdrodyvnamiecs (SPI). it is possible to. simultaneously ollow the evolution of a dissipative. continuous intergalactic medium.," Moreover, with the introduction of techniques such as smoothed particle hydrodynamics (SPH), it is possible to simultaneously follow the evolution of a dissipative, continuous intergalactic medium."
+ Llowever. there are several drawbacks to. this approach: N-bods simulations are very time-consuming. hey have a limited dvnamiucal range and they are very inflexible when trving to model the physical processes iippening on small scales (with small numbers of particles).," However, there are several drawbacks to this approach: N-body simulations are very time-consuming, they have a limited dynamical range and they are very inflexible when trying to model the physical processes happening on small scales (with small numbers of particles)."
+ For example. it is likely that the interstellar medium in a owotogalaxyv will contain a mixture of hold. and. cold. gas as well as stars with a variety of ages and dark matter.," For example, it is likely that the interstellar medium in a protogalaxy will contain a mixture of hold and cold gas as well as stars with a variety of ages and dark matter."
+ Simulations which can handle such situations are only just »'einning to appear., Simulations which can handle such situations are only just beginning to appear.
+ Thus it is highly desirable to set up a simple but ellicient Monte-Carlo. procedure which mimics the general features of the hierarchical clustering process ancl can be used to carey out a laree parameter investigation with Little timie-consumption., Thus it is highly desirable to set up a simple but efficient Monte-Carlo procedure which mimics the general features of the hierarchical clustering process and can be used to carry out a large parameter investigation with little time-consumption.
+ The first mocdel to be presented in those lines was the oof Cole Ixaiser (1989). used first to study the abundance of clusters and. subsequently some aspects. of. galaxy formation (Cole 1991. Cole 1994).," The first model to be presented in those lines was the of Cole Kaiser (1989), used first to study the abundance of clusters and subsequently some aspects of galaxy formation (Cole 1991, Cole 1994)."
+ Ho starts with a large cuboidal block. with sides in the ratio 1:2Les and subdivides it into two sub-blocks of the same shape.," It starts with a large cuboidal block, with sides in the ratio $1:2^{1/3}:2^{2/3}$, and subdivides it into two sub-blocks of the same shape."
+" Lf the initial block has an overdensity 9. (drawn from a gaussian with variance o(A)). then the two sub-blocks will inherit the same overdensity with an extra perturbation. added to one of them ancl subtracted from the other. drawn from a eaussian. with. variance. X. where M2=o7(M)c""(Mjf2)."," If the initial block has an overdensity $\delta$ (drawn from a gaussian with variance $\sigma(M)$ ), then the two sub-blocks will inherit the same overdensity with an extra perturbation, added to one of them and subtracted from the other, drawn from a gaussian with variance $\Sigma$, where $\Sigma^2=\sigma^2(M)-\sigma^2(M/2)$."
+ This quadratic procedure is applied. iteratively to each of 10 sub-blocks until the imposed mass resolution is achieved., This quadratic procedure is applied iteratively to each of the sub-blocks until the imposed mass resolution is achieved.
+ The advantage of the method comes from the fact that the relative position of all sub-blocks is known at all times so that it is simple to follow the merger history of any. halo etected at any stage of the simulation., The advantage of the method comes from the fact that the relative position of all sub-blocks is known at all times so that it is simple to follow the merger history of any halo detected at any stage of the simulation.
+ Ixaullmann White (1993) adopt a dillerent approach which makes use of the conditional merging probabilities erived by Bower (1991)., Kauffmann White (1993) adopt a different approach which makes use of the conditional merging probabilities derived by Bower (1991).
+ Given that a halo has a particular mass at some redshift. then one can work out the probability istribution for the mass of the halo (centred on the same point) at some earlier redshift.," Given that a halo has a particular mass at some redshift, then one can work out the probability distribution for the mass of the halo (centred on the same point) at some earlier redshift."
+ By generating a large number of representative halos. say 100 or more. it. is possible to allocate sub-halos with the correct. spectrum: of masses.," By generating a large number of representative halos, say 100 or more, it is possible to allocate sub-halos with the correct spectrum of masses."
+ This method. gives a wider mass-spectrum for halos (not restricted to powers of two) but restricts halo formation to occur at specific redshifts and is much more complicated to implement than theMocel., This method gives a wider mass-spectrum for halos (not restricted to powers of two) but restricts halo formation to occur at specific redshifts and is much more complicated to implement than the.
+. Llere we present a new method for following halo evolution which is much closer in spirit to the N-box= simulations without compromising the simplicity ancl speed of the above analytical techniques., Here we present a new method for following halo evolution which is much closer in spirit to the N-body simulations without compromising the simplicity and speed of the above analytical techniques.
+ It allows a continuous spectrum of halo masses (above a minimum. of S unit cells) anc a variable collapse time., It allows a continuous spectrum of halo masses (above a minimum of 8 unit cells) and a variable collapse time.
+ We start with a full, We start with a full
+The typical classical nova (CN) system consists of a white dwarf (WD) primary and a main-sequence secondary in a close binary orbit: the secondary star fills its Roche lobe.,The typical classical nova (CN) system consists of a white dwarf (WD) primary and a main-sequence secondary in a close binary orbit; the secondary star fills its Roche lobe.
+ Material lost by the secondary accumulates onto the surface of the WD via an accretion disk around the WD which lies in the orbital plane., Material lost by the secondary accumulates onto the surface of the WD via an accretion disk around the WD which lies in the orbital plane.
+ In a small number of cases - the - accretion is channelled to the magnetic poles of the WD and no appreciable disk is present., In a small number of cases - the - accretion is channelled to the magnetic poles of the WD and no appreciable disk is present.
+ Nuclear burning can initiate within this accreted. degenerate. material. leading to a thermonuclear runaway which will eject some or all of this accreted laver.," Nuclear burning can initiate within this accreted, degenerate, material, leading to a thermonuclear runaway which will eject some or all of this accreted layer."
+ A nova eruption can reach My=—10 (Shafteretal.2009) and may eject 10—107 Me of matter (e.g.Bode2010.andref-erencestherein ).., A nova eruption can reach $M_{V}\simeq-10$ \citep{2009ApJ...690.1148S} and may eject $10^{-5}-10^{-4}$ $_{\sun}$ of matter \citep[e.g.][and references therein]{2010AN....331..160B}.
+ The closely related recurrent novà (RN) systems exhibit recurrence timescales of 10—100 years., The closely related recurrent nova (RN) systems exhibit recurrence timescales of $10-100$ years.
+ This ts attributed to a combination of a high mass WD and high aceretion rate., This is attributed to a combination of a high mass WD and high accretion rate.
+ The elevated accretion rate 1s caused by the presence of a evolved secondary: either a sub-giant (U Sco-class) or a red giant (RS Oph-class)., The elevated accretion rate is caused by the presence of an evolved secondary; either a sub-giant (U Sco-class) or a red giant (RS Oph-class).
+ However. the T Pyx-class of recurrents contain lower mass WDs and main sequence secondaries. but other than their short recurrence timescales. they are more aki to CNe (Anupama2008)..," However, the T Pyx-class of recurrents contain lower mass WDs and main sequence secondaries, but other than their short recurrence timescales, they are more akin to CNe \citep{2008ASPC..401...31A}."
+ V249] Cyent (Nova Cygni 2008 #22) was discovered 1 outburst on 2008 April 10.8 UT with an unfiltered magnitude m=7.7 (Nakanoetal.2008).., V2491 Cygni (Nova Cygni 2008 2) was discovered in outburst on 2008 April 10.8 UT with an unfiltered magnitude $m=7.7$ \citep{2008IAUC.8934....1N}.
+ Its nature was confirmec spectroscopically by Ayant&Matsumoto(2008.HaFHWM~4.500kmο} and was latter classified as a He/ nova (Heltonetal.2008;Munart2011)..," Its nature was confirmed spectroscopically by \citet[H$\alpha$ FHWM $\sim4,500$~km~s$^{-1}$ and was latter classified as a He/N nova \citep{2008CBET.1379....1H,2010arXiv1009.0822M}."
+ The rapid declie from maximum light classifies V2491 Cyg as a nova., The rapid decline from maximum light classifies V2491 Cyg as a nova.
+ The optical decline of V2491 Cyg exhibited a secondary maximum at day ~15 (Munarietal.2011).., The optical decline of V2491 Cyg exhibited a secondary maximum at day $\sim15$ \citep{2010arXiv1009.0822M}.
+ Such significant have been seen in a number of other novae (V2362CygandVI493Αα:Stropeetal.2010).. but are unusual. and still poorly understood.," Such significant have been seen in a number of other novae \citep[V2362~Cyg and V1493~Aql;][]{2010AJ....140...34S}, but are unusual, and still poorly understood."
+ Hachisu&Kato(2009) have proposed magnetic activity as an additional energy source and a possible cause of the re-brightening. a requirement being a highly magnetised WD (a polar).," \citet{2009ApJ...694L.103H} have proposed magnetic activity as an additional energy source and a possible cause of the re-brightening, a requirement being a highly magnetised WD (a polar)."
+ Following the outburst. it was discovered that there was à pre-existing X-ray source coincident with the position of V249] Cyg (Ibarra&Kuulkers2008;Ibarraetal.2008. 2009).. ," Following the outburst, it was discovered that there was a pre-existing X-ray source coincident with the position of V2491 Cyg \citep{2008ATel.1473....1I,2008ATel.1478....1I,2009A&A...497L...5I}. ."
+This is only the second nova for which pre-outburst X-ray emission has been detected. after V2487 Oph (Hernanz&Sala2002)..," This is only the second nova for which pre-outburst X-ray emission has been detected, after V2487 Oph \citep{2002Sci...298..393H}."
+ V2487 Oph was classified as a recurrent nova when a previous outburst was indentified from 1900 (Pagnottaet2009 ).. Bak, V2487 Oph was classified as a recurrent nova when a previous outburst was indentified from 1900 \citep{2009AJ....138.1230P}.
+lanovetal.(2008) reported a variation in the B- and V- bands between 10 and 20 days after outburst. these had a period of 0.09580(5) days and amplitude of 0.03—0.05 mag.," \citet{2008ATel.1514....1B} reported a variation in the $B$ - and $V$ -bands between 10 and 20 days after outburst, these had a period of 0.09580(5) days and amplitude of $0.03-0.05$ mag."
+ This period has been taken by some authors to be the orbital period of the system., This period has been taken by some authors to be the orbital period of the system.
+ However. Baklanov (private communication) have not made this claim and also note that this period became negligible as the system waned.," However, Baklanov (private communication) have not made this claim and also note that this period became negligible as the system waned."
+ Raganetal.(2010) reported similar R-band periodicity between day 28 and 34., \citet{2010arXiv1004.0419R} reported similar $R$ -band periodicity between day 28 and 34.
+ The first indication that V2491 Cyg may be recurrent in nature came from Tomovetal.(2008) who noted the spectral similarities to the RNe U Sco and V394 CrA early after outburst., The first indication that V2491 Cyg may be recurrent in nature came from \citet{2008ATel.1485....1T} who noted the spectral similarities to the RNe U Sco and V394 CrA early after outburst.
+ Additionally. Bodeetal.(2009) noted that the V249] Cye spectra were similar to the early spectra of RS Oph and in particular those of M31N 2007-12b. a RS Oph-class RN candidate in M31.," Additionally, \citet{2009ApJ...705.1056B} noted that the V2491 Cyg spectra were similar to the early spectra of RS Oph and in particular those of M31N 2007-12b, a RS Oph-class RN candidate in M31."
+ The progenitor system of V249] Cyg was identified as USNO-B1.0 1223-042965 (Henden&Munari2008) and there is no evidence of previous outbursts at this position (Jurdana-Sepic&Μιπα, The progenitor system of V2491 Cyg was identified as USNO-B1.0 1223-042965 \citep{2008IBVS.5834....1H} and there is no evidence of previous outbursts at this position \citep{2008IBVS.5839....1J}.
+ 2008).. Rudyetal.(2008) determined a reddening to V249] Cyg of Eg.=0.43 via ΟΙ line ratios., \citet{2008IAUC.8938....2R} determined a reddening to V2491 Cyg of $E_{B-V}=0.43$ via OI line ratios.
+ Heltonetal.(2008) used the CN MMRD relation (dellaValle&Livio1995) with this reddening to determine d=10.5 kpe., \citet{2008CBET.1379....1H} used the CN MMRD relation \citep{1995ApJ...452..704D} with this reddening to determine $d=10.5$ kpc.
+ Munari=etal.(2011) independently derived both the extinction and distance to V2491 Cyg using the interstellar Nal line and the vandenBergh&Younger(1987) relation. finding Ep-y=0.23+0.01 and d=14 kpe.," \citet{2010arXiv1009.0822M} independently derived both the extinction and distance to V2491 Cyg using the interstellar NaI line and the \citet{1987A&AS...70..125V} relation, finding $E_{B-V}=0.23\pm0.01$ and $d=14$ kpc."
+" Ribeiroetal.(2011) reported an ejecta morphology consisting of polar blobs and an equatorial ring. with expansion velocities of ~3.000 km s! and an melination of 8077, degrees. i.e. close to edge on."," \citet{Ribeiro2010} reported an ejecta morphology consisting of polar blobs and an equatorial ring, with expansion velocities of $\sim3,000$ km $^{-1}$ and an inclination of $80^{+3}_{-12}$ degrees, i.e. close to edge on."
+ Pageetal.(2010) reported the results of extensive X- and UV observations of the outburst. these data indicated that the WD in the system may be close to the Chandrasekhar mass.," \citet{2010MNRAS.401..121P} reported the results of extensive X-ray and UV observations of the outburst, these data indicated that the WD in the system may be close to the Chandrasekhar mass."
+ By assuming the pre-outburst X-ray emission was due to aceretion Pageetal.(2010) deduced that V249] Cyg would recur on longer than typical timescales. centuries rather than decades. if the WD masswas indeed high.," By assuming the pre-outburst X-ray emission was due to accretion \citet{2010MNRAS.401..121P} deduced that V2491 Cyg would recur on longer than typical timescales, centuries rather than decades, if the WD masswas indeed high."
+ This paper noted in, This paper noted in
+Iluminosity) since zcO.8 in an Einstein-de Sitter Universe.,luminosity) since $z\simeq0.8$ in an Einstein-de Sitter Universe.
+ Despite this inherent. degeneracy the Kormencdy relations can firmly exelude very strong growth. (> 50%) from mergers in the redshift range 0.2<2<0.8., Despite this inherent degeneracy the Kormendy relations can firmly exclude very strong growth $\ge 50 \%\ $ ) from mergers in the redshift range $0.2<z<0.8$.
+ The results from a thorough re-examination of the BLK LUST images of a sample of 2~1 38CR. radio galaxies have been presented., The results from a thorough re-examination of the BLR HST images of a sample of $z\sim1$ 3CR radio galaxies have been presented.
+ It has been shown that. in terms of scalelength. absolute magnitudes and IxXormendy. relation there are no significant dillerences between 2~O.S and," It has been shown that, in terms of scalelength, absolute magnitudes and Kormendy relation there are no significant differences between $z\sim0.8$ and"
+"2010),, a better display of the data can be done by averaging the two Fe features, like in a composite index which actually sums up the equivalent width of both features.",", a better display of the data can be done by averaging the two Fe features, like in a composite index which actually sums up the equivalent width of both features."
+" An even cleaner result could also be secured for our dataset by taking advantage of the full observation of the 5 “Fe” indices comprised in the Lick system within the rrange (see, again, Table 8)), excluding the Fe5015 feature possibly affected in case of [ΟΠήροοτ emission."," An even cleaner result could also be secured for our dataset by taking advantage of the full observation of the 5 “Fe” indices comprised in the Lick system within the range (see, again, Table \ref{t8}) ), excluding the Fe5015 feature possibly affected in case of $_{5007}$ emission."
+ A further composite index can therefore be built up as Figure 18 gives a summary of our results., A further composite index can therefore be built up as Figure \ref{f18} gives a summary of our results.
+" Again, for reader’s better reference, our observed galaxy sample is matched in all the plots of the figure with the Rampazzoetal.(2005) and Annibalietal.(2006) standard ellipticals, as well as with the compilation of M31 and Galactic globulars according to Trageretal.(1998)."," Again, for reader's better reference, our observed galaxy sample is matched in all the plots of the figure with the \citet{rampazzo05} and \citet{annibali06} standard ellipticals, as well as with the compilation of M31 and Galactic globulars according to \citet{trager98}."
+". The (fe)2 index can also be easily computed for the theoretical models of Buzzoni,Mantegazza,&Gariboldi(1994) and it is displayed in the upper plot of the figure as a reference locus for 10 Gyr old SSPs to compare with the observed distribution.", The $\langle Fe\rangle _2$ index can also be easily computed for the theoretical models of \citet{buzzoni94} and it is displayed in the upper plot of the figure as a reference locus for 10 Gyr old SSPs to compare with the observed distribution.
+"estimated to vary from tens to hundreds of seconds depending on the model used for the fit (i.e. a Gaussian or a power law respectively), we thus assume for AGILE and GLAST integration times of 100 s and 500 s. We find that for a flare luminosity L=2x10*°ergs! the (predicted) high energy flare emission can be detected by AGILE and GLAST up to a maximum redshift Za,~1.2 and ζω~1.4 for an integration time of 500 s, and up to ζω~0.6 and z,4,~0.7 for an integration time of 100 s. We then calculate the emission expected for X-ray flare luminosity L’=10°°ergs! as a function of the fireball Lorentz factor T, at z=1 (Fig. 8)).","estimated to vary from tens to hundreds of seconds depending on the model used for the fit (i.e. a Gaussian or a power law respectively), we thus assume for AGILE and GLAST integration times of 100 s and 500 s. We find that for a flare luminosity $L=2 \times 10^{49}~erg~s^{-1}$ the (predicted) high energy flare emission can be detected by AGILE and GLAST up to a maximum redshift $z_{max} \sim 1.2$ and $z_{max} \sim 1.4$ for an integration time of 500 s, and up to $z_{max} \sim 0.6$ and $z_{max} \sim 0.7$ for an integration time of 100 s. We then calculate the emission expected for X-ray flare luminosity $L'=10^{50}~erg~s^{-1}$ as a function of the fireball Lorentz factor $\Gamma$, at $z=1$ (Fig. \ref{flarelum2e50z1}) )."
+ With this X-ray flares luminosity the predicted high energy emission can be detected both by AGILE and GLAST at z=1., With this X-ray flares luminosity the predicted high energy emission can be detected both by AGILE and GLAST at $z=1$.
+" If we take as mean flare luminosity (L/)=10°°ergs-!. the SSC flare emission can be detected by AGILE and GLAST up to maximum redshift Zmax2.4 and Zmax~2.8 for an integration time of 500 s, and Up tO ζην~1.2 and Zax—1.4 for an integration time of 100 s."," If we take as mean flare luminosity $\langle L' \rangle = 10^{50}~erg~s^{-1}$, the SSC flare emission can be detected by AGILE and GLAST up to maximum redshift $z_{max} \sim 2.4$ and $z_{max} \sim 2.8$ for an integration time of 500 s, and up to $z_{max} \sim 1.2$ and $z_{max} \sim 1.4$ for an integration time of 100 s."
+Rieger Mannheim 2003).,Rieger Mannheim 2003).
+ It is this conjecture which we wish to explore here in more detail., It is this conjecture which we wish to explore here in more detail.
+" According to the VHE variability analysis for July 2006. there are two major findings which need explanation: (1) First. the extreme short-term variability on a timescale as low as f,~200 sec (in the observer'sframe)!.. and (11) secondly. the log-normal distribution of fluxes. which points to an underlying. multiplicative process."," According to the VHE variability analysis for July 2006, there are two major findings which need explanation: (i) First, the extreme short-term variability on a timescale as low as $t_v \sim 200$ sec (in the observer's, and (ii) secondly, the log-normal distribution of fluxes, which points to an underlying, multiplicative process."
+ Consider the first finding If the luminosity of the host galaxy of PKS 2155-304 is indeed as high as reported. 1e.. M(R)=—24.4 mag (Kotilainen et al.," Consider the first finding If the luminosity of the host galaxy of PKS 2155-304 is indeed as high as reported, i.e., $M(R) \simeq 
+-24.4$ mag (Kotilainen et al."
+" 1998). its central BH mass 15 (even if one takes the scatter in the μη - Logs; relation into account) expected to exceed (MeLure Dunlop 2002) M,presence 2x I0M.."," 1998), its central BH mass is (even if one takes the scatter in the $M_{\rm 
+BH}$ - $L_{\rm bulge}$ relation into account) expected to exceed (McLure Dunlop 2002) M_t 2 10^8."
+. Accounting for the observed rapid variations based on conventional assumptions has thus proved challenging (see. e.g.. Levinson 2007 and Begelman et al.," Accounting for the observed rapid variations based on conventional assumptions has thus proved challenging (see, e.g., Levinson 2007 and Begelman et al."
+" 2008 for detailed discussions): In the galaxy’s rest frame. the characteristic scale of any disturbance produced by the central engine cannot be smaller than its minimum size r,. so that the minimum variability timescale for temporal fluctuations of the ejected fluid is r/c. where ryGM/c7 is the gravitational radius (Levinson 2008)."," 2008 for detailed discussions): In the galaxy's rest frame, the characteristic scale of any disturbance produced by the central engine cannot be smaller than its minimum size $r_g$, so that the minimum variability timescale for temporal fluctuations of the ejected fluid is $r_g/c$, where $r_g=GM/c^2$ is the gravitational radius (Levinson 2008)."
+ Only if parts of the jet flow would collide with a small external disturbance. e.g. a cloud. may smaller variability times get imprinted.," Only if parts of the jet flow would collide with a small external disturbance, e.g. a cloud, may smaller variability times get imprinted."
+" Yet. in this case the bulk power that can be tapped for producing radiation will be correspondingly smaller. and the noted PSD shape and lognormal flux distribution may still be in need of explanation,"," Yet, in this case the bulk power that can be tapped for producing radiation will be correspondingly smaller, and the noted PSD shape and lognormal flux distribution may still be in need of explanation."
+" If the observed minimum indeed reflect the size of the central engine. ptisf,dme.~vousderr.jc. the allowed maximum BH mass would be one of magnitude smaller than the one inferred from the host galaxy luminosity relation. i.e.. 4% 10 Me."," If the observed minimum variability would indeed reflect the size of the central engine, $t_v \sim (1+z) r_g/c$, the allowed maximum BH mass would be almost one order of magnitude smaller than the one inferred from the host galaxy luminosity relation, i.e., M_v 4 10^7 )."
+ Similar results (Gi)are obtained if the variability is taken to be produced by the collision of shells moving at differentspeeds (as in classical internal shock scenarios)., Similar results are obtained if the variability is taken to be produced by the collision of shells moving at different speeds (as in classical internal shock scenarios).
+" In all these cases. the inferred minimum source size (in the rest frame of the becomes independent of any Doppler or galaxy)Interestingly. similar small BH masses M,Lorenzx10M. for PKS 2155-304 have been previously inferred based on its X-ray variability properties (Hayashida et al."," In all these cases, the inferred minimum source size (in the rest frame of the galaxy) becomes independent of any Doppler or Lorentz Interestingly, similar small BH masses $M_v \lppr 3 \times 10^7 M_{\odot}$ for PKS 2155-304 have been previously inferred based on its X-ray variability properties (Hayashida et al."
+ 1998: Czerny et al., 1998; Czerny et al.
+ 2001)., 2001).
+ Moreover. Lachowiez et al. (," Moreover, Lachowicz et al. ("
+2009) have recently suggested that the detection of a possible ~4.6 hr X-ray quasi-periodic oscillation (ΟΡΟ) with XMM-Newton in 2006 might be related to a hot spot orbiting in the disk close to the innermost stable orbit.,2009) have recently suggested that the detection of a possible $\sim4.6$ hr X-ray quasi-periodic oscillation (QPO) with XMM-Newton in 2006 might be related to a hot spot orbiting in the disk close to the innermost stable orbit.
+ If this interpretation is correct. is would imply a firm upper limit on the BH mass of €3xI0/M. (Schwarzschild) and <2.1x105M. (extreme Kerr). One obvious possibility to reconcile these apparently divergent central mass estimates is related to the putative presence of a close binary BH system where the jet that dominates the high-energy emission originates from the less massive (secondary)BH.," If this interpretation is correct, is would imply a firm upper limit on the BH mass of $\leq 3 \times 10^7 M_{\odot}$ (Schwarzschild) and $\leq 2.1\times 10^8 M_{\odot}$ (extreme Kerr), One obvious possibility to reconcile these apparently divergent central mass estimates is related to the putative presence of a close binary BH system where the jet that dominates the high-energy emission originates from the less massive (secondary) BH."
+ In such a case. host galaxy observations may only be indicative of the total (primary and secondaryreveal) BH mass. whereas high-energy observations could actually signatures of the less massive BH.," In such a case, host galaxy observations may only be indicative of the total (primary and secondary) BH mass, whereas high-energy observations could actually reveal signatures of the less massive BH."
+ We note that the possible of close binary BH systems in blazar-type AGN sources has been repeatedly invoked to account for a variety of observational findings. most notably for the detection of mid-and long-term quasi-periodic variability (e.g.. see Rieger 2007 for a If the observed VHE output in PKS 2155-305 were be dominated by emission from the secondary BH. however. its mass cannot be supposed to be too small if one also wishes to account for the required jet power.," We note that the possible presence of close binary BH systems in blazar-type AGN sources has been repeatedly invoked to account for a variety of observational findings, most notably for the detection of mid-and long-term quasi-periodic variability (e.g., see Rieger 2007 for a If the observed VHE output in PKS 2155-305 were be dominated by emission from the secondary BH, however, its mass cannot be supposed to be too small if one also wishes to account for the required jet power."
+" Suppose that the maximum possible jet power Is constrained by the Eddington accretion limit Lye~—107(Mggu/109M). erg/s. The maximum spin power. for example. which can be extracted from the rotational energy of a BH (via a Blandford-Znajek-type process) is on the orde rof Ly~107(Mgg/10MÁY(By/10'GY ere/s. With an Eddington equipartition magnetic field strength By,x1/ VMpu. the extractable spin power thus becomes comparable to Lg. te. Ly~Ly/3."," Suppose that the maximum possible jet power is constrained by the Eddington accretion limit $L_E 
+\sim 10^{46} (M_{\rm BH}/10^8 M_{\odot})$ erg/s. The maximum spin power, for example, which can be extracted from the rotational energy of a BH (via a Blandford-Znajek-type process) is on the order of $L_s 
+\sim 10^{44} (M_{\rm BH}/10^8 M_{\odot})^2 (B_H/10^4 \mathrm{G})^2$ erg/s. With an Eddington equipartition magnetic field strength $B_H \propto 1/\sqrt{M_{\rm BH}}$ , the extractable spin power thus becomes comparable to $L_E$, i.e., $L_s \sim L_E/3$."
+" The average VHE gamma-ray flux detected by H.E.S.S. implies an isotropic. VHE luminosity Li,(8—10)x1076 erg/s (Aharonian et al.", The average VHE gamma-ray flux detected by H.E.S.S. implies an isotropic VHE luminosity $L_{\rm iso} \sim (5-10) \times 10^{46}$ erg/s (Aharonian et al.
+ 2007)., 2007).
+" If the radiation is produced within two. oppositely directed. conical jets of an opening angle 8.=0.1 rad. the required (minimum) lubelin.arM, power would be L;~(1-2)Eddineton107(9/0.1X erg/s. For a BH Jetacereting at close to the limit(see below) or with maximal spin power. the required minimum mass thus becomes 10(Gir)?.. which would be consistent with the variability constraints (Eq.??))"," If the radiation is produced within two, oppositely directed, conical jets of an opening angle $\theta \lppr 0.1$ rad, the required (minimum) jet power would be $L_j \sim (1-2)\times10^{45} (\theta/0.1)^2$ erg/s. For a BH accreting at close to the Eddington limit (see below) or with maximal spin power, the required minimum mass thus becomes 10^7, which would be consistent with the variability constraints (Eq. \ref{m_var}) )"
+ and be suggestive of a binary! mass ratio η»0.01. In reality. the variability constraint of Eq. (22))," and be suggestive of a binary mass ratio $m_2/m_1 \simeq M_{\rm BH}/M_t > 0.01$ In reality, the variability constraint of Eq. \ref{m_var}) )"
+ is probably too optimistic. in. particular if one wishes to account for the second finding (11) (e.. the log-normal VHE flux distribution) based on fluctuations in the accretion flow that feeds the jet.," is probably too optimistic, in particular if one wishes to account for the second finding (ii) (i.e., the log-normal VHE flux distribution) based on fluctuations in the accretion flow that feeds the jet."
+ Consider again the fluctuating disk model (Lyubarskit 1997). where fluctuations of the disk parameters atsome radius. which oceur on local viscous timescale. can lead to variations 1n the aecretion rate at smaller radii that are of the flicker- or red- type.," Consider again the fluctuating disk model (Lyubarskii 1997), where fluctuations of the disk parameters at some radius, which occur on local viscous timescale, can lead to variations in the accretion rate at smaller radii that are of the flicker- or red-noise type."
+ Then the relevant timescale to be employed ts the, Then the relevant timescale to be employed is the
+assuming isotropic scattering by dust grains.,assuming isotropic scattering by dust grains.
+" The calculation also assumes that the dust follows a Gaussian radial density distribution in a shell of radius Ry. and with a FWHM of AR,details).."," The calculation also assumes that the dust follows a Gaussian radial density distribution in a shell of radius $R_{\rm{sh}}$, and with a FWHM of $\Delta R_{\rm{sh}}$."
+ Following the procedure of?.. we have performed a simple analysis to calculate the total dust mass. My. of the observed structures.," Following the procedure of, we have performed a simple analysis to calculate the total dust mass, $M_{\rm{d}}$, of the observed structures."
+" If an optically thin dust envelope is assumed. the scattered flux. Εις. is given by integrating the product of the stellar flux (at the observed distance from the star). Fy/4«R. the number of scatterers. NM... and the scattering cross section. σι. over the observed wavelength range: Εις and F, were measured in the total-intensity images."," If an optically thin dust envelope is assumed, the scattered flux, $F_{\rm{sc}}$, is given by integrating the product of the stellar flux (at the observed distance from the star), $F_{\star}/4\pi R^2$ , the number of scatterers, $N_{\rm{sc}}$ , and the scattering cross section, $\sigma_{\rm{sc}}$ , over the observed wavelength range: $F_{\rm{sc}}$ and $F_{\star}$ were measured in the total-intensity images."
+ Fy was estimated by fitting à Gaussian to the star and summing all counts within 3o., $F_{\star}$ was estimated by fitting a Gaussian to the star and summing all counts within $\sigma$.
+ This was then corrected for the 5-magnitude dampening ofthe mask., This was then corrected for the 5-magnitude dampening ofthe mask.
+" The scattering cross sectio for spherical grains is given by where Q.. is the scattering efficiency and ais thegran radius. for simplicity assumed to be constant at jum. The ratio between Fy. and F, gives Ν.Α through Eq. 7.."," The scattering cross section for spherical grains is given by where $Q_{\rm{sc}}$ is the scattering efficiency and $a$ is thegrain radius, for simplicity assumed to be constant at $\mu$ m. The ratio between $F_{\rm{sc}}$ and $F_{\star}$ gives $N_{\rm{sc}}$ through Eq. \ref{F}."
+ The mass of each grain. my. is given by the volume and density of the dust grains.," The mass of each grain, $m_{\rm{d}}$, is given by the volume and density of the dust grains."
+ For the silicate and carbon grains we assumed typical values for the grain density of 3 anc 2eecem™. respectivelyrespectively).," For the silicate and carbon grains we assumed typical values for the grain density of 3 and $^{-3}$ , respectively."
+". The total dust mass. My. is then giver by M,xmay."," The total dust mass, $M_{\rm{d}}$, is then given by $N_{\rm{sc}}\times m_{\rm{d}}$."
+ For W ΑΙ. astronomical silicates with optical constants from were used to calculate Ος..," For W Aql, astronomical silicates with optical constants from were used to calculate $Q_{\rm{sc}}$."
+ For the carbor stars. U Camand DR Ser. amorphous carbon grains from were used.," For the carbon stars, U Camand DR Ser, amorphous carbon grains from were used."
+ The estimated total dust mass depends strongly on the assumed grain size. My«a/Q. where O.«a?. and 2€B<3 In the wavelength range of our observations.," The estimated total dust mass depends strongly on the assumed grain size, $M_{\rm{d}}\propto a/Q_{\rm{sc}}$, where $Q_{\rm{sc}}\propto a^{\beta}$, and $2\lesssim\beta\lesssim3$ in the wavelength range of our observations."
+ The assumption of a constant grain size further adds to the uncertainty in the mass estimate., The assumption of a constant grain size further adds to the uncertainty in the mass estimate.
+ In addition. there are a number of uncertainties and simplifications in this method and the estimates are therefore order-of-magnitude estimates.," In addition, there are a number of uncertainties and simplifications in this method and the estimates are therefore order-of-magnitude estimates."
+ To study the circumstellar dust distribution around W Agl the R-band images acquired with the coronographic mask were used., To study the circumstellar dust distribution around W Aql the R-band images acquired with the coronographic mask were used.
+ Already in the total-intensity image (Fig. 3.. /eff)," Already in the total-intensity image (Fig. \ref{deg}, )"
+ the brightness distribution around W Aql appears asymmetric., the brightness distribution around W Aql appears asymmetric.
+ The scattered light is more intense to the south-west (SW) where the emission extends out to about from the center. compared to about in the north and the east directions.," The scattered light is more intense to the south-west (SW) where the emission extends out to about from the center, compared to about in the north and the east directions."
+ The polarization angle and the degree of polarization are shown as polarization vectors in Fig., The polarization angle and the degree of polarization are shown as polarization vectors in Fig.
+ 3. (ight)., \ref{deg} ).
+ The vectors are overlaid on an image showing the polarized intensity., The vectors are overlaid on an image showing the polarized intensity.
+ Figure 3. (right) shows that the polarizing dust is distributed all around the star. however. the SW enhancement appears clearly.," Figure \ref{deg} ) shows that the polarizing dust is distributed all around the star, however, the SW enhancement appears clearly."
+ The image has been smoothed by a 3x3-pixel Gaussian to reduce the noise., The image has been smoothed by a $\times$ 3-pixel Gaussian to reduce the noise.
+ The integrated polarization degree across the image (corresponding to what would be measured if the source was unresolved) is about 10%., The integrated polarization degree across the image (corresponding to what would be measured if the source was unresolved) is about .
+.. In the SW quadrant the mean polarization degree is around across the feature., In the SW quadrant the mean polarization degree is around across the feature.
+ The maximum polarization degree is found to be just above in the SW part of the image., The maximum polarization degree is found to be just above in the SW part of the image.
+ Figure 4+ left) shows the polarized intensity within of W Agl., Figure \ref{prof} ) shows the polarized intensity within of W Aql.
+ To investigate the distribution and extent of the SW feature. the image was divided into four quadrants: north-east (NE). north-west (NW). south-east (SE). and south-west (SW).," To investigate the distribution and extent of the SW feature, the image was divided into four quadrants: north-east (NE), north-west (NW), south-east (SE), and south-west (SW)."
+ By calculating AARPs of the differentquadrants. we can compare the brightness distribution across the image (Fig. 4.. right).," By calculating AARPs of the differentquadrants, we can compare the brightness distribution across the image (Fig. \ref{prof}, )."
+ The image tis clearly brighter in the SW. while the other three quadrants look similar.," The image is clearly brighter in the SW, while the other three quadrants look similar."
+ A 4th degree polynomial (dashed line. Fig. 4.. right)," A 4th degree polynomial (dashed line, Fig. \ref{prof}, )"
+ is well-fitted to the log-log AARP of the three weaker-emission quadrants (solid line. Fig. [4.. right).," is well-fitted to the log-log AARP of the three weaker-emission quadrants (solid line, Fig. \ref{prof}, ),"
+ and by subtracting the fit from the polarized-intensity image. the location and extent of the SW asymmetry was found (Fig. [4.. right).," and by subtracting the fit from the polarized-intensity image, the location and extent of the SW asymmetry was found (Fig. \ref{prof}, )."
+ The area close to the mask («4/55 offset) and the outer parts offset) were not taken into account in the fitting., The area close to the mask $<$ 5 offset) and the outer parts $>$ offset) were not taken into account in the fitting.
+" The SW brightness enhancement seems to start already at the edge of the mask and it is nearly constant out to about7"".", The SW brightness enhancement seems to start already at the edge of the mask and it is nearly constant out to about.
+". It then declines until it disappears at approximately12"".", It then declines until it disappears at approximately.
+. Figure 4. left) shows the image after the fit to the weaker-emission quadrants has been subtracted., Figure \ref{prof} ) shows the image after the fit to the weaker-emission quadrants has been subtracted.
+ W Agl was also observed without the coronograph both in the R- and V-band., W Aql was also observed without the coronograph both in the R- and V-band.
+ The polarized-intensity images are shown in Fig. 5..," The polarized-intensity images are shown in Fig. \ref{nomask},"
+ overlaid by polarization vectors., overlaid by polarization vectors.
+ In the R-band image the star was saturated and any structures in the very inner parts of the image can therefore not be trusted., In the R-band image the star was saturated and any structures in the very inner parts of the image can therefore not be trusted.
+ This also explains the read-out tracks seen in Fig., This also explains the read-out tracks seen in Fig.
+ 5. (/efr)., \ref{nomask} ).
+ The polarized intensity of the V-band image isshown in Fig., The polarized intensity of the V-band image isshown in Fig.
+ 5. (right)., \ref{nomask} ).
+ To interpret the very inner structure is not straightforward as the secondary star contributes significantly in the V-band and the binary system is not resolved., To interpret the very inner structure is not straightforward as the secondary star contributes significantly in the V-band and the binary system is not resolved.
+ Both images in Fig., Both images in Fig.
+ 5 show the same polarization pattern as can be seen in Figs 3.. and 4. and the SW feature is also clearly confirmed.," \ref{nomask} show the same polarization pattern as can be seen in Figs \ref{deg}, and \ref{prof} and the SW feature is also clearly confirmed."
+ The amount of dust in the SW brightness enhancement was calculated assuming optically thin dust scattering (Sect. 5.2))., The amount of dust in the SW brightness enhancement was calculated assuming optically thin dust scattering (Sect. \ref{ss:caldus}) ).
+ By measuring the ratio between the scattered flux and the stellar flux in the total-intensity image. M. was obtained according to Eqn. 7..," By measuring the ratio between the scattered flux and the stellar flux in the total-intensity image, $N_{\rm{sc}}$ was obtained according to Eqn. \ref{F}."
+ The scattered flux was estimated byadding all counts in the SW quadrant of the image from 4’ (to avoid distortions due to the edge ofthemask) outto 12., The scattered flux was estimated byadding all counts in the SW quadrant of the image from $\arcsec$ (to avoid distortions due to the edge ofthemask) outto $\arcsec$ .
+" The stellar flux. F,. was estimated as described in Sect. 5.2.."," The stellar flux, $F_{\star}$ , was estimated as described in Sect. \ref{ss:caldus}. ."
+ The flux ratio was found to be Γι Γι=3x 107. and the dust mass was found to be My ~ 1«107° MMo.," The flux ratio was found to be $F_{\rm{sc}}$ $F_{\star} = 3 \times 10^{-3}$ , and the dust mass was found to be $M_{\rm{d}}$ $\sim$ $1 \times 10^{-6}$$_{\sun}$ ."
+ Assuming a ratio of 12x107 (2).. this corresponds to a total mass of 107 MMo.," Assuming a dust-to-gas ratio of $\times 10^{-3}$ , this corresponds to a total mass of $^{-3}$ $_{\sun}$ ."
+From (he observed vector magnetic field as described in Section 2.. we exirapolate to obtain three (vpes of coronal magnetic field. each of which has the vertical component of the maenetic [field imposed at the photopshere: In order to have energv values which can be compared. we have imposed (he same closed conditions on tlie side and top boundaries for each model.,"From the observed vector magnetic field as described in Section \ref{sec:obs}, we extrapolate to obtain three types of coronal magnetic field, each of which has the vertical component of the magnetic field imposed at the photopshere: In order to have energy values which can be compared, we have imposed the same closed conditions on the side and top boundaries for each model."
+ To satisfv these conditions. we surround (he vector magnetic field observed by MEES/IVM by weak field measurements provided bx SOIIO/MDI Bne-ol-sight observations.," To satisfy these conditions, we surround the vector magnetic field observed by MEES/IVM by weak field measurements provided by SOHO/MDI line-of-sight observations."
+ The active region fields are then confined by a surrounding potential field and (he magnetic field decreases [rom the center of the active region (compatible with the field vanishing at infinity)., The active region fields are then confined by a surrounding potential field and the magnetic field decreases from the center of the active region (compatible with the field vanishing at infinity).
+ The magnetic flux is balanced., The magnetic flux is balanced.
+ From the 3D coronal magnetic configurations. we can derive (he magnetic energy. [or the clifferent active regions and different models: in a volume O.," From the 3D coronal magnetic configurations, we can derive the magnetic energy for the different active regions and different models: in a volume $\Omega$."
+ The free magnetic energy is derived from the nonlinear (»Jff) lorce-DIree [lield using either the potential (pof) or linear force-Iree (/ff) field as reference field:, The free magnetic energy is derived from the nonlinear ) force-free field using either the potential ) or linear force-free ) field as reference field:
+After the detection of 6 high-frequency peaked BL Lac objects (HIDEs) with air Ceerenkov telescope facilities. the field of extragalactic GeV TeV astronomy is currently one of the most rapidly expauding research areas in astrophysics.,"After the detection of 6 high-frequency peaked BL Lac objects (HBLs) with ground-based air Čeerenkov telescope facilities, the field of extragalactic GeV – TeV astronomy is currently one of the most rapidly expanding research areas in astrophysics."
+ The steadily improving [lux sensitivities of the new generation of air Ceerenkov telescope arrays (IXonopelkoetal.2002) ancl their decreasing energv thresholds. provides a growing," The steadily improving flux sensitivities of the new generation of air Čeerenkov telescope arrays \citep{konopelko99,weekes02} and their decreasing energy thresholds, provides a growing"
+"value is indistinguishable from that eiven above for spiral spheroids. although is probably an overestimate since it includes a simall nunaber of chigh-5,"" ealaxics located in the cores of the Virgo aud Fornax clusters.","value is indistinguishable from that given above for spiral spheroids, although is probably an overestimate since it includes a small number of $S_n$ ” galaxies located in the cores of the Virgo and Fornax clusters."
+ The spirals. bv contrast. are almost invariably located in loose eroups.," The spirals, by contrast, are almost invariably located in loose groups."
+ In order to gauge the possible importance of local environment. we have calculated for cach object the local ealaxy density. py. using the ΠΟ of galaxies brightcr than Mp=16 coutained within a shell of radius 0.5 Alpe using the Nearby Galaxy. Catalog of Tully (1985).," In order to gauge the possible importance of local environment, we have calculated for each object the local galaxy density, $\rho_0$, using the number of galaxies brighter than $M_B = -16$ contained within a shell of radius 0.5 Mpc using the Nearby Galaxy Catalog of Tully (1988)."
+" The 29 git clliptical galaxies have 0.1<py9.2 I whereas the 11 spiral ealaxies span the range 0,0<<3.7 72"," The 29 giant elliptical galaxies have $0.1 \lae \rho_0 \lae 9.2$ $^{-3}$, whereas the 11 spiral galaxies span the range $0.0 \lae \rho_0 \lae 3.7$ $^{-3}$."
+ A total of 17 eunt elliptical galaxies have pyzi3.7. aud are thus located iu cuvirommenuts of conrparable deusitv.," A total of 17 giant elliptical galaxies have $\rho_0 \lae 3.7$, and are thus located in environments of comparable density."
+" The mean specific frequency of these galaxies is $,,=3.902.5. which is indistinguishable from that found for the spiral spheroids."," The mean specific frequency of these galaxies is ${\overline{S_n}} = 3.9\pm2.5$, which is indistinguishable from that found for the spiral spheroids."
+ The dramatic recent increase in the number of galaxies having accurately measured globular cluster metallicity distributions is due primarily to the use of metallicitysensitive color indices (6.g.. Ostrov. Geisler Forte 1993) and the excellent dnaging capabilities of the Wubble Space Telescope (e.g... Neilsen Tsvatanoy 1999: (ιομάνα Wissler-Patig 1999: Kumi 1999).," The dramatic recent increase in the number of galaxies having accurately measured globular cluster metallicity distributions is due primarily to the use of metallicity-sensitive color indices $e.g.$, Ostrov, Geisler Forte 1993) and the excellent imaging capabilities of the Hubble Space Telescope $e.g.$, Neilsen Tsvatanov 1999; Gebhardt Kissler-Patig 1999; Kundu 1999)."
+ These studies lave demonstrated that the majority of eiaut elliptical ealaxics show evidence for the preseuce of multiple. cheumicallv-distinct globular cluster populations.," These studies have demonstrated that the majority of giant elliptical galaxies show evidence for the presence of multiple, chemically-distinct globular cluster populations."
+ This imultiauodalitv is not unique to giaut elliptical ealaxies: the Calactic globular cluster svstoeia has long con known to show a bimodal distribution in metallicity. showing distinct peaks at |Fo/I]] ~1.6 aud —0.6. (," This multi-modality is not unique to giant elliptical galaxies: the Galactic globular cluster system has long been known to show a bimodal distribution in metallicity, showing distinct peaks at [Fe/H] $\simeq -1.6$ and $-0.6$. ("
+Morezu 1959: αμα 1959: Zinn 1985).,Morgan 1959; Kinman 1959; Zinn 1985).
+ Aualvsis of he elobular cluster svstem associated with MI suggests hat it too appears significantly iuultiauodal (Ashman Bird 1993: Barby et al., Analysis of the globular cluster system associated with M31 suggests that it too appears significantly multi-modal (Ashman Bird 1993; Barmby et al.
+ 1999)., 1999).
+ Unfortunately. the siuuple of spiral galaxies having accurate globular cluster uetallicity distributions ids limited to just these two objects. although recent work has demonstrated that oinodal elobulay cluster metallicity distributions are observed in SO ealaxies as well (INissler-Patig et al.," Unfortunately, the sample of spiral galaxies having accurate globular cluster metallicity distributions is limited to just these two objects, although recent work has demonstrated that bimodal globular cluster metallicity distributions are observed in S0 galaxies as well (Kissler-Patig et al."
+ L997): Kundu Whitinore 1998)., 1997b; Kundu Whitmore 1998).
+ At present. the ouly galaxies which show no evidence of multianodal globular cluster metallicity distributions are dE/dSphl galaxies of low aud intermediate huninositv (50. Mym— l1)," At present, the only galaxies which show no evidence of multi-modal globular cluster metallicity distributions are dE/dSph galaxies of low and intermediate luminosity $i.e.$, $M_V \gae -17$ )."
+ Such objects contain only metal-poor clusters (see Figure 1 of Cotté et al., Such objects contain only metal-poor clusters (see Figure 1 of Côtté et al.
+ 1998)., 1998).
+ All elobular cluster svstenis studied to date have radial profiles which are simular to. or shallower than. that of the uuderling ealaxy light (Iluris 1986: Iissler-Patie 1997: Duell et al.," All globular cluster systems studied to date have radial profiles which are similar to, or shallower than, that of the underlying galaxy light (Harris 1986; Kissler-Patig 1997; Durrell et al."
+ 1996)., 1996).
+ As Terris (1991) poiuts out. no case has been found iu which the elobular clusters svstenà is more centrally concentrated than the galaxy itself.," As Harris (1991) points out, no case has been found in which the globular clusters system is more centrally concentrated than the galaxy itself."
+ The well-studied Virgo eiaut elliptical galaxies M19 and AIS? are known to contain two chemically-distinct elobular cluster populations: in both cases. the metalrich clusters appear to be more centrally concentrated than their metalpoor counterparts (Geisler et al.," The well-studied Virgo giant elliptical galaxies M49 and M87 are known to contain two chemically-distinct globular cluster populations; in both cases, the metal-rich clusters appear to be more centrally concentrated than their metal-poor counterparts (Geisler et al."
+ 1996: Neilsen. Tsvatanoy Ford 1998: Lec. Ini Ceisler 1998).," 1996; Neilsen, Tsvatanov Ford 1998; Lee, Kim Geisler 1998)."
+ In the case of M9. Lee et al. (," In the case of M49, Lee et al. ("
+1997) report a 6-0 difference in the measured density profile slopes for the imietalaich aud metal-poor elobular clusters (although sce MeLaughliu 1999 for a dissenting view).,1997) report a $\sigma$ difference in the measured density profile slopes for the metal-rich and metal-poor globular clusters (although see McLaughlin 1999 for a dissenting view).
+ Du addition. Lee et al. (," In addition, Lee et al. ("
+1998) note that the metalaich clusters trace the underlving galaxy light in both radial profile aud. ellipticitv. whereas the metalpoor clusters comprise a more extended and spherically distributed population.,"1998) note that the metal-rich clusters trace the underlying galaxy light in both radial profile and ellipticity, whereas the metal-poor clusters comprise a more extended and spherically distributed population."
+ At the present time. the only spiral galaxy which has been studied im sufficient detail that it is possible to investigate reliably the separate distributions of metal-rich and metal-poor clusters is the Milkv Way.," At the present time, the only spiral galaxy which has been studied in sufficient detail that it is possible to investigate reliably the separate distributions of metal-rich and metal-poor clusters is the Milky Way."
+" The Galactic elobular cluster system as a whole is known to obey a threc-dimensional distribution of the form p,~7 where my,23.5 (Ibis 1976).", The Galactic globular cluster system as a whole is known to obey a three-dimensional distribution of the form ${\rho}_h \sim R_G^{-m_h}$ where $m_h \simeq 3.5$ (Harris 1976).
+" This is in close agreement with the values of im,~3.0δν derived from halo RR Lyrac and blue horizoutal brauch stars (Saha 1985: Preston. Scheetiman Beers 1991)."," This is in close agreement with the values of $m_h \simeq 3.0 - 3.5$ derived from halo RR Lyrae and blue horizontal branch stars (Saha 1985; Preston, Schectman Beers 1991)."
+ Iis (1976) noted. however. that uo metalrich clusters are found bevond Ree~T kpc. and suggested that most of these clusters are associated with the ceutral bulee of the Galaxy. cousisteut with the Ποιος of several receut studies (Miuniti 1995: Barbuy. Bica Oxrtolami: Cotté 1999: cf," Harris (1976) noted, however, that no metal-rich clusters are found beyond $R_G \sim 7$ kpc, and suggested that most of these clusters are associated with the central bulge of the Galaxy, consistent with the findings of several recent studies (Minniti 1995; Barbuy, Bica Ortolani; Côtté 1999; cf."
+ Ziun 1985)., Zinn 1985).
+" For conrparisou. the deusitv profile of the Galactic bulge is roughly py~Re!"" where 3.65«ois<L2 (CToerndrup 1988. Dlauco Terudrup 1989)."," For comparison, the density profile of the Galactic bulge is roughly ${\rho}_b \sim R_G^{-m_b}$ where $3.65 \lae m \lae 4.2$ (Terndrup 1988, Blanco Terndrup 1989)."
+" Moreover. the mean Calactocentrie radius of the metaliich cluster sample discussed by Cotté (1999) is Rep=3.2£2.0 kpe. which is xinilar to the effective radius of the Galactic bulee CR,=2.7 kpe: Ciliiore. Ning vau der I&kxuit 1990)."," Moreover, the mean Galactocentric radius of the metal-rich cluster sample discussed by Côtté (1999) is ${\overline{R_G}} = 3.2\pm2.0$ kpc, which is similar to the effective radius of the Galactic bulge $R_{e} = 2.7$ kpc; Gilmore, King van der Kruit 1990)."
+ Iu sununary. the Galactic globular cluster system. like those of all giant elliptical galaxies studied to date. has a spatial distributions which is comparable to. or slightlv more extended. than that of the unuderbliug spheroicd helt.," In summary, the Galactic globular cluster system, like those of all giant elliptical galaxies studied to date, has a spatial distributions which is comparable to, or slightly more extended, than that of the underlying spheroid light."
+ Moreover. the metalaich globular clusters are more centrally concentrated than their metal-poor counterparts.," Moreover, the metal-rich globular clusters are more centrally concentrated than their metal-poor counterparts."
+ Given these sinulavitics. it is natural to ask if a sinele model can explain the observed properties of elobular clusters in both eiaut elliptical aud spiral galaxies.," Given these similarities, it is natural to ask if a single model can explain the observed properties of globular clusters in both giant elliptical and spiral galaxies."
+ For instance. models which seek to explain the formation of eiut elliptical galaxies and their associated svstenis of auctalrich elobular clusters iu spiralespiral merecrs (ug Sclaviezer 1986: Ashinan Zepf 1992) suffer from the obvious difficulty that the spirals themselves appear to show multi-modal globular clusters metallicity distributions.," For instance, models which seek to explain the formation of giant elliptical galaxies and their associated systems of metal-rich globular clusters in spiral-spiral mergers $e.g.$, Schwiezer 1986; Ashman Zepf 1992) suffer from the obvious difficulty that the spirals themselves appear to show multi-modal globular clusters metallicity distributions."
+ Tn the hierarchical picture of Cotté et al (1998). the metalaich elobular clusters in giant elliptical galaxies represent the clusters mitiallv associated with the most niassive prote-galactic fragment.," In the hierarchical picture of Côtté et al (1998), the metal-rich globular clusters in giant elliptical galaxies represent the clusters initially associated with the most massive proto-galactic fragment."
+ By contrast. the metapoor globular clusters now associated with the giaut elliptical galaxy are identified as those which have becu acquired during the accretion and disruption of microtiy. dwarf galaxies aud proto-galactic fragiments (460.. fant systems which are known to coutaiu only metal-poor clusters).," By contrast, the metal-poor globular clusters now associated with the giant elliptical galaxy are identified as those which have been acquired during the accretion and disruption of numerous dwarf galaxies and proto-galactic fragments $i.e.$, faint systems which are known to contain only metal-poor clusters)."
+ The relative number of metal-rich aud meta elobular clusters in the giant elliptical galaxy is then assumed to reflect the Iuuinosities (or. alternatively. masses) of the dominant proto-galactic building bloc- and the acereted population of smaller proto-galactic," The relative number of metal-rich and metal-poor globular clusters in the giant elliptical galaxy is then assumed to reflect the luminosities (or, alternatively, masses) of the dominant proto-galactic building block and the accreted population of smaller proto-galactic"
+and star-formation in the host galaxy.,and star-formation in the host galaxy.
+ Attributing the submm-Iuminosity of high-recdshift quasars to star-formation herefore accords with the speculations about the origins of he subnim sources., Attributing the submm-luminosity of high-redshift quasars to star-formation therefore accords with the speculations about the origins of the submm sources.
+ 3elore one can estimate physical quantities. such as dust temperature and star-formation rate. it is necessary o characterise the form of the dust-emission spectrum of hese objects.," Before one can estimate physical quantities, such as dust temperature and star-formation rate, it is necessary to characterise the form of the dust-emission spectrum of these objects."
+ It is also important to quantify variations of he SED with object class. luminosity or redshift. otherwise. if one adopts a single template for all sources. it is not egitimate to adopt. more than one derived parameter. as an indicator of physical status.," It is also important to quantify variations of the SED with object class, luminosity or redshift, otherwise, if one adopts a single template for all sources, it is not legitimate to adopt more than one derived parameter as an indicator of physical status."
+ In fitting thermal SEDs. lowever. we are constrained by the paucity of photonietric data. and often it is inappropriate to fit both the emissivity and the temperature as free parameters. when we have only hree or four photometric points per source.," In fitting thermal SEDs, however, we are constrained by the paucity of photometric data, and often it is inappropriate to fit both the emissivity and the temperature as free parameters, when we have only three or four photometric points per source."
+ Previous work. which considered object-by-object temperature fits. adopted an assumed value for the emissivity index (Bulleyal... in ep.).," Previous work, which considered object-by-object temperature fits, adopted an assumed value for the emissivity index (Buffey, in prep.)."
+ In contrast. the purposes of this Letter are. first. to investigate the possibility of co-adcding data from a sample of luminous. 2z4 quasars. to find a self-consistent. overall rest-Lit temperature ancl emissivity: and to characterise the uncertainties involved in doing so. whether they are due to experimental errors or to. physical dillerences.," In contrast, the purposes of this Letter are, first, to investigate the possibility of co-adding data from a sample of luminous, $z\ga4$ quasars, to find a self-consistent, overall best-fit temperature and emissivity; and to characterise the uncertainties involved in doing so, whether they are due to experimental errors or to physical differences."
+ Finally. we discuss applications of the mean SED. and try to justify the inferred parameters in terms of a simple evolutionary model.," Finally, we discuss applications of the mean SED, and try to justify the inferred parameters in terms of a simple evolutionary model."
+ VPhroughout. we assume a .A-clominateck cosmology 120.3. O4—0.7. £y=65 ," Throughout, we assume a $\Lambda$ -dominated cosmology $\Omega_M$ =0.3, $\Omega_{\Lambda}$ =0.7, $H_0$ $^{-1}$."
+The zc4 datasets are from our work with LAA and JCNUE (SCUBAyam...jim.. aancd jm: UII jmn)) as DEMin companion papers (Isaak 1994: Omontal... 1996a: Bulleal... in prep.):," The $z>4$ datasets are from our work with IRAM ) and JCMT (SCUBA, and ; UKT14 ) as described in companion papers (Isaak 1994; Omont, 1996a; Buffey, in prep.);"
+ and CCSO photometryv is. from Benford (1999), and CSO photometry is from Benford (1999).
+ Additionallv. ου include the 5=X91 quasar | 5255) (Lewis 1998).," Additionally, we include the $z=3.91$ quasar $+$ 5255 (Lewis 1998)."
+ ALL the sources have been detected at least four cillerent wavelengths., All the sources have been detected at least four different wavelengths.
+ An important caveat is that the quoted statistical uncertainties do not rellect the systematic errors due to Παν calibration., An important caveat is that the quoted statistical uncertainties do not reflect the systematic errors due to flux calibration.
+ 'lherefore. we assume calibration errors. combined in quacrature with the random uncertainties. of20 percent.," Therefore, we assume calibration errors, combined in quadrature with the random uncertainties, of 20 percent."
+" The rest-frame Far-infrared.(PLR) SED can be described by a ΠούΠο. blackbody. spectrum. assuming that the grains are in equilibrium. with the radiation field ancl that. the is optically-thin at submm. wavelengths: S,cxdust. cloudNote. however. that some ultraluminous starbursts (such as Arp220: Rowan-Robinson Efstathiou. 1993) are optically thick at the shortest rest-wavelengths sampled by the present cata (605)."," The rest-frame far-infrared (FIR) SED can be described by a modified blackbody spectrum, assuming that the grains are in equilibrium with the radiation field and that the dust cloud is optically-thin at submm wavelengths: $S_{\nu} \propto \frac{\nu^{3+\beta}}{e^{\frac{h\nu}{kT}}-1}.$ Note, however, that some ultraluminous starbursts (such as Arp220; Rowan-Robinson Efstathiou, 1993) are optically thick at the shortest rest-wavelengths sampled by the present data $\mu$ m)."
+ TEhis method determines the most ellectiveparameterisation of the lone-wavelcneth emission. a simple. convenient description that is frequently adopted Benford et al..," This method determines the most effective parameterisation of the long-wavelength emission, a simple, convenient description that is frequently adopted Benford et al.,"
+. 199!* Lisenfeld. Isaak Llills. 0: Dunne et al..," 1999; Lisenfeld, Isaak Hills, 2000; Dunne et al.,"
+ 2000) though in reality. a distribution of dust temperatures is likely.," 2000), though in reality, a distribution of dust temperatures is likely."
+ Shorter wavelength radiation. in the near-to-mid infrared region (10 rrest-[rame). is likely to be contaminated by the hotter dust hat absorbs the harder continuum clue to the AGN itself ltowan-IHobinson. 1995).," Shorter wavelength radiation, in the near-to-mid infrared region $\sim$ rest-frame), is likely to be contaminated by the hotter dust that absorbs the harder continuum due to the AGN itself Rowan-Robinson, 1995)."
+ The contribution [rom star ormation. on the other hand. emerges in the submillimetre region.," The contribution from star formation, on the other hand, emerges in the submillimetre region."
+ To assess the contamination of the submillimetre uminosity due to the AGN component requires better mid-infrared constraints than exist for all but a few high-redshif quasars (see. for example. Itowan-Itobinson. 2000).," To assess the contamination of the submillimetre luminosity due to the AGN component requires better mid-infrared constraints than exist for all but a few high-redshift quasars (see, for example, Rowan-Robinson, 2000)."
+ We herefore simply assume that the component extracted. by our fitting procedure is appropriate for the extrapolation of a far-infarect luminosity (Lip) from which a star formation rate is to be calculated., We therefore simply assume that the component extracted by our fitting procedure is appropriate for the extrapolation of a far-infared luminosity $L_{\rm FIR}$ ) from which a star formation rate is to be calculated.
+ We caution. however. that. the unforeseen presence of a hot component. could skew the fi to higher temperatures.leading us to overestimate the star formation rate.," We caution, however, that the unforeseen presence of a hot component could skew the fit to higher temperatures,leading us to overestimate the star formation rate."
+ Adopting ZY and 3 as free parameters. we employ a ο procedure to determine the best. single-component fit to the data. which is coadded: in ᾱ ποconsistent wav.," Adopting $T$ and $\beta$ as free parameters, we employ a $\chi^2$ -minimisation procedure to determine the best single-component fit to the data, which is coadded in a self-consistent way."
+ X subtlety arises. because [or à given observed: wavelength. the range of objects! redshifts causes us to sample their spectra at. different rest. wavelengths.," A subtlety arises because for a given observed wavelength, the range of objects' redshifts causes us to sample their spectra at different rest wavelengths."
+ Therefore. if we choose to normalise the data at a given rest wavelength. a spectrum must be assumed. so that an extrapolation to this reference from the nearest data point can be performed.," Therefore, if we choose to normalise the data at a given rest wavelength, a spectrum must be assumed so that an extrapolation to this reference from the nearest data point can be performed."
+ The best-fit spectrum thus determined. then itself becomes the template for the extrapolation. and the process is iterated.," The best-fit spectrum thus determined then itself becomes the template for the extrapolation, and the process is iterated."
+ In. practice. this procedure converges to à solution rapidly without roaming around parameter space) so for these datasets ab least. we conclude that the overnlé x7-mininiised. [it is. incecd. the best self-consistent solution.," In practice, this procedure converges to a solution rapidly without roaming around parameter space) so for these datasets at least, we conclude that the $\chi^2$ -minimised fit is, indeed, the best self-consistent solution."
+increments of a monotonic function of the form where (he ος are appropriate non-negative constants.,increments of a monotonic function of the form where the $c$ 's are appropriate non-negative constants.
+ Near (he centre (he requirement of equal increments of f. will lead to equal increments of à=m?77., Near the centre the requirement of equal increments of $f$ will lead to equal increments of $x=m^{2/3}$.
+ The second term of f. will force equal increments of the hydrogen mass fraction where Xj; changes rapidly (at an shell)., The second term of $f$ will force equal increments of the hydrogen mass fraction where $X_H$ changes rapidly (at an H-burning shell).
+" The third will lead to equalsteps of Inp towards the surface. where i/Azz1 anc Nyy, is uniform: and the last term will cause a fine subdivision around 2—e;z20.000. where the opacity varies rapidly over several orders of magnitude."," The third will lead to equalsteps of $\ln p$ towards the surface, where $m/M\approx 1$ and $X_H$ is uniform; and the last term will cause a fine subdivision around $T=c_4\approx 20,000 K$, where the opacity varies rapidly over several orders of magnitude."
+ The variables (5.m.L.p.T.3;) ave represented by arrays over a grid of 7=I.....n. where i=1 corresponds to the centre. aud /=n {ο the surface (photosphere).," The variables $(s,m,L,\rho,T,p,Y_j)$ are represented by arrays over a grid of $i=1,\dots,n$, where $i=1$ corresponds to the centre, and $i=n$ to the surface (photosphere)."
+ Thus eqs., Thus eqs.
+ become the difference equations There is one such pair of equations for each ¢=2..... n.," become the difference equations There is one such pair of equations for each $i=2,\dots,n$ ."
+" Together with the boundary conditions s,=0 and (Ape),=(1—E,)g,7.. these add up to 20 equations."," Together with the boundary conditions $s_1=0$ and $(\kappa p_G)_n=(1-\Gamma_n)g_n\tau_s$, these add up to $2n$ equations."
+ The variables V. L and Fj. related to the energy and. particle fluxes. are replaced. by arravs (hal refer to (hemidpoints /x1/2.," The variables $\nabla$, $L$ and $F_j$, related to the energy and particle fluxes, are replaced by arrays that refer to the $i\pm 1/2$."
+ Thus eqs., Thus eqs.
+ become where. in (he last pair of equations. we have suppressed (he index j that refers to the nuclear species.," become where, in the last pair of equations, we have suppressed the index $j$ that refers to the nuclear species."
+ The coellicients V;;2 anda;joa are evaluated by using the arithmetic means of the 5eric-point arguments.5 for example r;iosài=(r;PPL424Hjr;))/2.," The coefficients $\nabla_{i-\frac{1}{2}}$ and $\sigma_{i+\frac{1}{2}}$ are evaluated by using the arithmetic means of the grid-point arguments, for example $r_{i-\frac{1}{2}}=(r_{i-1}+r_i)/2$."
+ Again.S there is one eq.," Again, there is one eq."
+" for each 7=2.....n. which. together with the boundary condition L,=4arzoT}. brings the number of equations up to 32."," for each $i=2,\dots,n$, which, together with the boundary condition $L_n=4\pi r_n^2\sigma T_n^4$, brings the number of equations up to $3n$ ."
+ Furthermore. there is one set of eqs.," Furthermore, there is one set of eqs."
+ and for each 7=1.....n(and for each one of the species).," and for each $i=1,\dots,n$(and for each one of the species)."
+ WJ is the number of, If $J$ is the number of
+campaign observing sctups and data reduction. flux calibration of the spectra. aud intercalibration of the data sets to form the single set of optical continu and eenission liue light curves shown for each object iu Figure 2..,"campaign observing setups and data reduction, flux calibration of the spectra, and intercalibration of the data sets to form the single set of optical continuum and emission line light curves shown for each object in Figure \ref{fig:lc4cc}."
+ Table 1 displays basic statistical parameters describing the final light curves shown in Figure 2.., Table \ref{tab:lcstats} displays basic statistical parameters describing the final light curves shown in Figure \ref{fig:lc4cc}.
+ Colma 1 eives the object. and Column 2 lists the spectral feature represented by each light cuve.," Column 1 gives the object, and Column 2 lists the spectral feature represented by each light curve."
+ The uuuber of data points in each lieht curve is shown iu Column 3. with the median sampling iuterval between these data poiuts eiven in Column 1.," The number of data points in each light curve is shown in Column 3, with the median sampling interval between these data points given in Column 4."
+ Cohuun 5 shows the mean fractional error in the faxes of cach time series., Column 5 shows the mean fractional error in the fluxes of each time series.
+ Colima 6 eives the excess variance. calculated as where o? is the variance of the observed fluxes. 6? is their mean square wucertainty. aud Cf? is the mean of the observed fixesLO97).," Column 6 gives the excess variance, calculated as where $\sigma^2$ is the variance of the observed fluxes, $\delta^2$ is their mean square uncertainty, and $\langle f \rangle$ is the mean of the observed fluxes."
+. Colum 7 is the ratio of the iiaxinumun to minima fux in the light curves. and Colum & gives the adopted mean ttime lag aud uncertainties determined through the primary time series analysis which utilized the full Ine profile aud is described iu detail for cach object by DOO).," Column 7 is the ratio of the maximum to minimum flux in the light curves, and Column 8 gives the adopted mean time lag and uncertainties determined through the primary time series analysis which utilized the full line profile and is described in detail for each object by D09b."
+ The lag iieasureimeuts between the continuum aud eenüssion listed in Table d. represent the average time delay across the BER. measured from the centroid of the cross correlation fiction produced during the time series analysis of the continua aud full line profile light. curves shown in Figure 2. (see Dü09a for details).," The lag measurements between the continuum and emission listed in Table \ref{tab:lcstats} represent the average time delay across the BLR, measured from the centroid of the cross correlation function produced during the time series analysis of the continuum and full line profile light curves shown in Figure \ref{fig:lc4cc} (see D09a for details)."
+ Here. we focus ou the velocitv-xresolved time series analysis we performed on cach target to investigate the potential for recovering velocity-depeudent time delays across the ecluission line m order to inter the kinematic structure of the line eittiug gas.," Here, we focus on the velocity-resolved time series analysis we performed on each target to investigate the potential for recovering velocity-dependent time delays across the emission line in order to infer the kinematic structure of the line emitting gas."
+ We divided the ecluission line into eight velocityv-space bins. whose boundaries were determined by the division of the ruis spectrmm of each object iuto cight bins of equal flux. as depicted in the top panels of Figure 23...," We divided the emission line into eight velocity-space bins, whose boundaries were determined by the division of the rms spectrum of each object into eight bins of equal flux, as depicted in the top panels of Figure \ref{fig:reslags}."
+ Compared to the line boundaries used for the full profile analysis leading to the Πο curves shown iu Figure 2.. those used for this analysis were slightly narrowed in the cases of 33516 and 33227 in order to mclude oulv the most variable portions of the line profile. aud boundarics were broadened for 55515 because the rus spectrum shows variability in the red wwing that extends beneath the ΠΙΑ1959. narrow cluission line.," Compared to the line boundaries used for the full profile analysis leading to the light curves shown in Figure \ref{fig:lc4cc}, those used for this analysis were slightly narrowed in the cases of 3516 and 3227 in order to include only the most variable portions of the line profile, and boundaries were broadened for 5548 because the rms spectrum shows variability in the red wing that extends beneath the $\lambda 4959$ narrow emission line."
+ Light curves were created from measurements of the inteerated fflux iu each bin aud then cross correlated with the contimmun light curves shown in Figure 2.., Light curves were created from measurements of the integrated flux in each bin and then cross correlated with the continuum light curves shown in Figure \ref{fig:lc4cc}.
+ The bottom panels of Fieure 3) show the lag micasurcments for cach of these bius., The bottom panels of Figure \ref{fig:reslags} show the lag measurements for each of these bins.
+ Error bars in the velocity direction represent the bin width., Error bars in the velocity direction represent the bin width.
+ The evidence for a velocity-stratified BLR respouse to σολπα variations is clear in all three cases., The evidence for a velocity-stratified BLR response to continuum variations is clear in all three cases.
+ Iuterestingly. cach case demonstrates a cifferent kinematic signature: (1) outflow is indicated In 33227. given the generally lounger lags frou shifted BLR eas compared to the eas on the bluc-shifted side of the line. (2) 33516 shows the opposite siguature. with the blue side of the line lagging the red side an indication that there is an infall component," Interestingly, each case demonstrates a different kinematic signature: (1) outflow is indicated in 3227, given the generally longer lags from red-shifted BLR gas compared to the gas on the blue-shifted side of the line, (2) 3516 shows the opposite signature, with the blue side of the line lagging the red side — an indication that there is an infall component"
+clip evens bevoud 3 7 Ton fre fit aud refit. uutil eve ids clipped.,"clip events beyond 3 $\sigma$ from the fit and refit, until no event is clipped."
+ We simulate I:vevents with a Caussian SCitter 0.15 nag arouid the Wubble line., We simulate Ia events with a Gaussian scatter 0.15 mag around the Hubble line.
+ For Ibe events. it Is uncear how appli18o to hem the brightcr-slower ald briieshter-bluer corrections for SNe Ia will affect t1ο alsolute magnitude «istribution.," For Ibc events, it is unclear how applying to them the brighter-slower and brighter-bluer corrections for SNe Ia will affect the absolute magnitude distribution."
+ We inieht eness that SOLe par of their brigitress scatter is due to extinctkna iu heir host galaxy. aid that |xiehter-bluer correctiols would narrow their disALCO DOCuhs distributio1.," We might guess that some part of their brightness scatter is due to extinction in their host galaxy, and that brighter-bluer corrections would narrow their distance modulus distribution."
+ We will COisider below two estiliaes of he The 1»xiehtness scatter. O1e as observed aud oue sieuificaurlv ower.," We will consider below two estimates of the Ibc brightness scatter, one as observed and one significantly lower."
+ Note that ? estimate mtrisie πο]ude scatter correcte for host ealaxv extiuction which «litter little from raw estimates.," Note that \cite{Richardson02}
+ estimate intrinsic magnitude scatter corrected for host galaxy extinction which differ little from raw estimates."
+ The contamunation depends ou the xieht eud of the The hnuunositv function which is uot well constrained: 7 propose a distribution extending well bevoud the Ta average brightuess. while ? brigitest The event ls 0.5 mag fainter than the average Ia. Both results are however compatible at the —10 CL. given the modest statisticsinvolved!.," The contamination depends on the bright end of the Ibc luminosity function which is not well constrained: \cite{Richardson02} propose a distribution extending well beyond the Ia average brightness, while \cite{Li10II} brightest Ibc event is 0.5 mag fainter than the average Ia. Both results are however compatible at the $\sim$ 10 CL, given the modest statistics."
+". We propose three scenarios for the Thc population. cach with two values for σι. detailed iu Table Lt: Contamination affects cosmologv bv biasine the average distance liio111,"," We propose three scenarios for the Ibc population, each with two values for $\sigma_{bc}$, detailed in Table \ref{tab:contamination}: Contamination affects cosmology by biasing the average distance modulus."
+ Ilowever. a redshift independent bias has no ciffecf on cosmologw. aud we report in Table |. the slope of the bias as a function of redshift dp/dz. which fairly escribes most of the effect.," However, a redshift independent bias has no effect on cosmology, and we report in Table \ref{tab:contamination} the slope of the bias as a function of redshift $d\,\delta\mu/dz$, which fairly describes most of the effect."
+ We find that uuder the τος scenarios. the effect on cosmology is small.," We find that under the three scenarios, the effect on cosmology is small."
+ Even in «»r “LILO scenario. where the coΓαλΓΙΟ. ds suspiciOlsly laree compared to the fractio1 of The ideutified by veh redshit SNe τα surveys. the effect on cosmologv renans below the systematic U1certainty σολ)=0.01 (defined later in Eq. 5 )," Even in our “L10” scenario, where the contamination is suspiciously large compared to the fraction of Ibc identified by high redshift SNe Ia surveys, the effect on cosmology remains below the systematic uncertainty $\sigma(e_M)=0.01$ (defined later in Eq. \ref{eq:M_depend_z}) )"
+ that we will consider as a |selije., that we will consider as a baseline.
+ The cliping process has a neeheible impact on the SNe Ta statistics., The clipping process has a negligible impact on the SNe Ia statistics.
+ A rough estimae of the size of the «feet we fiid can be readilv conued. for e.g. the first ine of the able: the average offset ¢ft Toc within the Ia ce wiow is Q.017. the fractio1 of the Ibe population within he same window amounts to ~20%.. he Ihe to Ia uuuber ratio at D-0 js asstuned to be 0.)6. and dtomereases with redshüft as {1|2).," A rough estimate of the size of the effect we find can be readily computed, for e.g. the first line of the table: the average offset of Ibc within the Ia $\pm 3\sigma$ window is 0.047, the fraction of the Ibc population within the same window amounts to $\sim$, the Ibc to Ia number ratio at $z=0$ is assumed to be 0.16, and it increases with redshift as $\sim(1+z)$."
+ The evoluion of the Ibe 1iduced distance bias reads 0.047xL20.16=0.0015., The evolution of the Ibc induced distance bias reads $0.047*0.2*0.16 = 0.0015$.
+ One mieht aree that a Crassi distribution is inadequate escrif the tails of he Toc distribution. but more populate tails also exhiat shallower slopes. aud both alterations lave a tendencvw to cancel cach other.," One might argue that a Gaussian distribution is inadequate to describe the tails of the Ibc distribution, but more populated tails also exhibit shallower slopes, and both alterations have a tendency to cancel each other."
+" Note that our iniates ignore the ]20chtial rejection from Πο curve soes and colours oThe events, which is certaiily a ¢“OUSCEVATIVe assuniptiou."," Note that our estimates ignore the potential rejection from light curve shapes and colours of Ibc events, which is certainly a conservative assumption."
+ Οιr results wielt look at odds wit1 other attempts., Our results might look at odds with other attempts.
+ 7. follows a siuilar approach and finds sizabe effects ol cosmologv. aud we attribute the difference to he asence of clipping.," \cite{Homeier05}
+ follows a similar approach and finds sizable effects on cosmology, and we attribute the difference to the absence of clipping."
+ More recently. ? xoposcetd a supernova photometric iceutificatioi challenge.," More recently, \cite{SnIdChallenge} proposed a supernova photometric identification challenge."
+ The Danned simulated sample contaius a laree fraction of events for which the phase coverage aux he signal to noise ratio are inadequate for a distance measurement. would hey be eeuuine Type Ta. Iun our sinitation. such events are excluded a priori aud do not couit as identification zilures.," The provided simulated sample contains a large fraction of events for which the phase coverage and the signal to noise ratio are inadequate for a distance measurement, would they be genuine Type Ia. In our simulation, such events are excluded a priori and do not count as identification failures."
+ We also benefit in our siniulaion from assuniug spectroscopic redshifts as this reduces t1e Thibble diagram scatter., We also benefit in our simulation from assuming spectroscopic redshifts as this reduces the Hubble diagram scatter.
+ Reearding the identification 1sing colours. Type Ta events were ecnerated in? with a colour smeariug of 0.1 mac. which we now know to be abott L times too large. ancl leads 0 overestimating mnis-idenutifications.," Regarding the identification using colours, Type Ia events were generated in \cite{SnIdChallenge} with a colour smearing of 0.1 mag, which we now know to be about 4 times too large, and leads to overestimating mis-identifications."
+ We eventually lenore the efficicuevloss due to photometric selection. auc integrate the effects of possible contanuiualous muto a drift of absohte magnitude (see Section 1.3.1 aud Eq. 5 )).," We eventually ignore the efficiencyloss due to photometric selection, and integrate the effects of possible contaminations into a drift of absolute magnitude (see Section \ref{sec:distance_estimation_parameters}
+ and Eq. \ref{eq:M_depend_z}) )."
+ Note that if spectroscopic redshifts o: host galaxies are acquired aud the wrong host IS assieue to a supernova. this eveut will likely fail the photometric typing cuts aud hence will not pollute the IIubble diagram.," Note that if spectroscopic redshifts of host galaxies are acquired and the wrong host is assigned to a supernova, this event will likely fail the photometric typing cuts and hence will not pollute the Hubble diagram."
+" The wide survey provides the possibility of coustructing a rest frame Lbaud Uubble diagrain out to : 0.9. with a statistics of about του events,"," The wide survey provides the possibility of constructing a rest frame I-band Hubble diagram out to $z\simeq 0.9$ , with a statistics of about 7000 events."
+ Encouraging, Encouraging
+null hypothesis (u— 0) cannot be rejected.,null hypothesis $u = 0$ ) cannot be rejected.
+" In other words, in this scenario inclusion of the gas inflow is insignificant."," In other words, in this scenario inclusion of the gas inflow is insignificant."
+" So, below we discuss two models - M20 and M7b."," So, below we discuss two models - M20 and M7b."
+ Let us compare them and test the null hypothesis that the corresponding variances do not differ significantly against the alternative one that they do., Let us compare them and test the null hypothesis that the corresponding variances do not differ significantly against the alternative one that they do.
+" For this, according to Martin (1971) we should check if the value F=(minew2o/mincyr)? lies in the range: F«F5 (ais the significance level and we adopt that the 1/F,,5degrees of freedom are approximately the same in the both alternatives)."," For this, according to Martin (1971) we should check if the value $F=(min \, \sigma_{M20}/min \, \sigma_{M7b})^2$ lies in the range: $1/F_{\alpha/2} < F < F_{\alpha/2}$ $\alpha$ is the significance level and we adopt that the degrees of freedom are approximately the same in the both alternatives)."
+" From Table 1 we derive Fz&1.35, whereas fora=0.02 Foj2=F(18,18,0.99)z3.14 (Draper Smith 1981)."," From Table 1 we derive $F \approx 1.35$, whereas for $\alpha = 0.02$ $F_{\alpha/2} = F(18,18,0.99) \approx 3.14$ (Draper Smith 1981)."
+" Hence, the above inequality is satisfied and we conclude: the difference between models M20 and M7b is undetected statistically (on the basis of available observational data of oxygen radial distribution and rather simple theoretical approach)."," Hence, the above inequality is satisfied and we conclude: the difference between models M20 and M7b is undetected statistically (on the basis of available observational data of oxygen radial distribution and a rather simple theoretical approach)."
+" At the moment, the only way ato discriminate between them is as it was discussed in Sec."," At the moment, the only way to discriminate between them is as it was discussed in Sec."
+ 4., 4.
+" Finally, our estimates for the target parameters are as follows: 1)M20: Qp= 33.2t1.4kms! ! (the corotation radius ro—7.02:0.3 km s! Κρε-1): B=0.016+0.001 Gyr."," Finally, our estimates for the target parameters are as follows: 1): $\Omega_P = 33.2 \pm 1.4$ km $^{-1}$ $^{-1}$ (the corotation radius $r_C = 7.0 \pm 0.3$ km $^{-1}$ $^{-1}$ ); $\beta = 0.016 \pm 0.001$ Gyr."
+ 2)M7b: Np=30.5:1 kms! kpc! (the corotation radius ro=7.6133 km s! kpc 8=0.036+0.002 Gyr., 2): $\Omega_P = 30.5 \pm 1$ km $^{-1}$ $^{-1}$ (the corotation radius $r_C = 7.6^{+0.3}_{-0.2}$ km $^{-1}$ $^{-1}$ ); $\beta = 0.036 \pm 0.002$ Gyr.
+ In Fig., In Fig.
+ 3 the best theoreticalD); oxygen distributions for the above two models and the distributions for the upper and lower values of Op and B are shown., 3 the best theoretical oxygen distributions for the above two models and the distributions for the upper and lower values of $\Omega_P$ and $\beta$ are shown.
+ From the figure it is seen that our theory explains the bimodal radial distribution of oxygen in the disc of our Galaxy rather well., From the figure it is seen that our theory explains the bimodal radial distribution of oxygen in the disc of our Galaxy rather well.
+" In the present paper, a new method for evaluation of the corotation radius in the Galaxy is developed."," In the present paper, a new method for evaluation of the corotation radius in the Galaxy is developed."
+ Our approach is based on a statistical analysis of the bimodal structure of oxygen radial distribution in the galactic disc determined over Cepheids., Our approach is based on a statistical analysis of the bimodal structure of oxygen radial distribution in the galactic disc determined over Cepheids.
+ By means of treatment of observational data we derive that the corotation resonance happens to be situated at rc.~7.0—7.6 kpc depending on the rate of intergalactic gas infall on to the galactic disc (the statistical error is ~0.3 kpc)., By means of treatment of observational data we derive that the corotation resonance happens to be situated at $r_C \sim 7.0 - 7.6$ kpc depending on the rate of intergalactic gas infall on to the galactic disc (the statistical error is $\sim$ 0.3 kpc).
+ The above value for the corotation radius is close to the solar galactocentric distance το=7.9 kpc., The above value for the corotation radius is close to the solar galactocentric distance $r_0 = 7.9$ kpc.
+ Simultaneously the constant for the rate of oxygen synthesis in the galactic disc 6~0.016—0.036 Gyrs was obtained., Simultaneously the constant for the rate of oxygen synthesis in the galactic disc $\beta \sim 0.016 - 0.036$ Gyrs was obtained.
+" We also argue in favour of a short time-scale formation of the galactic disc, namely: ty~2 Gyrs."," We also argue in favour of a short time-scale formation of the galactic disc, namely: $t_f \sim 2$ Gyrs."
+" This scenario enables to solve the problem of the lack of intergalactic gas infall, i.e. the very low present integrated rate of gas infall on to the disc, M4~0.1—0.2 Mo year! observed by Sancisi Fraternali (2008) and Bregman (2009), whereas the integrated star formation rate is expected to be M,~1—2 Me year! (Robitaille Whitney 2010)."," This scenario enables to solve the problem of the lack of intergalactic gas infall, i.e. the very low present integrated rate of gas infall on to the disc, $\dot M_f \sim 0.1 - 0.2$ $_{\odot}$ $^{-1}$ observed by Sancisi Fraternali (2008) and Bregman (2009), whereas the integrated star formation rate is expected to be $\dot M_s \sim 1 - 2$ $_{\odot}$ $^{-1}$ (Robitaille Whitney 2010)."
+" Higher level of the current global star formation rate, M.~5 Mo year""!, needs too high rate of gas infall on to the galactic disc or unrealistically high gaseous and stellar current densities at the solar galactocentric distance."," Higher level of the current global star formation rate, $\dot M_s \sim 5$ $_{\odot}$ $^{-1}$, needs too high rate of gas infall on to the galactic disc or unrealistically high gaseous and stellar current densities at the solar galactocentric distance."
+" Authors thank to referee for vary important comments, to Profs."," Authors thank to referee for vary important comments, to Profs."
+ Yu., Yu.
+Shchekinov and A.Zasov for helpful discussions.,Shchekinov and A.Zasov for helpful discussions.
+ The work was supported in part by grants No., The work was supported in part by grants No.
+ 02.740.0247 and P685 of Federal agency for science and innovations., 02.740.0247 and P685 of Federal agency for science and innovations.
+" IAA thanks to the Russian funds for basic research, grant No."," IAA thanks to the Russian funds for basic research, grant No."
+ 11-02-90702., 11-02-90702.
+tvpe binaries can lead to svstematic shifts in the atmospheric parameters (e.g.Leberοἱal.Foretal.2010:Müller 2010).,"type binaries can lead to systematic shifts in the atmospheric parameters \citep[e.g.][]{heber04,for10,mueller10}."
+. The equality of the individual SDSS spectra. which cover most of the orbital phase is not good enough (o resolve (his ellect.," The quality of the individual SDSS spectra, which cover most of the orbital phase is not good enough to resolve this effect."
+ The statistical errors are higher (han the expected modulations., The statistical errors are higher than the expected modulations.
+ Drechseletal.(2001). adoptecl errors of 900Ix in ζω ancl 0.1dex in logy to account for this effect in the case of the WWVir tvpe binary 007054-6700.," \citet{drechsel01} adopted errors of $900\,{\rm K}$ in $T_{\rm eff}$ and $0.1\,{\rm dex}$ in $\log{g}$ to account for this effect in the case of the Vir type binary $+$ 6700."
+ Since the orbital period as well as the atmospheric parameters of (his binary are similar to (he ones of JOS8205-+0008. we adopt similar uncertainties here.," Since the orbital period as well as the atmospheric parameters of this binary are similar to the ones of J08205+0008, we adopt similar uncertainties here."
+ Systematic errors introduced by different model grids are typically smaller (han that etal. 2007).," Systematic errors introduced by different model grids are typically smaller than that \citep{lisker05,geier07}."
+. While the inclination of the svstem is well constrained bv the light. curve analvsis. there remains a degeneracy between the masses and (he radii of the components.," While the inclination of the system is well constrained by the light curve analysis, there remains a degeneracy between the masses and the radii of the components."
+ The surface gravitv of the sclB determined in the quantitative spectroscopic analvsis provides an additional constraint since it only depends on the mass aud the radius of (he subdwarl., The surface gravity of the sdB determined in the quantitative spectroscopic analysis provides an additional constraint since it only depends on the mass and the radius of the subdwarf.
+ To proceed. however. we need to constrain the sdB mass from evolutionary models.," To proceed, however, we need to constrain the sdB mass from evolutionary models."
+ In Fig., In Fig.
+ 3 we compare the position of the star in the (Tar. logg)-plane to other IW. Vir stars as well as to (wo sels of models.," \ref{tefflogg} we compare the position of the star in the $T_{\rm eff}$, $\log{g}$ )-plane to other HW Vir stars as well as to two sets of models."
+ The first one represents the canonical picture of EIID evolution. while the second one recalls post-RGB evolution. which means that the sdB star has left the RGB early and cid not ignite helium in the core at all.," The first one represents the canonical picture of EHB evolution, while the second one recalls post-RGB evolution, which means that the sdB star has left the RGB early and did not ignite helium in the core at all."
+ According to Fig., According to Fig.
+ 3. the star is situated on the Extreme Horizontal Braneh (EIB) consistent with being a core helium-burning star as are the other WW Vir stars (except AA Dor)., \ref{tefflogg} the star is situated on the Extreme Horizontal Branch (EHB) consistent with being a core helium-burning star as are the other HW Vir stars (except AA Dor).
+ Since the orbital period is short. it was formed via common envelope ejection.," Since the orbital period is short, it was formed via common envelope ejection."
+" Population svuthesis models (lanetal.2002.2003). predict a mass range of O48A/. with a sharp peak at 0.47AM, for sdDs in binaries formed in this wav."," Population synthesis models \citep{han02,han03} predict a mass range of $M_{\rm sdB}=0.37-0.48\,M_{\rm \odot}$ with a sharp peak at $0.47\,M_{\rm \odot}$ for sdBs in binaries formed in this way."
+ Even lower sdD masses (down to 0.3 M.) are possible. when a more massive progenitor star (2—3. ) ignites core helium-burning under non-degenerate conditions.," Even lower sdB masses (down to $0.3\,M_{\rm \odot}$ ) are possible, when a more massive progenitor star $2-3\,M_{\rm \odot}$ ) ignites core helium-burning under non-degenerate conditions."
+ Since this formation channel is predicted to be rare. (he results of the quantitative spectroscopic analysis are filly consistent with an EIID model of the canonical mass.," Since this formation channel is predicted to be rare, the results of the quantitative spectroscopic analysis are fully consistent with an EHB model of the canonical mass."
+ Hot subcdwarf masses derived [rom asteroseismic analvses (e.g.vanGrooteletal.2010). or [rom analyses of eclipsing binaries are in general agreement wilh (his picture., Hot subdwarf masses derived from asteroseismic analyses \citep[e.g.][]{vangrootel10} or from analyses of eclipsing binaries are in general agreement with this picture.
+ In spite of the consistency discussed. above. it may still be premature to adopt. the canonical mass.," In spite of the consistency discussed above, it may still be premature to adopt the canonical mass."
+ It has been pointed out that the sdD stars in DDor. 223334-3927 and 1188112 might not burn helium in their cores and may (herelore be of lower mass (Rauch 2003).," It has been pointed out that the sdB stars in Dor, $+$ 3927 and 188112 might not burn helium in their cores and may therefore be of lower mass \citep{rauch00,heber04,heber03}."
+. Such close binaries are expected to form whenever the RGB evolution is interrupted by the ejectüón of a common envelope before the core has reached (he mass required for helium ignition., Such close binaries are expected to form whenever the RGB evolution is interrupted by the ejection of a common envelope before the core has reached the mass required for helium ignition.
+sslew survey (Readetal.2005)854021.,slew survey \citep[][]{xmmslew05}.
+.. The source was detected with the European Photon Lhmaging Camera (EPIC) PN instrument at a count of rate2.24270.53countss+ (0.212 keV). which converts into a 0.5.10 keV unabsorbed [ux of (icl)10Icreemο7s (again using with Ng=O382l2 ↓∪≼↛⊔↓⋜⋯∠⇂⋜↧↓≻∪∖∖⊽∢⋅↓⋅↓⋜↧∖∖⊽↓⊔∠⇂∢⊾⇀∖↓∶−≽⋅⇀∫≻↴⊳ : .−⋅ ," The source was detected with the European Photon Imaging Camera (EPIC) PN instrument at a count rateof $2.24 \pm 0.53~\cnts$ (0.2–12 keV), which converts into a 0.5–10 keV unabsorbed flux of $\sim(4 \pm 1) \times 10^{-11}~\flux$ (again using with $N_{\mathrm{H}}=9.3 \times 10^{21}~\mathrm{cm}^{-2}$ and a powerlaw index $\Gamma=2.3$ )."
+The corresponding NIE count rate is 0.10countss , The corresponding /XRT count rate is $\sim 0.40 \pm 0.10~\cnts$ .
+Ilies within the 3 aremin error box of the unclassified. hard X-ray source citepsee.Fig. 3: this coincidence was also noted]rodriguezü9. which appears in the aall-skyv survey catalogue (Ixrivonos 2010)..," lies within the 3 arcmin error box of the unclassified hard X-ray source \\citep[see Fig.~\ref{fig:images}; this coincidence was also noted, which appears in the all-sky survey catalogue \citep{krivonos07,bird09}."
+ We used the publicly available dedata to search for N-rav bursts from the location of3539., We used the publicly available data to search for X-ray bursts from the location of.
+. Phis region has been covered by regular observations of the ssatellito (Winkleretal.2003) since the beginning of 2003. in particular at. low energy (320 keV) with the Joint European X-ray. Monitor (JIZM-N:Lundetal.2003).. module |. and 2. and at high energies (17.100 keV) with the SSolt Gamma-ray lmager (ISCGRI:Lebrunetal.2003).. mounted on the Imager onDoard the SSatellite (IBIS:Ubertinietal.2003).," This region has been covered by regular observations of the satellite \citep{winkler2003} since the beginning of 2003, in particular at low energy (3–20 keV) with the Joint European X-ray Monitor \citep[JEM-X;][]{lund2003}, module 1 and 2, and at high energies (17–100 keV) with the Soft Gamma-ray Imager \citep[ISGRI;][]{lebrun2003}, mounted on the Imager onBoard the Satellite \citep[IBIS;][]{ubertini2003}."
+. The data are clivicled into individual pointings called. Science. Windows (ScW). hemselves grouped into revolutions of the satellite.," The data are divided into individual pointings called Science Windows (ScW), themselves grouped into revolutions of the satellite."
+ wwas not pointing towards the source field when the X-ray ourst. picked up by BAT occurred., was not pointing towards the source field when the X-ray burst picked up by BAT occurred.
+ In the archival public data. there are 7359 [BIS ScW yetween revolutions 37 and 674. pointing less than 12 degrees from the source. anc 650 JIZM-N. ScW between revolutions 46 and 0601. pointing less than 3.5 degrees from he source.," In the archival public data, there are 7359 IBIS ScW between revolutions 37 and 674, pointing less than 12 degrees from the source, and 650 JEM-X ScW between revolutions 46 and 661, pointing less than 3.5 degrees from the source."
+ These data are spread over a time range of ive vears. from 2003 February 1 to 2008 April 20. for cllective exposures of 16 and 0.76 Ms for 1BIS and. JIM-X respectively.," These data are spread over a time range of five years, from 2003 February 1 to 2008 April 20, for effective exposures of 16 and 0.76 Ms for IBIS and JEM-X respectively."
+ Phe dillerence of exposure is due to the [act hat IBIS has a Larger field of view than JIZM-N. and thus 1appened to observe mmore often., The difference of exposure is due to the fact that IBIS has a larger field of view than JEM-X and thus happened to observe more often.
+ We have analysed this data set with the standard Ollline Science Analysis software (OSA: v. 7.0). distributed by the SScience Data Center (ISDC:Courvoisieretal.2003). and based. on algorithms described in Goldwurmetal.(2003) for LBIS and Westergaardetal.(2003). Lor JEM-X. The total collapsed. mosaic of the LBES imagesreveals a weak but significant (7.70) excess at the position of the source.," We have analysed this data set with the standard Offline Science Analysis software (OSA; v. 7.0), distributed by the Science Data Center \citep[ISDC;][]{courvoisier2003} and based on algorithms described in \citet{goldwurm2003} for IBIS and \citet{westergaard2003} for JEM-X. The total collapsed mosaic of the IBIS imagesreveals a weak but significant $\sigma$ ) excess at the position of the source."
+ lis Hux in the 1740 keV band is 410.4?creem?s.5., Its flux in the 17–40 keV band is $\sim 4 \times 10^{-12}~\flux$.
+ We have searched for X-ray bursts in the IBIS data with the BBurst Alert Svstem (IBAS:Alercehettietal.2003).. vet no X-ray burst was detected.," We have searched for X-ray bursts in the IBIS data with the Burst Alert System \citep[IBAS;][]{mereghetti2003}, yet no X-ray burst was detected."
+ We have also explored the JIZM-N data. more suitable to look for such events since these are usually soft.," We have also explored the JEM-X data, more suitable to look for such events since these are usually soft."
+ However. again. no X-ray burst was found.," However, again, no X-ray burst was found."
+" We investigated the Swiff//BAT transient monitor results of J1735.. provided. by the ΟΛΗ team""."," We investigated the /BAT transient monitor results of , provided by the /BAT ."
+.. No other X-ray bursts are detected with a limiting llux of ~1.4 (1550 keV) for a single pointing (which, No other X-ray bursts are detected with a limiting flux of $\sim1.4 \times 10^{-9}~\flux$ (15–50 keV) for a single pointing (which
+results. but for sheets and filaments respectively.,"results, but for sheets and filaments respectively."
+ Similar (trends are seen in all three panels. indicating that sheets and filaments grow hierarchically in much the same wav as halos.," Similar trends are seen in all three panels, indicating that sheets and filaments grow hierarchically in much the same way as halos."
+ These panels show that. at 2=0. more (han half of the cosmic mass is in sheets will masses exceeding LOMAL...," These panels show that, at $z=0$, more than half of the cosmic mass is in sheets with masses exceeding $10^{13}M_{\odot}$."
+ This fraction is about 30% for filaments in the same mass range. and is only 13% for halos.," This fraction is about $30\%$ for filaments in the same mass range, and is only $13\%$ for halos."
+ At the present time. more than 99% of the cosmic mass is contained in sheets more massive than LOMM... in contrast to virialized halos which. in this mass range. contain only about half of the cosmic mass.," At the present time, more than $99\%$ of the cosmic mass is contained in sheets more massive than $10^{10}M_{\odot}$, in contrast to virialized halos which, in this mass range, contain only about half of the cosmic mass."
+ Al z~B3. more than of the cosmic mass was already assembled in sheets more massive than that of the Milky Way halo (~LOMAL. ).," At $z\sim 3$, more than of the cosmic mass was already assembled in sheets more massive than that of the Milky Way halo $\sim 10^{12}M_\odot$ ),"
+To illustrate our argument we take an isothermal sphere with a core radius as our lens model (Scheinderetal.1992).,To illustrate our argument we take an isothermal sphere with a core radius as our lens model \cite{Sch92}.
+. Ehe lens parameters are then given by where d=(1|(rfro)7) and r is the distance to the lens centre., The lens parameters are then given by where $d=(1+(r/r_0)^{2})$ and $r$ is the distance to the lens centre.
+ In addition to the values of ry. ancl Py. the coordinates of the centre of svnunetry. introduce two further parameters.," In addition to the values of $r_0$, and $\Phi_0$, the coordinates of the centre of symmetry, introduce two further parameters."
+ The centre of the lens can be inferred in various wavs (see for example (Ixochanek 1990))) and we shall assume in the rest of this article that it is already known., The centre of the lens can be inferred in various ways (see for example \cite{Koch90}) ) and we shall assume in the rest of this article that it is already known.
+" We have used a lens with ro=1 and d,=0.9 (Fig.3)).", We have used a lens with $r_0=1$ and $\Phi_0=0.9$ \ref{sphere_iso}) ).
+ If we assume values for the pair of parameters (ry.994). and the image parameters are assumed to have been determined exactiv. then it is possible to reconstruct the parameters for the source.," If we assume values for the pair of parameters $(r_0,\Phi_0)$, and the image parameters are assumed to have been determined exactly, then it is possible to reconstruct the parameters for the source."
+ I£. we assume also that the only information that we have about the source is from the measurement of the polarization. which provides us with the orientation of the source galaxies with a standard. error of e. then we can find the lens parameters that best fit the observed. polarization direction.," If we assume also that the only information that we have about the source is from the measurement of the polarization, which provides us with the orientation of the source galaxies with a standard error of $\epsilon$, then we can find the lens parameters that best fit the observed polarization direction."
+ The natural way to do this is to carry out a chi-squared minimisation. i.e. we minimise the following expression over the lens paranietCrs.," The natural way to do this is to carry out a chi-squared minimisation, i.e. we minimise the following expression over the lens parameters."
+ The summation in equation (16)) is carried out over the set of galaxies whose polarization has been determined., The summation in equation \ref{chi}) ) is carried out over the set of galaxies whose polarization has been determined.
+ a) is the direction which has been reconstructed from the parameters of the lens model. and à. is the measured direction.," $\alpha_{s}^{r}$ is the direction which has been reconstructed from the parameters of the lens model, and $\alpha_{s}^{m}$ is the measured direction."
+ We have reconstructed lens parameters for the cases when the number of galaxies. N. for which the polarization has been measured is given by IN= 50. 10 and 5.," We have reconstructed lens parameters for the cases when the number of galaxies, $N$, for which the polarization has been measured is given by $N=50$ , $10$ and $5$."
+ For cach of these cases we consider the error on the polarization direction. e. to be given by 2. 5. and. 4n order to get an idea of the conlidence one should have in the values of ry and do obtained in this wav. we have for. each pair of values (GN.e). generated 500 dillerent samples of galaxies. and from these used the chi-sequare estimator of the values of (ro.Po).," For each of these cases we consider the error on the polarization direction, $\epsilon$, to be given by 2, 5, and .In order to get an idea of the confidence one should have in the values of $r_0$ and $\Phi_0$ obtained in this way, we have for, each pair of values $(N,\epsilon)$, generated $500$ different samples of galaxies, and from these used the chi-square estimator of the values of $(r_0,\Phi_0)$."
+ Figures 4 and 5. show ellipses in the (ry.Po) plane which contain TOM and 90% of the reconstructed values.," Figures \ref{pol70} and \ref{pol90} show ellipses in the $(r_0,\Phi_0)$ plane which contain $70\%$ and $90\%$ of the reconstructed values."
+ In order to quantify the information that the measurement of the polarization provides. we have also carried. out a reconstruction of the lens (that is determined rg and $0) assuming no knowledge of the orientation of the source galaxies.," In order to quantify the information that the measurement of the polarization provides, we have also carried out a reconstruction of the lens (that is determined $r_0$ and $\Phi_0$ ) assuming no knowledge of the orientation of the source galaxies."
+ We did this bv comparingthe distribution of source parameters inferred from an assumed. pair of lens parameters and sample of, We did this by comparingthe distribution of source parameters inferred from an assumed pair of lens parameters and sample of
+values of dust lanes and patches. iu contrast with what happens for spiral galaxies.,"values of dust lanes and patches, in contrast with what happens for spiral galaxies."
+ The authors suggestoo that this .lnass discrepancy.» lay be explained. bv the existence. of; a diffuse⋅⋅ component (within⋅⋅⊀↽ 2 kpc from. the ceuter). which is not detectable at optical wavelengths.," The authors suggest that this “mass discrepancy” may be explained by the existence of a diffuse component (within 2 Kpc from the center), which is not detectable at optical wavelengths."
+" Ou the other haud. Tsai Mathews (1996) sugeestedOO that. while. the distributed⋅⋅ dust component is⋅ associated⋅ with⋅⋅⊀↻↥⋅∪⋉∖↥⋅↑↕↸∖↴∖↴≺∐↕≼∐∖↴⋝↥⋅⋜⋯≼↧↓∩≺∖⋅≻⋟∙↖↖⇁↕∐↕↸∖⋜↕⋯↕∪↥⋅↑↸∖∐∏⋉∖↥⋅⋜↧⊓∐⋅↸∖↕↴∖↴ dust recently ejected. .from evolving⋅⋝ stars. another ∙⇁""extra conrponeut of. dust is. present in⋅⋅ ellipticals⋅ both iu≼∐∖↥⋅↕↖↽↸∖≺⊔↥⋅∪⋯↑∐∖⋮∐↴↸∖∐∐↴∖↴↴∖↴↕∪∐∙↕≻↕↕−↥↸∖↥⋅↸∖∐↑⋜⋯↑∐∪↥⋅↴∖↴∙∐∪↖↖⇁↸∖↖↽↸∖↥⋅⋅ dust anes and rines and/or in other ealactic regions."," On the other hand, Tsai Mathews (1996) suggested that, while the distributed dust component is associated with dust recently ejected from evolving stars, another “extra” component of dust is present in ellipticals both in dust lanes and rings and/or in other galactic regions."
+" Iu xwtieulu. they postulated that a substantial nass of cold gas remains ""observationally elusive without forming completely iuto stars”."," In particular, they postulated that a substantial mass of cold gas remains “observationally elusive without forming completely into stars”."
+ Ifthe extra dust is optically thin iu he visible it should be located far from the ealactic core region. where the intensity of the starlight and therefore he erain heating is reduced.," If the extra dust is optically thin in the visible it should be located far from the galactic core region, where the intensity of the starlight and therefore the grain heating is reduced."
+ The dust spatial distributi1 sugeested by GJO5 aud w Tsai Mathews (1996) support the two most popular scenarios respectively: theflow picture and 1 picture., The dust spatial distribution suggested by GJ95 and by Tsai Mathews (1996) support the two most popular scenarios respectively: the picture and the picture.
+. In the sccuario. ⋅⋅ ⋅⋅ ∐∖⋜↧↸∖≼∏⋝⋅↖↽↑∐↸∖↥⋅⋜⊔∐⋜↧↕∪∐∏↸∖∐∙↖↖⇁↕∐↸⊳↕⊔∐⊓∐⋅∐≺∐↰⋉∖∐≼↧↴∖↴∪∐↑↕∐∖ ∐∖↸⊳↕≺∏≼↧↴∖↴∪↕≼∏↴∖↴↑⋜↧∐↧∶↴∙⊾⋜↧↴∖↴↸⊳↿∐⋅↥⋅↸∖∐↑↕∙↖↽∪↴⋝↴∖↴↸∖↥⋅↖↽↸∖≼↧∐⊔∖∐∏≻⊓↸⊳⋜↧↕↴∖↴ ⋯↖↽↸∖⋯⋜↧↕∐↕↖↽⋜↧↕↸∖↘↑↸∖↥⋅∐⋜↧↕∪∏∩⊾↕∐∙↴⋝↸∖↕∐∩⊾⋜↧↴∖↴↴∖↴∪↸⊳↕⋜↧↑↸∖≼↧↑∪↸∖↖↽↸∖∐↑↴∖↴° 1e f galaxy interaction and/or mergers aud being heated by rwermmal conduction iu the hot gas (de Jong et al.," In the scenario the clouds of dust and gas currently observed in ellipticals have mainly an external origin, being associated to events of galaxy interaction and/or mergers and being heated by thermal conduction in the hot gas (de Jong et al."
+ 1990. Sparks et al.," 1990, Sparks et al."
+ 1989)., 1989).
+ On the contrary flow). the oeiternal origin of gas aud dust av be explained with oth red oneiant: winds. (ντα-an ot |al.," On the contrary ), the internal origin of gas and dust may be explained with both red giant winds (Knapp et al."
+ 1992) Dand by the ⋅⋅ ealaxy. is heated byIowan παράπονα explosions aud x& collisions between expanding stellar euvelopes during (T/T jie formation stage. and then cools and condeuses Wt) = al.1991).," 1992) and by the mechanism, in which mass loss from stars within the galaxy is heated by supernova explosions and by collisions between expanding stellar envelopes during the galaxy formation stage, and then cools and condenses (Fabian et al.1991)."
+∠↙∣ UL .( elliptical.n the. presencegit) of dust and=0 the dust “mass discrepancy are related with the dust spatial distribution. which depends ou the nature aud the evolution of these systems.," Therefore, in ellipticals, the presence of dust and the dust “mass discrepancy” are related with the dust spatial distribution, which depends on the nature and the evolution of these systems."
+" Deuce. iu order to investigate the dust content. dust ass evaluations as accurate as possible are required to estimate the amount of the ""ass discrepancy."," Hence, in order to investigate the dust content, dust mass evaluations as accurate as possible are required to estimate the amount of the “mass discrepancy”."
+ Since both scenarios sugecst the presence of different dust componcuts and the very nature of the galactic cuviromments involves the existence of dust grains at differeut. tempcratures. it is necessary to take iuto account a dust temperature distribution for the dust mass evaluations.," Since both scenarios suggest the presence of different dust components and the very nature of the galactic environments involves the existence of dust grains at different temperatures, it is necessary to take into account a dust temperature distribution for the dust mass evaluations."
+ Ίσσα & Xie (1992) sugeest a theoretical approach to take iuto account the effects of the dust temperature distribution iu the dust mass evaluation., Kwan $\&$ Xie (1992) suggest a theoretical approach to take into account the effects of the dust temperature distribution in the dust mass evaluation.
+ I present here an application of their method which is discussed and Huplemented in Section 2., I present here an application of their method which is discussed and implemented in Section 2.
+ The results obtained for the of siuupleellipticals. introducedin Section 3. are prescuted and discussed in Section [| by comparing the FIR with the visual dust mass evaluation and by discussing the correlation between blue aud FIR huuinosities.," The results obtained for the sample of ellipticals, introduced in Section 3, are presented and discussed in Section 4 by comparing the FIR with the visual dust mass evaluation and by discussing the correlation between blue and FIR luminosities."
+ The: dust mass can be derived: both from- optical: aud ποια7 FIR.; observatious., The dust mass can be derived both from optical and from FIR observations.
+. The value of. the mass depends ou the plysical-chemucal properties of the solid particles (1.6. grain radius. erain density aud emissivitv).," The value of the mass depends on the physical-chemical properties of the solid particles (i.e. grain radius, grain density and emissivity)."
+ The usual ⋜∪∐⋅∪⋯⊳∐↸⊳∪∐↴∖↴↕↴∖↴↑↴∖↴∪↕⋟⋜↧↴∖↴↴∖↴⋯⊔↕∐∩⊾↴∖↴∪↕⊔↸∖↖↽⋜↧↕⋯∖↴∖↴↕⋟∪↥⋅↑∐↸∖∩⊾↥⋅⋜↕↕∐Pi ↽∖∙≻⊔⋅ D ⋅ ⋅ ⋅ ≻⋅⋅ um .," The usual approach consists of assuming some values for the grain properties (Hildebrand 1983), while a color temperature is derived from the FIR emission."
+ present slightly. different⋅⋅ foriulae. for. the evaluation. of. the dust mass (Thronson Telesco 1986. Creenhouse et al.," Different authors, however, present slightly different formulae for the evaluation of the dust mass (Thronson Telesco 1986, Greenhouse et al."
+ 1988. Young et al.," 1988, Young et al."
+ 1989. Roberts et al.," 1989, Roberts et al."
+ 1991 aud Thuan Sauvage 1992)., 1991 and Thuan Sauvage 1992).
+ The differcuces between these relations are ouly due to different assumptions on the erain parameters and to different derivations of the color temperature., The differences between these relations are only due to different assumptions on the grain parameters and to different derivations of the color temperature.
+ It should be noticed that. within a flux uncertainty of 1054 Qvhich is typical of the high quality IRAS data). all the dust mass values obtained by using the different formulae are in. agreement.," It should be noticed that, within a flux uncertainty of $\%$ (which is typical of the high quality IRAS data), all the dust mass values obtained by using the different formulae are in agreement."
+! A single temperature model is a rough approximation in describing a ealactic cuviroument., A single temperature model is a rough approximation in describing a galactic environment.
+ The dust is in fact . ⋅⋅ ⋅ ⋅⋅ ⋅ ⋅ ⋅⋅⋅⋅ ↴∖↴≺∏∐⋅↸⊳↸∖↴∖↴∪↕↕∏∐∐∐∪↴∖↴↕⋅↖⇁⋜⋯≼↧∪∐↑↕∐∖∐⋅↴∖↴↻⋜↧↑⋯↕≼∐↴∖↴↑∐↴⋝∏⊓∪⊔∐ Galaxy.," The dust is in fact heated by the radiation field, which in turn depends on the sources of luminosity and on their spatial distribution in the galaxy."
+ The total↴ FIR cutission is⋅⋅⋅ thus Likely due⋅ ≻���∐↸∖ ⋯∐↥⋅∏∏↕↑↕∪∐∪↕⋟≼⊔↕↴∖↴⋜↧↑≼∐∐∎↸∖↥⋅↸∖∐↑↑↸∖∐∏⋉∖↥⋅⋜↧⊓∐⋅↸∖, The total FIR emission is thus likely due to the contribution of dust at different temperatures.
+↴∖↴∙⋀∖↕∪↥⋅↸∖∪↖↽↸∖↥⋅∙ ic TRAS FIR iieasurements are not adequate to detect ie enudissijon connues from cold dust (10—20 EK) which oeaks at wavelengths° between 200 and 300 sau. ↕⋜∥↧∪↻↑↑∐↸∖≼↧∏↴∖↴↑↑↸∖∐∏⋉∖↥⋅⋜↧⊓∐⋅↸∖≼∐↴∖↴⊓⋅↕↴⋝∏⊓∪∐∶↴∙⊾↕↖⇁↸∖∐↴⋝∙↖⇁ ⋅⋅⋅ ⋅ Xie within gHthe Τι and Ἡ are the lower and upper limits of the temperature FT: 3 aud > are free parameters that deteriune the shape of the distribution.," Moreover, the IRAS FIR measurements are not adequate to detect the emission coming from cold dust $10-20$ K) which peaks at wavelengths between 200 and 300 $\mu$ m. I adopt the dust temperature distribution given by Kwan Xie (1992): $T_L$ and $T_U$ are the lower and upper limits of the temperature $T$; $\beta$ and $\gamma$ are free parameters that determine the shape of the distribution."
+ The equations relating the temperature distribution q(T) to the luminosity emitted by the dust aud to the dust mass are detailed in Kiwan Xie (1992)., The equations relating the temperature distribution $g(T)$ to the luminosity emitted by the dust and to the dust mass are detailed in Kwan Xie (1992).
+ Due to the observed spectral range. I adopt T;j27 IK aud Τζος Ik. taking into account ouly temperature distributions peaking at a value intermediate between Tp aud Tr aud excluding those pairs of . aud > which produce unrealistic distributious (as. for instance. nonotomically increasing or decreasing functions).," Due to the observed spectral range, I adopt $T_L$ =7 K and $T_U$ =60 K, taking into account only temperature distributions peaking at a value intermediate between $T_L$ and $T_U$ and excluding those pairs of $\beta$ and $\gamma$ which produce unrealistic distributions (as, for instance, monotonically increasing or decreasing functions)."
+ The main problem iu the preseut approach is to select the proper values for the parameters ο) aud >. the choice being constrained by the ratio of the flix deusities at two cifferent waveleneths.," The main problem in the present approach is to select the proper values for the parameters $\beta$ and $\gamma$, the choice being constrained by the ratio of the flux densities at two different wavelengths."
+ I first ideutifv a rauge of J aud > pairs which produce the observed flux ratio., I first identify a range of $\beta$ and $\gamma$ pairs which produce the observed flux ratio.
+ Since the same flux ratio can be obtained by functions having quite different shapes. a further coustraiut is needed.," Since the same flux ratio can be obtained by functions having quite different shapes, a further constraint is needed."
+ Unfortunately. for noue of he galaxies iu the sample submillimeter observations are," Unfortunately, for none of the galaxies in the sample submillimeter observations are"
+an extended tail is detected towards the northeast of the galactic disk.,an extended tail is detected towards the northeast of the galactic disk.
+ This N-rav tail is associated with the Io plume reported by Yoshida et al. (2002)), This X-ray tail is associated with the $\alpha$ plume reported by Yoshida et al. \cite{yoshida2002}) )
+ aud the beginning of the huge tail discovered by Oosterloo van Corkom (2005) (Fie. 3))., and the beginning of the huge tail discovered by Oosterloo van Gorkom (2005) (Fig. \ref{fig:n4388x}) ).
+ At distances 25’ from the ealactic disk the N-rav emission starts to be dominated x the bright halo of M 86., At distances $> 5'$ from the galactic disk the X-ray emission starts to be dominated by the bright halo of M 86.
+ We performed a spectral analysis of the nucleus. he disk aud nucleus. aud the part of the tail cOse TO he ealactic disk.," We performed a spectral analysis of the nucleus, the disk and nucleus, and the part of the tail close to the galactic disk."
+ The spectra (Fig. 1)), The spectra (Fig. \ref{fig:n4388spec}) )
+ are fitted with wo sinele temperature thermal plasina MIENAL models with an acditional powerluw compoucut to account for uidetected point sources., are fitted with two single temperature thermal plasma MEKAL models with an additional powerlaw component to account for undetected point sources.
+ Iu the fit of the disk axd t1C wicleus. a contribution from the ACN was iicluded: by lding to the model an absorbed powerlaw coniponueut id the 6.1 keV iron line (see Table 5)).," In the fit of the disk and the nucleus, a contribution from the AGN was included by adding to the model an absorbed powerlaw component and the 6.4 keV iron line (see Table \ref{tab:models}) )."
+ The te11]yeratures xd deusitics are eiven in Table 3.., The temperatures and densities are given in Table \ref{tab:tempdens}.
+ The nucleis has a gas deusity of ~301)°° 2.," The nucleus has a gas density of $\sim 30\,\eta^{-0.5}$ $^{-3}$."
+" The cold aud lot components the disk have temperatures of ~0.2 keV aud 0.5 keV. respectively,"," The cold and hot components of the disk have temperatures of $\sim 0.2$ keV and $\sim 0.8$ keV, respectively."
+ The exceptionally high temperature of the hot compoucut in the disk is probably due to the nuclear outflow (Veilleux et al., The exceptionally high temperature of the hot component in the disk is probably due to the nuclear outflow (Veilleux et al.
+ 1999. Yoshida et al.," 1999, Yoshida et al."
+ 2001)., 2004).
+ In the tail the temperature of the cold component is comparable, In the tail the temperature of the cold component is comparable
+Recent gas infall should also influence abundances of voung population in the LAIC for (he elements other than N. Here we discuss (his issue by comparing the LMC abundance with those of HIIVC's or the SAIC.,Recent gas infall should also influence abundances of young population in the LMC for the elements other than N. Here we discuss this issue by comparing the LMC abundance with those of HVCs or the SMC.
+ Supereiants in both the LMC and the SAIC! basically exhibit the solar abundance ratios for a-elements and iron-peak elements (lll 2004)., Supergiants in both the LMC and the SMC basically exhibit the solar abundance ratios for $\alpha$ -elements and iron-peak elements (Hill 2004).
+ In addition. both AMC's exhibit a similar overabundance for neutron-capture elements (11Η 2004).," In addition, both MCs exhibit a similar overabundance for neutron-capture elements (Hill 2004)."
+ In (his way. two galaxies have a very similar present-day elemental abundance pattern.," In this way, two galaxies have a very similar present-day elemental abundance pattern."
+ It implies that eas intall [rom (he SAIC would dilute the N abundance with little imprint in other elemental ratios., It implies that gas infall from the SMC would dilute the N abundance with little imprint in other elemental ratios.
+ On the other hand. the information on abundances of HIIVCs for the elements other iui N and O is very restricted.," On the other hand, the information on abundances of HVCs for the elements other than N and O is very restricted."
+ Complex C seems (o exhibit an essentially solar pattern for a-elements (O. 5. Si) in comparison with Fe but with a possible slightlv-enhanced [a- (Collins et al.," Complex C seems to exhibit an essentially solar pattern for $\alpha$ -elements (O, S, Si) in comparison with Fe but with a possible slightly-enhanced $\alpha$ -elements/Fe] (Collins et al."
+ 2007)., 2007).
+ If the future observation reveals the clear SN I-like [a- for IIVCs. we will be able to conclude Chat the dilution by them will change ihe LMC abundance into the one at odds with the observed LAIC abundance.," If the future observation reveals the clear SN II-like $\alpha$ -elements/Fe] for HVCs, we will be able to conclude that the dilution by them will change the LMC abundance into the one at odds with the observed LMC abundance."
+ IIowever. at 1 moment. it ean be said that large impact on abundances by infall of IIVCS is likely (o be seen only in the deficiency of the N abundance.," However, at the moment, it can be said that large impact on abundances by infall of HVCs is likely to be seen only in the deficiency of the N abundance."
+ The proposed accretion scenario in the present paper mieht well be one of possible scenarios., The proposed accretion scenario in the present paper might well be one of possible scenarios.
+ Whatever alternative scenario is proposed. it would need to discuss the origin of the unusually low [N/II] in the context of an unique LAIC’s environment (e.g.. interaction with the Galaxy aud the SAIC).," Whatever alternative scenario is proposed, it would need to discuss the origin of the unusually low [N/H] in the context of an unique LMC's environment (e.g., interaction with the Galaxy and the SMC)."
+ We here discuss (μου alternative scenarios (hat could possibly explain the observed unusually low [N/H]., We here discuss three alternative scenarios that could possibly explain the observed unusually low [N/H].
+ The first is that the stellar winds of AGB stars (i.e.. rich. sources of nitrogen) can be efficiently and continuously stripped from the LMC owing to hvdrodyvnamical interaction between the LMC and the Galactic warm halo gas so that the present [N/II] can be rather low.," The first is that the stellar winds of AGB stars (i.e., rich sources of nitrogen) can be efficiently and continuously stripped from the LMC owing to hydrodynamical interaction between the LMC and the Galactic warm halo gas so that the present [N/H] can be rather low."
+ This scenario with selective loss of ACB ejecta would have a problem of explaining other observational results on chemical abundances of stars (e.g.. Che observed abundances of s-process elements).," This scenario with selective loss of AGB ejecta would have a problem of explaining other observational results on chemical abundances of stars (e.g., the observed abundances of s-process elements)."
+ llowever. it is possible that stellar ejecta only [rom massive AGB stars (i.e... rich source ol nitrogen) are removed owing to the stronger winds whereas those [rom low-mass ones (i.e.. rich sources of s-process elements) are not.," However, it is possible that stellar ejecta only from massive AGB stars (i.e., rich source of nitrogen) are removed owing to the stronger winds whereas those from low-mass ones (i.e., rich sources of s-process elements) are not."
+ If (his is the case. the above problem associated wilh the s-process elements would not be so serious for the selective loss scenario.," If this is the case, the above problem associated with the s-process elements would not be so serious for the selective loss scenario."
+ We consider that (his scenario is highlv unlikely. because it needs to explain why gas ου supernovae," We consider that this scenario is highly unlikely, because it needs to explain why gas from supernovae"
+Sudden energy releases in the solar atmosphere are known to generate large scale global waves propagating over long distances (see. e.g. Moreton and Ramsey 1960. Uchida 1970. Thompson et al.,"Sudden energy releases in the solar atmosphere are known to generate large scale global waves propagating over long distances (see, e.g. Moreton and Ramsey 1960, Uchida 1970, Thompson et al."
+" 1999, Ballai et al."," 1999, Ballai et al."
+ 2005)., 2005).
+ The energy stored in these waves can be released by traditional dissipative mechanisms. but it could also be transferred to magnetic structures which may come in contact with global waves.," The energy stored in these waves can be released by traditional dissipative mechanisms, but it could also be transferred to magnetic structures which may come in contact with global waves."
+" This scenario is true not only for coronal structures. but applies to all magnetic entities in the solar atmosphere that can serve as waveguides (see. e.g. Wills-Davey and Thompson 1999, Patsourakos and Vourlidas 2009. Liu et al."," This scenario is true not only for coronal structures, but applies to all magnetic entities in the solar atmosphere that can serve as waveguides (see, e.g. Wills-Davey and Thompson 1999, Patsourakos and Vourlidas 2009, Liu et al."
+ 2010. ete).," 2010, etc)."
+ In the corona. EIT waves generated by coronal mass ejections (CMEs) and/or flares could interact with coronal loops. resulting in the generation of kink modes. re. oscillations which exhibit periodic movement about the loop's symmetry axis.," In the corona, EIT waves generated by coronal mass ejections (CMEs) and/or flares could interact with coronal loops, resulting in the generation of kink modes, i.e. oscillations which exhibit periodic movement about the loop's symmetry axis."
+ Global waves can also interact with prominence fibrils as observed by. e.g. Ramsey and Smith (1966). and more recently by Eto et al. (," Global waves can also interact with prominence fibrils as observed by, e.g. Ramsey and Smith (1966), and more recently by Eto et al. ("
+2002). Jing et al. (,"2002), Jing et al. ("
+2003). Okamoto et al. (,"2003), Okamoto et al. ("
+2004). Isobe and Tripathi (2007). Pintérr et al. (,"2004), Isobe and Tripathi (2007), Pintérr et al. ("
+2008).,2008).
+ Oscillations of magnetic structures were (and are currently) used as a basic ingredient in one of solar physics most dynamically expanding fields. namelyseismology. where observations of wave characteristics (amplitude. wavelength. propagation speed. damping time/length) are corroborated with theoretical modelling (dispersion and evolutionary equations. as well as MHD models) in order to derive quantities that cannot be directly or indirectly measured (magnetic field magnitude and sub-resolution structuring. transport coefficients. heating functions. thermal state of the plasma. stratification parameters. etc.).," Oscillations of magnetic structures were (and are currently) used as a basic ingredient in one of solar physics most dynamically expanding fields, namely, where observations of wave characteristics (amplitude, wavelength, propagation speed, damping time/length) are corroborated with theoretical modelling (dispersion and evolutionary equations, as well as MHD models) in order to derive quantities that cannot be directly or indirectly measured (magnetic field magnitude and sub-resolution structuring, transport coefficients, heating functions, thermal state of the plasma, stratification parameters, etc.)."
+ Considerable advances have been achieved in diagnosing the state of the field and plasma (see. e.g-Roberts et al.," Considerable advances have been achieved in diagnosing the state of the field and plasma (see, e.g.Roberts et al."
+ 1984. Nakariakov et al.," 1984, Nakariakov et al."
+" 1999, Nakariakov and Ofman 2001. Ofman and Thompson 2002. Ruderman and Roberts 2002. Andries et al."," 1999, Nakariakov and Ofman 2001, Ofman and Thompson 2002, Ruderman and Roberts 2002, Andries et al."
+ 2005. 2009. Ballai et al.," 2005, 2009, Ballai et al."
+ 2005. 2011. Gruszecki et al.," 2005, 2011, Gruszecki et al."
+ 2006. 2007. 2008. Banerjee et al.," 2006, 2007, 2008, Banerjee et al."
+ 2007. Ofman 2007. 2009. MeLaughlin and Ofman 2008. Verth et al.," 2007, Ofman 2007, 2009, McLaughlin and Ofman 2008, Verth et al."
+ 2007. Ballai 2007. Ruderman et al.," 2007, Ballai 2007, Ruderman et al."
+ 2008. Van Doorsselaere et al.," 2008, Van Doorsselaere et al."
+ 2008. Verth and Erdéllyi 2008. Verth et al.," 2008, Verth and Erdéllyi 2008, Verth et al."
+ 2008. Morton and Erdéllyi 2009. Ruderman and Erdéllyti 2009. Andries et al.," 2008, Morton and Erdéllyi 2009, Ruderman and Erdéllyi 2009, Andries et al."
+ 2009. Selwa and Ofman 2009. Selwa et al.," 2009, Selwa and Ofman 2009, Selwa et al."
+ 2010)., 2010).
+ It is highly likely that higher resolution observations made possible recently by space satellites such as STEREO. Hinode. SDO (and future missions) will further help understanding the complicated reality of the solar plasma environment.," It is highly likely that higher resolution observations made possible recently by space satellites such as STEREO, Hinode, SDO (and future missions) will further help understanding the complicated reality of the solar plasma environment."
+ Indeed since their launch. data provided by these satellites are already shedding light on numerous aspects of coronal seismology. e.g. Verwichte et al. (," Indeed since their launch, data provided by these satellites are already shedding light on numerous aspects of coronal seismology, e.g. Verwichte et al. ("
+2009) used STEREO data to determine the three-dimensional geometry of the loop. SDO/AIA data was used by Aschwanden and Schrijver (2011) to prove the coupling of the kink mode and eross-sectional oscillations that could be explained as a consequence of the loop length variation in the vertical polarization mode.,"2009) used STEREO data to determine the three-dimensional geometry of the loop, SDO/AIA data was used by Aschwanden and Schrijver (2011) to prove the coupling of the kink mode and cross-sectional oscillations that could be explained as a consequence of the loop length variation in the vertical polarization mode."
+ Finally. based on Hinode data. Ofman and Wang (2008) provided the first evidence for transverse waves in coronal multithreaded loops with cool plasma ejected from the chromosphere flowing along the threads.," Finally, based on Hinode data, Ofman and Wang (2008) provided the first evidence for transverse waves in coronal multithreaded loops with cool plasma ejected from the chromosphere flowing along the threads."
+ On the other hand the development of even the fundamental mode is not always guaranteed. as was shown observationally by. e.g. Aschnwanden et al. (," On the other hand the development of even the fundamental mode is not always guaranteed, as was shown observationally by, e.g. Aschnwanden et al. ("
+2002) and later using MHD modelling by Selwa and Ofman (2010) and Selwa et al. (,2002) and later using MHD modelling by Selwa and Ofman (2010) and Selwa et al. (
+2011 a.b).,"2011 a,b)."
+ The dispersion relations for many simple (and some quite complicated) plasma waves under the assumptions of ideal magnetohydrodynamics (MHD) are well known: they were derived long before accurate EUV observations were available (see. e.g. Edwir and Roberts 1983. Roberts et al.," The dispersion relations for many simple (and some quite complicated) plasma waves under the assumptions of ideal magnetohydrodynamics (MHD) are well known; they were derived long before accurate EUV observations were available (see, e.g. Edwin and Roberts 1983, Roberts et al."
+ 1984) using simplified models within the framework of ideal and linear MHD., 1984) using simplified models within the framework of ideal and linear MHD.
+ Although the realistic interpretation of many observations 1s nade difficult especially by the poor spatial resolution of present satellites not being quite sufficient. considerable amount of information about the thermodynamical and dynamical state of the plasma. and the structure and magnitude of the coronal magnetic field. can still be obtained.," Although the realistic interpretation of many observations is made difficult especially by the poor spatial resolution of present satellites not being quite sufficient, considerable amount of information about the thermodynamical and dynamical state of the plasma, and the structure and magnitude of the coronal magnetic field, can still be obtained."
+ The mathematical description of waves and oscillations in solar structures is. in general. given by equations whose coefficients vary in space and time.," The mathematical description of waves and oscillations in solar structures is, in general, given by equations whose coefficients vary in space and time."
+ It has been recognised recently by. e.g. Andries et al. (," It has been recognised recently by, e.g. Andries et al. ("
+2005) that the longitudinal stratification (1.e. along the longitudinal symmetry axis of the tube that coincides with the direction of the magnetic,2005) that the longitudinal stratification (i.e. along the longitudinal symmetry axis of the tube that coincides with the direction of the magnetic
+The final total intensity images are shown in Fig. l..,The final total intensity images are shown in Fig. \ref{fig:figure1}.
+ A codename (band+array) is given to each image cf., A codename (band+array) is given to each image cf.
+ Table 3.., Table \ref{tab:table3}.
+ The high resolution images (LB and CC) show details of the jets. while in the relatively low resolution images (LC and CD) the diffuse lobe structures are better seen.," The high resolution images (LB and CC) show details of the jets, while in the relatively low resolution images (LC and CD) the diffuse lobe structures are better seen."
+ 334.47 has a remarkably bright core. unlike typical double-lobed 3CR quasars whose cores generally contribute <10% of the total 6 or 20 em flux density (e.g..2)..," 34.47 has a remarkably bright core, unlike typical double-lobed 3CR quasars whose cores generally contribute $\la 10\%$ of the total 6 or 20 cm flux density \citep[e.g.,][]{Bridle1994}."
+ Table 4. summarizes the resulting imaging figures., Table \ref{tab:table4} summarizes the resulting imaging figures.
+ While the theoretical noise was not reached. acceptable figures were obtained.," While the theoretical noise was not reached, acceptable figures were obtained."
+ The resulting angular resolution is ~4” and ~ respectively. cf.," The resulting angular resolution is $\sim$ and $\sim$ respectively, cf."
+ expectation., expectation.
+" Given that the overall source angular size is Just over4""... primaryy beam corrections are not necessaryy (the VLA 6cm primary beam width R=1.22-,U/Dx 98)."," Given that the overall source angular size is just over, primary beam corrections are not necessary (the VLA 6cm primary beam width $R \simeq 1.22 \cdot \lambda / D \simeq 9{\farcm}8$ )."
+ The images show that both hot spots andthe core are well aligned. to better than 1°.," The images show that both hot spots andthe core are well aligned, to better than $^{\circ}$."
+ The high resolution 6em image (CC) displays a beautiful straight one-sided jet. which stretches almost continously from the core to the southern hot spot — a (projected) distance of kkpe.," The high resolution 6cm image (CC) displays a beautiful straight one-sided jet, which stretches almost continously from the core to the southern hot spot – a (projected) distance of kpc."
+" For comparison. the straight part of the 3315 jet measures 310kkpe (e.g..22).. ""m"," For comparison, the straight part of the 315 jet measures kpc \citep[e.g.,][]{Bridle1979, Mack1997}."
+"""To our best knowledge. the Jet in 334.47 is the record straight jet."," To our best knowledge, the jet in 34.47 is the record longest straight jet."
+ Even at relativistic speeds. the particle travel time from core to southern hot spot well exceeds one million years.," Even at relativistic speeds, the particle travel time from core to southern hot spot well exceeds one million years."
+ This southern hot spot is more compact than the northern one: it is also more distant from the core than the northern one ts., This southern hot spot is more compact than the northern one; it is also more distant from the core than the northern one is.
+ No counterjet (towards the northern lobe) is seen., No counterjet (towards the northern lobe) is seen.
+ As alluded to already. the intensity of the core is high: its relative contribution is ~35% of the 20cm and - of the 6cm total intensity.," As alluded to already, the intensity of the core is high: its relative contribution is $\sim$ of the 20cm and $\sim$ of the 6cm total intensity."
+ Its spectral index «ae-0.3 (S.x ," Its spectral index $\alpha^{5GHz}_{1.4GHz}
+= -0.3$ $_{\nu} \propto \nu^{\alpha}$ )."
+From the low resolution data we infer à total GGHz flux density of 1.4JJy: taking out the probably relativistically boosted core emission then yields a GGHz integrated radio source luminosity of 1.1x107 W/Hz., From the low resolution data we infer a total GHz flux density of Jy; taking out the probably relativistically boosted core emission then yields a GHz integrated radio source luminosity of $1.1 \times 10^{26}$ W/Hz.
+ With reference to ?.. we note the resemblance of 334.47 to FR2 quasar 447. concerning their overall radio morphologies and the knot strueture in their one-sided jets (but we keep in mind that the latter is roughly a factor twenty more luminous than the former).," With reference to \citet{Fernini1991}, we note the resemblance of 34.47 to FR2 quasar 47, concerning their overall radio morphologies and the knot structure in their one-sided jets (but we keep in mind that the latter is roughly a factor twenty more radio-luminous than the former)."
+ Images at the same frequency but with different resolutions can be combined to obtain a multi-resolution image., Images at the same frequency but with different resolutions can be combined to obtain a multi-resolution image.
+ These then show both the detailed features of the high resolution and the low surface brightness features of the low resolution image., These then show both the detailed features of the high resolution and the low surface brightness features of the low resolution image.
+ Employing core subtraction as deseribed above. we obtained such images at 20em and at 6em: these are shown in Fig.2 (LB+LC) and Fig.3 (CC+CD). respectively.," Employing core subtraction as described above, we obtained such images at 20cm and at 6cm; these are shown in \ref{fig:figure2} (LB+LC) and \ref{fig:figure3}
+ (CC+CD), respectively."
+ A fixed circular clean beam of 775 was adopted. in order to facilitate intercomparison of the images. incl.," A fixed circular clean beam of 5 was adopted, in order to facilitate intercomparison of the images, incl."
+ radio spectral behaviour., radio spectral behaviour.
+ The images reach noise levels of ~0.12 and ~0.06 mJy/beam respectively., The images reach noise levels of $\sim 0.12$ and $\sim 0.06$ mJy/beam respectively.
+ The straight jet stands out nicely in the 6cm image., The straight jet stands out nicely in the 6cm image.
+ It displays a number of knots and connects to the more distant of the two hot spots. eq.," It displays a number of knots and connects to the more distant of the two hot spots, cq."
+ radio lobes., radio lobes.
+ No counter-jet is observed., No counter-jet is observed.
+ The variability of the radio core is noteworthy., The variability of the radio core is noteworthy.
+ We compare the measured 1984 core strength with earlier measurements made with the Westerbork Synthesis Radio Telescope (WSRT). in 1973/74 (?) and in 1974/79 (?)..," We compare the measured 1984 core strength with earlier measurements made with the Westerbork Synthesis Radio Telescope (WSRT), in 1973/74 \citep{Conway1977} and in 1974/79 \citep{Jagers1982}."
+ The results are summarized in Table 5.., The results are summarized in Table \ref{tab:table5}.
+ The data indicate a — decrease in ten years at GGHz. and ~25% at GGHz.," The data indicate a $\sim$ decrease in ten years at GHz, and $\sim$ at GHz."
+ It should be noted that the associated optical QSO 17214343 Vas also found to be variable (?).., It should be noted that the associated optical QSO 1721+343 was also found to be variable \citep{McGimsey1978}. .
+ Relativistic beaming effects. invoked to explain the measured superluminal motion (??).. provide a natural explanation for the behaviour of the thermal) core emission.," Relativistic beaming effects, invoked to explain the measured superluminal motion \citep{Barthel1989a, Hooi1992}, provide a natural explanation for the behaviour of the (non-thermal) core emission."
+ Fig., Fig.
+ + presents the spectral index map («n ) for 334.47. taken from its high resolution images. with capital letters marking positions of interest.," \ref{fig:figure4} presents the spectral index map $\alpha^{5GHz}_{1.4GHz}$ ) for 34.47, taken from its high resolution images, with capital letters marking positions of interest."
+ Both hot spots (positions A and D) display steep emission. with values —0.86 and —0.88. respectively.," Both hot spots (positions A and I) display steep emission, with values $-0.86$ and $-0.88$, respectively."
+ Lobe emission ts also steep. although not extreme: we measure indices —0.77 and —0.84 at locations B and H. As alluded to before. the core (C) has a fairly flat index —0.29; jet knots display values between —0.42 to —0.79 (Dto G).," Lobe emission is also steep, although not extreme: we measure indices $-0.77$ and $-0.84$ at locations B and H. As alluded to before, the core (C) has a fairly flat index $-0.29$ ; jet knots display values between $-0.42$ to $-0.79$ (Dto G)."
+ The spectral index error is estimated to be 20.05 at most., The spectral index error is estimated to be $\pm0.05$ at most.
+ A summary of the measurementsis given in the Table above the spectral index map., A summary of the measurements is given in the Table above the spectral index map.
+"of these equations, the agreement with the numerical results is surprisingly good.","of these equations, the agreement with the numerical results is surprisingly good."
+" After specifying the basic ingredients of our second-order dynamical model, we attempt to apply it to y-Cephei."," After specifying the basic ingredients of our second-order dynamical model, we attempt to apply it to $\gamma$ -Cephei."
+" As mentioned in the introduction, this is probably the best-studied tight binary system with a known planetary body."," As mentioned in the introduction, this is probably the best-studied tight binary system with a known planetary body."
+" Since the main secular parameters g and e; strongly depend on the stellar masses and orbital elements of the secondary star, we begin our discussion by reviewing the accuracy of these parameters."," Since the main secular parameters $g$ and $e_f$ strongly depend on the stellar masses and orbital elements of the secondary star, we begin our discussion by reviewing the accuracy of these parameters."
+" Throughout this work we refer to the more massive stellar component by y-Cephei-A, while y-Cephei-B is used to identify the less massive star."," Throughout this work we refer to the more massive stellar component by $\gamma$ -Cephei-A, while $\gamma$ -Cephei-B is used to identify the less massive star."
+" The giant planet orbiting y- Cephei-A is called y-Cephei-b. The masses of each body are denoted by ma, mg, and mp, in that order."," The giant planet orbiting $\gamma$ -Cephei-A is called $\gamma$ -Cephei-b. The masses of each body are denoted by $m_A$, $m_B$, and $m_p$, in that order."
+" Several years before the discovery of the first planetary body around a main sequence star (Mayor and Queloz 1995), Campbell et al. ("," Several years before the discovery of the first planetary body around a main sequence star (Mayor and Queloz 1995), Campbell et al. ("
+"1988) suggested the presence of a mass object in a 2.7 yr orbit around y-Cephei-A. The authors, however, remained cautious about claiming a true planetary detection, since the observed periodic variations in radial velocity (RV) were at the very limit of the instrumental resolution.","1988) suggested the presence of a Jupiter-mass object in a $2.7$ yr orbit around $\gamma$ -Cephei-A. The authors, however, remained cautious about claiming a true planetary detection, since the observed periodic variations in radial velocity (RV) were at the very limit of the instrumental resolution."
+" To complicate the problemeven further, the variations in RV attributed to the Jovian planet, with a semi-amplitude of only about 25 m/s, were superimposed on a much larger variation caused by a previously unnoticed stellar companion with a much longer orbital period."," To complicate the problemeven further, the variations in RV attributed to the Jovian planet, with a semi-amplitude of only about $25$ m/s, were superimposed on a much larger variation caused by a previously unnoticed stellar companion with a much longer orbital period."
+ The planetary interpretation was questioned later when changes in the chromospheric activity were observed with similar period (Walker et al., The planetary interpretation was questioned later when changes in the chromospheric activity were observed with similar period (Walker et al.
+ 1992)., 1992).
+" Thus, it was proposed that the observed changes in RV were spurious and probably only due to changes in the spectral line profiles caused by surface inhomogeneities (spots)."," Thus, it was proposed that the observed changes in RV were spurious and probably only due to changes in the spectral line profiles caused by surface inhomogeneities (spots)."
+" The existence of a binary component (i.e. y-Cephei-B) was only reevaluated several years later, when Griffin et al. ("," The existence of a binary component (i.e. $\gamma$ -Cephei-B) was only reevaluated several years later, when Griffin et al. ("
+2002) combined several historical sources of radial velocities that include epochs from 1896 to 1980.,2002) combined several historical sources of radial velocities that include epochs from 1896 to 1980.
+" This data set consisted of 88 RV observations, although many of them did not contain proper uncertainties, and a gap of some 50 years was present in the data set."," This data set consisted of 88 RV observations, although many of them did not contain proper uncertainties, and a gap of some 50 years was present in the data set."
+" Even so, the authors proposed a secondary stellar mass in the system with an orbital period of P~66 yrs."," Even so, the authors proposed a secondary stellar mass in the system with an orbital period of $P \sim 66$ yrs."
+" The presence of a third body, this time a planet around y- was only confirmed by Hatzes et al. ("," The presence of a third body, this time a planet around $\gamma$ -Cephei-A, was only confirmed by Hatzes et al. ("
+2003) after incorporating new high-precision velocity observations from the McDonald Observatory.,2003) after incorporating new high-precision velocity observations from the McDonald Observatory.
+" They show convincingly that the 2.5 yr variation in radial velocity was coherent in phase and amplitude throughout the entire 20 yr interval, as would be expected for Keplerian motion, and that no changes were Observed in the spectral-line bisectors."," They show convincingly that the $2.5$ yr variation in radial velocity was coherent in phase and amplitude throughout the entire $20$ yr interval, as would be expected for Keplerian motion, and that no changes were observed in the spectral-line bisectors."
+" More recently Torres (2007) has again analyzed the historical sources of radial velocities using the extensive Harvard-Smithsonian Center for Astrophysics (CfA) database consisting of ~250,000 spectra."," More recently, Torres (2007) has again analyzed the historical sources of radial velocities using the extensive Harvard-Smithsonian Center for Astrophysics (CfA) database consisting of $\sim 250,000$ spectra."
+ Torres pointed out that some of the historical radial velocities showed large internal discrepancies when compared with other data taken at similar times and were consequently not reliable., Torres pointed out that some of the historical radial velocities showed large internal discrepancies when compared with other data taken at similar times and were consequently not reliable.
+ The author constructed a reliable data set consisting of 30 RV observations., The author constructed a reliable data set consisting of 30 RV observations.
+ The complete sets of radial velocities (four sets by Hatzes et al. (, The complete sets of radial velocities (four sets by Hatzes et al. (
+"2003) and one by Torres (2007)) are shown in Figure 4,, where the errors bars indicate the uncertainty on each numerical value.","2003) and one by Torres (2007)) are shown in Figure \ref{fig4}, where the errors bars indicate the uncertainty on each numerical value."
+" The difference in precision is remarkable, showing how the incorporation of modern techniques in RV measurements lead to the detection of the planetary mass."," The difference in precision is remarkable, showing how the incorporation of modern techniques in RV measurements lead to the detection of the planetary mass."
+ This increase in precision also allowed the mean anomaly M and longitude of pericenter c of the binary component to be accurately defined., This increase in precision also allowed the mean anomaly $M$ and longitude of pericenter $\varpi$ of the binary component to be accurately defined.
+" Independently and without attempting to identify any planetary body, Gontcharov et al. ("," Independently and without attempting to identify any planetary body, Gontcharov et al. ("
+2000) studied the mass and orbital parameters of the binary system using astrometric observations from several sources.,2000) studied the mass and orbital parameters of the binary system using astrometric observations from several sources.
+ They obtained an orbital period of ~45 yr and a total mass of 3Mo., They obtained an orbital period of $\sim 45$ yr and a total mass of $3 M_{\odot}$ .
+" Unfortunately, the individual masses were not specified."," Unfortunately, the individual masses were not specified."
+" In his work, Torres (2007) also combined a total of 140 astrometric measurements"," In his work, Torres (2007) also combined a total of $140$ astrometric measurements"
+" (AT <100 10°°—1075 (107 46°: 2.02-0.] 9 ~37 30” (45” 60” X7 ~28 ~40 Dec.2—37738/55"" ~4’ 5"" (7.3) X ", $kT$ $<100$ $10^{36}-10^{38}$ $10^{38-40}$ $^{\circ}$ $2.0\pm0.1$ $\sim 9$ $\sim 37$ $30''$ $45''$ $60''$ $\chi^2$ $\sim 28$ $\sim 40$ $-37^{\circ}38'55''$ $\sim 4'$ $5''$ $7.3''$ $\chi^2$ 
+Dust is the cornerstone of the active galactic nuclei (AGN) Unified Scheme that postulates that the diversity of the observed. properties of AGN are merely a result. of the different lines of sight with respect to obscuring material surrounding the active nucleus (Antonucci 1993: Urri&Padovani 1995: Tacdhunter 2008)).,Dust is the cornerstone of the active galactic nuclei (AGN) Unified Scheme that postulates that the diversity of the observed properties of AGN are merely a result of the different lines of sight with respect to obscuring material surrounding the active nucleus \citealt{antonucci93}; ; \citealt{urri95}; \citealt{tadhunter08}) ).
+ The elements necessary for understanding the nature of AGN and the variety of their properties are in fact the geometry of the cireumnuclear dust. ane the amount of obscuration., The elements necessary for understanding the nature of AGN and the variety of their properties are in fact the geometry of the circumnuclear dust and the amount of obscuration.
+ The latter is also needed. in order to understand the physical conditions of the dust cnshroudec nuclear region and to estimate the intrinsic UV-to-optical properties of the AGN., The latter is also needed in order to understand the physical conditions of the dust enshrouded nuclear region and to estimate the intrinsic UV-to-optical properties of the AGN.
+ Obscuring clusty material should. re-radiate in the infrared (11) the fraction of aceretion Luminosity it absorbs. providing thus a direct. metric for the study of the medium.," Obscuring dusty material should re-radiate in the infrared (IR) the fraction of accretion luminosity it absorbs, providing thus a direct metric for the study of the medium."
+ Indeed. all spectroscopically confirmed ACN ving in regions covered by the various Spitzer Space Telescope (Spitzer) surveys show significant Hi emission down to the detectionLimits., Indeed all spectroscopically confirmed AGN lying in regions covered by the various Spitzer Space Telescope (Spitzer) surveys show significant IR emission down to the detectionlimits.
+and dup=50mas+15mas.,"and $d_{\rm UD} = 50\,{\rm mas}\pm 
+15\,{\rm mas}$."
+" These diameters correspond to pc + pc and pc + pe, respectively."," These diameters correspond to pc $\pm$ pc and pc $\pm$ pc, respectively."
+" An alternative model interpretation of the decreasing visibility function is, for example, an object slightly larger than the above diameters plus an unresolved central object or an object which is unresolved in one direction (~ EW), but resolved in the direction perpendicular to it (~ NS)."," An alternative model interpretation of the decreasing visibility function is, for example, an object slightly larger than the above diameters plus an unresolved central object or an object which is unresolved in one direction $\sim$ EW), but resolved in the direction perpendicular to it $\sim$ NS)."
+ Figure 3 shows the diffraction-limited speckle masking reconstruction of NGC11068., Figure \ref{rek} shows the diffraction-limited speckle masking reconstruction of 1068.
+ The resolution of the reconstructed image is 76mmas., The resolution of the reconstructed image is mas.
+" It shows an elongated structure in the northern direction, ie. approximately the position of the radio jet."," It shows an elongated structure in the northern direction, i.e. approximately the position of the radio jet."
+ Figure 4 shows the azimuthally averaged radial plots of the reconstructions., Figure \ref{rad} shows the azimuthally averaged radial plots of the reconstructions.
+ Photometry was performed by comparing the integral intensities of the long-exposure images of 11068 and the photometric standard star 1110609., Photometry was performed by comparing the integral intensities of the long-exposure images of 1068 and the photometric standard star 110609.
+ It yields for the flux Εν of 11068 in the K-band the value Fx=650mJy200mJy.," It yields for the flux $F_{\nu}$ of 1068 in the K-band the value $F_{\rm K} = 650\,{\rm mJy} \pm 200\,{\rm mJy}$ ."
+" ;From this value we have to subtract the flux of the underlying faint extended component in order to get the flux ΕΠΟ""*5 of only the mmas component."," >From this value we have to subtract the flux of the underlying faint extended component in order to get the flux $F_{\rm K}^{\rm 30\,mas}$ of only the mas component."
+ We determine the contribution of this extended component from the zero-frequency visibility peak discussed above to + of Fk., We determine the contribution of this extended component from the zero-frequency visibility peak discussed above to $\pm$ of $F_{\rm K}$.
+ This results in a flux from the mmas component of F20™s=520mJy+210mJy.," This results in a flux from the mas component of $F_{\rm K}^{\rm 30\,mas} = 
+520\,{\rm mJy} \pm 210 \,{\rm mJy}$."
+" MPH96 showed that the central ~2"" have substructures on a scale of mmas requiring observations with a resolution of better than 100mas to separate the true nuclear spectrum from that of surrounding sources."," MPH96 showed that the central $\sim 2^{\prime\prime}$ have substructures on a scale of mas requiring observations with a resolution of better than $100 \,{\rm mas}$ to separate the true nuclear spectrum from that of surrounding sources."
+ Only a few published radio flux determinations of the nucleus of 11068 have a sufficiently high angular resolution to allow the separation of individual nuclear components and thus to use it for physical investigations of the nucleus' property., Only a few published radio flux determinations of the nucleus of 1068 have a sufficiently high angular resolution to allow the separation of individual nuclear components and thus to use it for physical investigations of the nucleus' property.
+" Fortunately, our speckle observations have the required high angular resolution (76 mas)."," Fortunately, our speckle observations have the required high angular resolution (76 mas)."
+" Therefore, our resolution would allow the separation of the individual core components discussed by MPH96 if present in the IR."," Therefore, our resolution would allow the separation of the individual core components discussed by MPH96 if present in the IR."
+ We assume that the single source that we have Observed in the K-band is the same as the true nucleus observed by MPH96., We assume that the single source that we have observed in the K-band is the same as the true nucleus observed by MPH96.
+ Our observations constitute an upper limit to the volume from which the above determined 520mJy+210 originate.," Our observations constitute an upper limit to the volume from which the above determined $520\,{\rm mJy} \pm 210 \,{\rm mJy}$ originate."
+" Usually, the K band flux from nuclei of 22 galaxies is attributed either to a warm dust torus or a compact nuclear stellar cluster or a combination of the two (e.g., Thatte et 11997)."," Usually, the K band flux from nuclei of 2 galaxies is attributed either to a warm dust torus or a compact nuclear stellar cluster or a combination of the two (e.g., Thatte et 1997)."
+" If our source is the torus that is held responsible for the different appearences of Seyfert11 and 2 galaxies, our observation constitute the first determination of a torus size."," If our source is the torus that is held responsible for the different appearences of 1 and 2 galaxies, our observation constitute the first determination of a torus size."
+" To clarify the nature of the radiating source, further spectroscopic and polarimetric measurements with a similarly high angular resolution are necessary."," To clarify the nature of the radiating source, further spectroscopic and polarimetric measurements with a similarly high angular resolution are necessary."
+" However, a combination of the flux measurements and nuclear source identification by MPH96 in the radio frequency regime with our observation makes it intriguing to speculate whether a sizable fraction of F;3°™*S could originate from the very nucleus of NGC 1068 rather than from the torus."," However, a combination of the flux measurements and nuclear source identification by MPH96 in the radio frequency regime with our observation makes it intriguing to speculate whether a sizable fraction of $F_{\rm K}^{\rm 30\,mas}$ could originate from the very nucleus of NGC 1068 rather than from the torus."
+ This could be achieved in a scattering halo above and below the nuclear torus., This could be achieved in a scattering halo above and below the nuclear torus.
+ In this halo a large part of the flux could be isotropically scattered rather than absorbed and thermalized in an opaque torus along our direct line-of-sight to the nucleus., In this halo a large part of the flux could be isotropically scattered rather than absorbed and thermalized in an opaque torus along our direct line-of-sight to the nucleus.
+ Fg*5 lies only about a factor of two above the extrapolated," $F_{\rm K}^{\rm 30\,mas}$ lies only about a factor of two above the extrapolated"
+below the radiation reaction limit.,below the radiation reaction limit.
+ To determine (he evolution of the orbits kev parameters. first consider one hallevele segment of the orbit. 2€[0.zi]. where σι is the next midplane-crossing.," To determine the evolution of the orbit's key parameters, first consider one half-cycle segment of the orbit, $z\in[0,z_1]$, where $z_1$ is the next midplane-crossing."
+ The motion is determined bv (wo parameters: 5o and (he midplane-crossing angle 6j., The motion is determined by two parameters: $\gamma_0$ and the midplane-crossing angle $\theta_0$.
+" We wish to calculate the slight changes in > and |@| over this segment. i.e.. the values >(2,) and |8(z4)|: from this we will derive the secular evolution of Jua."," We wish to calculate the slight changes in $\gamma$ and $|\theta|$ over this segment, i.e., the values $\gamma(z_1)$ and $|\theta(z_1)|$; from this we will derive the secular evolution of $\ymax$."
+" The s-lactor follows from enerey conservation: e(z)=egy+εις. ie. 42) =4, where 7=z/p.."," The $\gamma$ -factor follows from energy conservation: $\epsilon(z) = \epsilon_0 + eE_0 z$, i.e., (z) = _0 +, where $\zbar \equiv z/\rho_c$."
+" The particle'strajectory is obtained from the y component of the relativistic equation of motion: d(v,)/dl=αςd(sv)/dzπω εις. integrating which we gel gU), etn, where 3,9=9sinfye@ since 2=v/e1."," The particle'strajectory is obtained from the $y$ component of the relativistic equation of motion: $d(\gamma v_y)/dt = - \omega_c v_z\Rightarrow d(\gamma v_y)/dz = 
+- \omega_c \equiv - eB_0/m_e c$ , integrating which we get ) = _0 - _0 -, where $\beta_{y0} = \beta\, \sin\theta_0 \simeq \sin\theta_0$ since $\beta = v/c \approx 1$."
+" This gives us the trajectorys apex 2, (where ey—0 and y— Yast! Z5=5o5nfs.", This gives us the trajectory's apex $\zbar_a$ (where $v_y=0$ and $y=\ymax$ ): $\zbar_a = \gamma_0\sin\theta_0$.
+" Next. we assume that |0|<0,«Ll. so that απας=34/9.~3,4. where 7=y/p.."," Next, we assume that $|\theta| \leq \theta_0 \ll 1$, so that $d\ybar/d\zbar = \beta_y/\beta_z \simeq \beta_y$, where $\ybar \equiv y/\rho_c$."
+" Then. integrating Equation (2)). we obtain an explicit expression for the particle trajectory: g(2)=eesinAy)In(1+3,POm6/4 τμ)."," Then, integrating Equation \ref{eq-gamma-beta_y}) ), we obtain an explicit expression for the particle trajectory: $\ybar(\zbar) = -\, \zbar/\beta_{\rm rec} + 
+(\gamma_0/\beta_{\rm rec}^2)\, (1+ \beta_{\rm rec} \sin\theta_0)\, \ln (1+ \beta_{\rm rec} \zbar/\gamma_0)$ ."
+" Anticipating the [ractional increase of the particles energv over τι lo be small. 2424/50~Go« 1. we expand g(2) ancl [ind Unas=US)c95002/2 and the segments length. =, (the distance to the next midplane"," Anticipating the fractional increase of the particle's energy over $z_1$ to be small, $\beta_{\rm rec}\zbar_1/\gamma_0 \sim \theta_0 \ll 1$ , we expand$\ybar(\zbar)$ and find $\ybar_{\rm max} \equiv \ybar (\zbar_a) \simeq \gamma_0 \theta_0^2/2$ and the segment's length $\zbar_1$ (the distance to the next midplane"
+jets (Eichler Cranot 2006: Cranot ct al.,jets (Eichler Granot 2006; Granot et al.
+ 2006: Panaitescu 2007: Yamazaki 2009: Nu IInaug 2010). or even as due to dust scattering (Shao Dai 2007: Shao et al.," 2006; Panaitescu 2007; Yamazaki 2009; Xu Huang 2010), or even as due to dust scattering (Shao Dai 2007; Shao et al."
+ 2008)., 2008).
+ According to our L-T-E correlation (Eq. (, According to our L-T-E correlation (Eq. (
+7)). the N-vay luminosity at the cud time of the plateau cau be expressed as a function of the end time aud the isotropic 2-rav energv release as. We believe that this relation can eive useful coustraiut ou the uuderlving plivsics.,"7)), the X-ray luminosity at the end time of the plateau can be expressed as a function of the end time and the isotropic $\gamma$ -ray energy release as, We believe that this relation can give useful constraint on the underlying physics."
+ For the enerex injection model. a natural mechanisui is the dipole radiation from the PAxuniue down of a magnetaro at the ceuter of the fireball.," For the energy injection model, a natural mechanism is the dipole radiation from the spinning down of a magnetar at the center of the fireball."
+ Note that the injected euergv -nay not be Povutius flux. but can be electrou-positrou pairs (Dai 2001).," Note that the injected energy may not be Poynting flux, but can be electron-positron pairs (Dai 2004)."
+ These pairs oeiteract with the fireball material. leacing to the formation of a relativistic wind bubble.," These pairs interact with the fireball material, leading to the formation of a relativistic wind bubble."
+ When the οποιον injection dominates the dynamical evolution of the external shock. the afterglow intensity should naturally be proportional to the cnerey injection power.," When the energy injection dominates the dynamical evolution of the external shock, the afterglow intensity should naturally be proportional to the energy injection power."
+ So. Lx is actually a measure of the energy. injection rate.," So, $L_{\rm X}$ is actually a measure of the energy injection rate."
+ According to Eq. (, According to Eq. (
+9). Lx is roughly inversely proportional to the timescale of the energy injection. Z4.,"9), $L_{\rm X}$ is roughly inversely proportional to the timescale of the energy injection, $T_{\rm a}$."
+ It hints that the euergv reservoir should be roughly a constant., It hints that the energy reservoir should be roughly a constant.
+ This is consistent with the euergv injection model. which usually assiunues that the central engine is a rapidly rotating millisecond magnetar.," This is consistent with the energy injection model, which usually assumes that the central engine is a rapidly rotating millisecond magnetar."
+ In different CRBs. the surface maguetic feld iuteusities of the ceutral magnetars may be quite different. leading to various energy injection Iuninosities and enerev injection timescales.," In different GRBs, the surface magnetic field intensities of the central magnetars may be quite different, leading to various energy injection luminosities and energy injection timescales."
+ But the total cucrev available for energy mjection is relatively constant (about rotational energev of the magnetar)., But the total energy available for energy injection is relatively constant (about rotational energy of the magnetar).
+ It is mainly constrained x the hlnuitiug augular velocity of the magnetar. which again is determined by the equation of state of neutron stars.," It is mainly constrained by the limiting angular velocity of the magnetar, which again is determined by the equation of state of neutron stars."
+ Additionally. according to Dai (2001). in order to oxoduce an obvious plateau in the afterglow lighteurve. the total injected energy mst ve comparable to the original fireball energy. (which may be comparable to E. πο)," Additionally, according to Dai (2004), in order to produce an obvious plateau in the afterglow lightcurve, the total injected energy must be comparable to the original fireball energy (which may be comparable to $E_{\gamma,\rm iso}$ )."
+ This : ⋅⋡⋅ ⋅ Eq.⋅ ⋅ ↥⋅↸∖≺∣∏∐⋅↸∖⋯↸∖∐↑↕↴∖↴⋜↧∶↴∙⋜↧∐⊔⋅∪∏∶↴∙∐⋅↖⇁↸⊳∪∐↴∖↴↕↴∖↴↑↸∖∐↑↖↖⇁↕↑∐↑∐↸∖↕↑↸∖⋯∪↕∑ sso in(," This requirement is again roughly consistent with the item of $E_{\gamma, \rm iso}^{0.88 \pm 0.08}$ in Eq. ("
+9).,9).
+ Based on he above analvsises. we argued that the L-T-E correlation strongly supports the energy injection model of maguetars.," Based on the above analysises, we argued that the L-T-E correlation strongly supports the energy injection model of magnetars."
+ It also indicates that the newly born millisecond magnuetars associated with GRBs provide a good standard caudle in our Universe., It also indicates that the newly born millisecond magnetars associated with GRBs provide a good standard candle in our Universe.
+ Thus the L-T-E correlation may potentially be usec to test the cosmological models., Thus the L-T-E correlation may potentially be used to test the cosmological models.
+ Our suuple contains 17 lone CRBs and 8 intermediate class CRBs., Our sample contains 47 long GRBs and 8 intermediate class GRBs.
+ From Figure 1. we see that both of these two classes are consistent with tle same L-T-E correlation.," From Figure 1, we see that both of these two classes are consistent with the same L-T-E correlation."
+ Towerer. note that they behave very differeutly in frame work of the two-parameter L-T correlation.," Howerer, note that they behave very differently in frame work of the two-parameter L-T correlation."
+ This is another important advantage of our three-parameter correlation., This is another important advantage of our three-parameter correlation.
+ It indicates that magnuetars mav also form in mterinediate class GRBs. aud their Iuuitiug spinuing is just similar to those maenetars born in lone CRBs.," It indicates that magnetars may also form in intermediate class GRBs, and their limiting spinning is just similar to those magnetars born in long GRBs."
+ Α natural problem will be raised as to whether short GRBs with plateau ράσο in the afterglow also obey the, A natural problem will be raised as to whether short GRBs with plateau phase in the afterglow also obey the
+Evidence shows that the evolution of black holes and that of their host galaxies appear to be closely coupled.,Evidence shows that the evolution of black holes and that of their host galaxies appear to be closely coupled.
+ Lt was found that there is strong correlation between the central black hole masses. Adi. and their bulge stellar velocity dispersion. 0.," It was found that there is strong correlation between the central black hole masses, $M_{bh}$, and their bulge stellar velocity dispersion, $\sigma$."
+ ‘Tremaine et al. (, Tremaine et al. (
+2002) investigated: this relationship in a sample of 31 nearby inactive galaxies and. gave a better expression as. λογο are many methods to estimate the central black vole masses (Bian Zhao 2003a and reference therein).,"2002) investigated this relationship in a sample of 31 nearby inactive galaxies and gave a better expression as, There are many methods to estimate the central black hole masses (Bian Zhao 2003a and reference therein)."
+ In hese methods. the reverberation method is thought to be more reliable.," In these methods, the reverberation method is thought to be more reliable."
+ Using the reverberation mapping method. the sizes of broad line regions (DLIts) and then the central black role masses were obtained for 3v ACGNSs (Llo 1998: Wancdel et al.," Using the reverberation mapping method, the sizes of broad line regions (BLRs) and then the central black hole masses were obtained for 37 AGNs (Ho 1998; Wandel et al."
+ 1999: Ixaspi et. al., 1999; Kaspi et al.
+ 2000)., 2000).
+ For some AGNSs with available )ulee velocity dispersion. and. the reverberation mapping mass. Gebhardt et al. (," For some AGNs with available bulge velocity dispersion and the reverberation mapping mass, Gebhardt et al. ("
+2000) ancl Ferrarese et al. (,2000) and Ferrarese et al. (
+"2001) also found that these AXGNSs also follow the Ady,(0 relation ounded in the nearby inactive galaxies.","2001) also found that these AGNs also follow the $M_{\rm
+bh}-\sigma$ relation founded in the nearby inactive galaxies."
+ As we know. it is dillieult to obtain the bulge velocity. dispersion of AGNs.," As we know, it is difficult to obtain the bulge velocity dispersion of AGNs."
+" In order to investigate this relation in a larger sample of AGNs. Nelson (2000) used the width of OL] line emitting from the narrow line region (NLl1is) to indicate the bulge velocity. dispersion. where o=FMHM(OI1])/2.35. and found that these 37 AXGNs with the reverberation mapping masses follow the Mui,6 relation."," In order to investigate this relation in a larger sample of AGNs, Nelson (2000) used the width of [OIII] line emitting from the narrow line region (NLRs) to indicate the bulge velocity dispersion, where $\sigma=FWHM([OIII])/2.35$, and found that these 37 AGNs with the reverberation mapping masses follow the $M_{\rm
+bh}-\sigma$ relation."
+ Wang Lu (2001) investigated this relation in a sample of NLSIs from Veron-Cotty Veron (2001)., Wang Lu (2001) investigated this relation in a sample of NLS1s from Veron-Cetty Veron (2001).
+ They. used the B band. magnitude and the 111 FWLIAT to estimate the black hole masses and he OIL] FWLIAL to indicate the bulge velocity dispersion., They used the B band magnitude and the $\beta$ FWHM to estimate the black hole masses and the [OIII] FWHM to indicate the bulge velocity dispersion.
+ Vhey found that NLSIs also follow the μι6 relation but with more scatter.," They found that NLS1s also follow the $M_{\rm
+bh}-\sigma$ relation but with more scatter."
+ We should notice that NLSIs deviated rom the correlation defined in the nearby inactive galaxies if we think OLLI] ENIM. is not overestimated because of he spectral resolution., We should notice that NLS1s deviated from the correlation defined in the nearby inactive galaxies if we think [OIII] FWHM is not overestimated because of the spectral resolution.
+ Using the Sloan Digital Sky Survey (SDSS). Boroson (2003) investigated the relation between he black hole mass via the Le? ENIM. and the stellar velocity dispersion. via ΟΠΗ ENIM. in a sample of 107 ow-redshift. raciio-quiet AGNs.," Using the Sloan Digital Sky Survey (SDSS), Boroson (2003) investigated the relation between the black hole mass via the $\beta$ FWHM and the stellar velocity dispersion via [OIII] FWHM in a sample of 107 low-redshift radio-quiet AGNs."
+ They found the correlation is consistent with that defined in nearby. galaxies ancl the OLLI] ENCLIM can predict black hole mass to a factor of 5., They found the correlation is consistent with that defined in nearby galaxies and the [OIII] FWHM can predict black hole mass to a factor of 5.
+ Vhere are only a few racüo-Ioud αλ in Boroson (2003)., There are only a few radio-loud AGNs in Boroson (2003).
+ Shields et al. (, Shields et al. (
+"2003) also investigated the Adi,oo relation as a [function of red:shift for an assembled sample of quasars.",2003) also investigated the $M_{\rm bh}-\sigma$ relation as a function of redshift for an assembled sample of quasars.
+ They suggested. that this correlation can be right out. to recdshift of 23., They suggested that this correlation can be right out to redshift of $z\approx 3$.
+ However. Shields et al. (," However, Shields et al. ("
+2003) noticed that the racio-loud ACGNs seem to deviate from this correlation.,2003) noticed that the radio-loud AGNs seem to deviate from this correlation.
+Crida et al.,Crida et al.
+ 2006). namely when the radius of the Hill sphere exceeds the disk semi-thickness: In the simulations presented here. this happens roughly when My23X107.," 2006), namely when the radius of the Hill sphere exceeds the disk semi-thickness: In the simulations presented here, this happens roughly when $m_p \ge 3\times 10^{-4}$."
+ From this time. the planet may be able to open a gap in the disk provided that the gravitational torques overwhelm the viscous forces (Bryden et al.," From this time, the planet may be able to open a gap in the disk provided that the gravitational torques overwhelm the viscous forces (Bryden et al."
+ 1999: Papaloizou et al., 1999; Papaloizou et al.
+" 2006). a condition which can be expressed as: where ©, is the angular velocity of the planet."," 2006), a condition which can be expressed as: where $\Omega_p$ is the angular velocity of the planet."
+" Since the value for the viscosity m the active region is correlated with the size of the dead-zone. the previous equation shows that the gap structure and consequently the accretion rate onto the planet may depend strongly on Hpz for m,=3x1077."," Since the value for the viscosity in the active region is correlated with the size of the dead-zone, the previous equation shows that the gap structure and consequently the accretion rate onto the planet may depend strongly on $\hdz$ for $m_p\ge 3\times 10^{-4}$."
+" Indeed. considering for instance the case mm,=|M;. Eq."," Indeed, considering for instance the case $m_p=1\;M_J$, Eq."
+ 13 predicts that gap opening should occur for a< 0.01., \ref{viscous} predicts that gap opening should occur for $\alpha \le 0.01$ .
+" Therefore. in the cases Where Hpy/H=1.5 and Hpz/H=2.3. we expect a |M, planet to be able to open a gap only in the dead-zone. possibly leaving the gas flowing over the top of the planet in the active region if the accretion flow there is rapid enough."," Therefore, in the cases where $\hdz/H=1.5$ and $\hdz/H=2.3$, we expect a $1\;M_J$ planet to be able to open a gap only in the dead-zone, possibly leaving the gas flowing over the top of the planet in the active region if the accretion flow there is rapid enough."
+ Snapshots of the disk density in the equatorial plane. and at the disk surface. are depicted in Fig. 3..," Snapshots of the disk density in the equatorial plane, and at the disk surface, are depicted in Fig. \ref{jupiter2}."
+" These plots correspond to atime when the mass of the accreting protoplanet has reached |M, and are for a model in which Zpz=2.3H.", These plots correspond to a time when the mass of the accreting protoplanet has reached $1\;M_J$ and are for a model in which $\hdz=2.3H$.
+ These reveal that the planet truncates the disk throughout its vertical extent. despite the high value for the viscosity in the active region.," These reveal that the planet truncates the disk throughout its vertical extent, despite the high value for the viscosity in the active region."
+ As can be seen in Fig. 4..," As can be seen in Fig. \ref{velocity},"
+ such a process arises because the planet is able to pull some of the gas from the live-zone down in toward the midplane as the latter tries to flow over the top of the planet., such a process arises because the planet is able to pull some of the gas from the live-zone down in toward the midplane as the latter tries to flow over the top of the planet.
+ This suggests that both the gap structure and the accretion rate onto the planet do not depend strongly on the presence of a dead-zone. but rather on the vertically-integrated accretion rate of gas through the disk.," This suggests that both the gap structure and the accretion rate onto the planet do not depend strongly on the presence of a dead-zone, but rather on the vertically-integrated accretion rate of gas through the disk."
+ In that case. we would expect the results from these 3D runs to not significantly differ from 2D calculations for a given value of the accretion rate M.," In that case, we would expect the results from these 3D runs to not significantly differ from 2D calculations for a given value of the accretion rate $\dot M$."
+ To investigate this issue in more detail. we have performed a suite of 2D runs with varying values of α from a=0 to a=14x107 and for which the initial surface density at the planet position is X(Ro)=Xo.," To investigate this issue in more detail, we have performed a suite of 2D runs with varying values of $\alpha$ from $\alpha=0$ to $\alpha=1.4\times 10^{-3}$ and for which the initial surface density at the planet position is $\Sigma(R_0)=\Sigma_0$."
+ An additional 2D calculation with w=2.8x107 and X(Ro)=0.55. has also been performed., An additional 2D calculation with $\alpha=2.8\times 10^{-3}$ and $\Sigma(R_0)=0.5\Sigma_0$ has also been performed.
+ The results of these calculations are shown in Fig. 5..., The results of these calculations are shown in Fig. \ref{2d3daccret}.
+ Also displayed are the results of two 3D simulations corresponding to: i) a fully active disk with a=[.4x107 throughout: ii) a dead disk with a=O everywhere., Also displayed are the results of two 3D simulations corresponding to: i) a fully active disk with $\alpha=1.4 \times 10^{-3}$ throughout; ii) a dead disk with $\alpha=0$ everywhere.
+ As expected. 2D runs with a#0 predict that the protoplanet grows faster as the viscosity increases.," As expected, 2D runs with $\alpha \ne 0$ predict that the protoplanet grows faster as the viscosity increases."
+ The run with «=2.8x107 shows slower growth at earlier times since the surface density is smaller in that case., The run with $\alpha=2.8\times 10^{-3}$ shows slower growth at earlier times since the surface density is smaller in that case.
+ However. once the planet opens a gap in the disk. it is clear that the growth rate becomes similar to that obtained from the run with a=1.4107 and for which the value for the mass accretion rate through the disk is the same.," However, once the planet opens a gap in the disk, it is clear that the growth rate becomes similar to that obtained from the run with ${\alpha=1.4\times10^{-3}}$ and for which the value for the mass accretion rate through the disk is the same."
+" In the case Where a=0. viscous supply of gas through the disk is inhibited and the planet mass can saturate once the feeding zone is empty of gas. which occurs when i,~0.5M, for the disk parameters used in this work."," In the case where $\alpha=0$, viscous supply of gas through the disk is inhibited and the planet mass can saturate once the feeding zone is empty of gas, which occurs when $m_p\sim 0.5 M_J$ for the disk parameters used in this work."
+ Clearly. for both ae=0 and Y=1.4x107. there is a good agreement between 2D and 3D calculations.," Clearly, for both $\alpha=0$ and $\alpha=1.4\times 10^{-3}$, there is a good agreement between 2D and 3D calculations."
+ This confirms that the accretion rate onto a gap-opening planet does not depend strongly on the details of the flow in the (RK.=) plane. but only on the vertically-averaged aceretion rate through the disk.," This confirms that the accretion rate onto a gap-opening planet does not depend strongly on the details of the flow in the $(R,z)$ plane, but only on the vertically-averaged accretion rate through the disk."
+ In order to study the evolution of a Saturn-mass planet embedded in a layered protoplanetary disk. we restarted the simulations described in the previous section when the mass of the accreting planet had reached the relevant value.," In order to study the evolution of a Saturn-mass planet embedded in a layered protoplanetary disk, we restarted the simulations described in the previous section when the mass of the accreting planet had reached the relevant value."
+ We then released the planet and let it evolve freely under the action of disk The evolution of Saturn's semi-major axis for the different models is shown in the upper panel of Fig. 6.., We then released the planet and let it evolve freely under the action of disk The evolution of Saturn's semi-major axis for the different models is shown in the upper panel of Fig. \ref{saturn}.
+ For each value of Hpz/H. Saturn does not undergo runaway migration (Masset Papaloizou 2003) because the disk is not massive enough.," For each value of $\hdz/H$, Saturn does not undergo runaway migration (Masset Papaloizou 2003) because the disk is not massive enough."
+ Instead. the migration rate 1s intermediate between the Type IL and Type I regimes.," Instead, the migration rate is intermediate between the Type II and Type I regimes."
+ This arises because the mass of Saturn is insufficient to enable a clean gap to form in the disk. as illustrated in Fig. 8..," This arises because the mass of Saturn is insufficient to enable a clean gap to form in the disk, as illustrated in Fig. \ref{saturn2}."
+ This figure displays. for the model with Hpz/H=2.3. a snapshot of the Saturn-induced gap structure in both the equatorial plane and at the disk surface.," This figure displays, for the model with $\hdz/H=2.3$, a snapshot of the Saturn-induced gap structure in both the equatorial plane and at the disk surface."
+ It ts worthwhile noticing that gap opening occurs in the live-zone. despite the high value for the viscosity there. indicating that both Jupiter and Saturn can pull gas from the active region down toward the disk For models with Zp57/Hx1.5. the evolution of the system proceeds quite similarly. with a slight tendeney for the migration rate to decrease as the size of the dead-zone increases.," It is worthwhile noticing that gap opening occurs in the live-zone, despite the high value for the viscosity there, indicating that both Jupiter and Saturn can pull gas from the active region down toward the disk For models with $\hdz/H \le 1.5$, the evolution of the system proceeds quite similarly, with a slight tendency for the migration rate to decrease as the size of the dead-zone increases."
+ For the calculation with Apy/H=2.3 however. we see that in ~200 orbits the semi-major axis has decreased by an amount that is nearly 20% less than for the other cases.," For the calculation with $\hdz/H=2.3$ however, we see that in $\sim 200$ orbits the semi-major axis has decreased by an amount that is nearly $20\%$ less than for the other cases."
+ This can be confirmed by examining the total disk torques exerted on the planet. which are displayed for each model in the lower panel of Fig. 6..," This can be confirmed by examining the total disk torques exerted on the planet, which are displayed for each model in the lower panel of Fig. \ref{saturn}."
+ It is well known that the total tidal torque. I. exerted on the planet can be written as the sum of a differential Lindblad torque. L;. and a total corotation torque. Lc (Tanaka Ward 2002).," It is well known that the total tidal torque, $\Gamma$, exerted on the planet can be written as the sum of a differential Lindblad torque, $\Gamma_L$ , and a total corotation torque, $\Gamma_C$ (Tanaka Ward 2002)."
+ In contrast to ΓΕ). which is basically independent of viscosity (Meyer-Vernet Sicardy 1987; Papaloizou Lin1984). [¢ depends strongly on v and can saturate in the low- limit.," In contrast to $\Gamma_L$, which is basically independent of viscosity (Meyer-Vernet Sicardy 1987; Papaloizou Lin1984), $\Gamma_C$ depends strongly on $\nu$ and can saturate in the low-viscosity limit."
+ In the present case. Lc can be decomposed as: where P24 (resp.," In the present case, $\Gamma_C$ can be decomposed as: where $\Gamma_C^{DZ}$ (resp."
+ L4) is the corotation torque due tothe librating fluid elements which originate from the dead (resp., $\Gamma_C^{AZ}$ ) is the corotation torque due tothe librating fluid elements which originate from the dead (resp.
+'Tidal interactions are important1 in determiningη the fate of short-periocd. extrasolar planets and the spinsο of their. host stars.,Tidal interactions are important in determining the fate of short-period extrasolar planets and the spins of their host stars.
+ nPhe extent of the spin-orbit⊀ evolution. that results from⋅ tides. depends on the dissipative∢⊀∢ properties. of the star and planet., The extent of the spin-orbit evolution that results from tides depends on the dissipative properties of the star and planet.
+ κ.These are usually parametrized. by a dimensionless. quality ⋠⋅factor for⋅ cach body. which ..is an inverse: measure of Le1ο dissipation.," These are usually parametrized by a dimensionless quality factor for each body, which is an inverse measure of the dissipation."
+"dissipati1 ""hisThis isi usuallylv defined to tolbe prolproportionaltional tto the ratio of the maximum energy stored in a tidal oscillation to the energy. dissipated over one cvele (e.g. Soter 1966)).", This is usually defined to be proportional to the ratio of the maximum energy stored in a tidal oscillation to the energy dissipated over one cycle (e.g. \citealt{GoldSot1966}) ).
+ In principal the modified quality (.E depends on tidal frequency. the internal structure of the. body. and the amplitude. of the tidal. forcing.," In principal the modified quality $Q^{\prime}$ depends on tidal frequency, the internal structure of the body, and the amplitude of the tidal forcing."
+"⋅⊀ ⇁⋅Unfortunately. the mechanisms. of ⋅⊀⊀⊀dissipationthat contribute⊀ to (∕ are poorlv usos"" L0 ⊔⊔"," Unfortunately, the mechanisms of dissipationthat contribute to $Q^{\prime}$ are poorly understood."
+∠⇂∢⊾↓⋅⊳∖↿∪⋯⇂⊳↓∪⊳∖↓⊔↓↓≻↓↓⇂∙∖⇁↿⇂⊔⋡∖∠∐∐∐↛⊔∐↓≻↓⋅∪∣⋡↓⋖⋅⊔↓⊳⊔∢⊾⋜⊔⋅↓∙∖⇁⋜↧∐qur studies. neglect any. amplitude. dependence of⋅ the dissipationREN (except. for: example Goodman..&Lackneronn 2009)).," To simplify this difficult problem, nearly all studies neglect any amplitude dependence of the dissipation (except, for example \citealt{GoodmanLackner2009}) )."
+" Such. studies. already exhibit""A a complicated. dependence of ⋅∕( on the tidal. frequency.⋅ (c.g Savonije↴»&Papaloizou.-1997:: 2005: Ogilvie&Lin 2004.. hereafter OLO4: Ogilvie&Lin 2007.. hereafter OLOT: Papaloizou&Ivanov 2010)). and the"," Such studies already exhibit a complicated dependence of $Q^{\prime}$ on the tidal frequency (e.g \citealt{SavPap1997};\citealt{Wu2005b}; \citealt{Gio2004}, , hereafter OL04; \citealt{Gio2007}, , hereafter OL07; \citealt{PapIv2010}) ), and the"
+the rectangular HEAO A-1 error boxes using the coordinates reported by ?/ and overplotted the EPIC pn fields together with the position of the AGN which constitute our sample.,the rectangular HEAO A-1 error boxes using the coordinates reported by \citet{wood84} and overplotted the EPIC pn fields together with the position of the AGN which constitute our sample.
+ The results are shown in Fig. 6.., The results are shown in Fig. \ref{heaofields}.
+ With the only exception of NGC 4968. all the sources fall largely outside of the LASS error box.," With the only exception of NGC 4968, all the sources fall largely outside of the LASS error box."
+ Note that the actual error region is much smaller than the latter. because the LMA identifications also take into consideration the MC pointings.," Note that the actual error region is much smaller than the latter, because the LMA identifications also take into consideration the MC pointings."
+ One of the sources. IRAS 13218+0552. was part of a 22-objects subsample of the LMA analyzed in detail by ?.. who included a figure very similar to the one presented here.," One of the sources, IRAS 13218+0552, was part of a 22-objects subsample of the LMA analyzed in detail by \citet{rem93}, who included a figure very similar to the one presented here."
+ Nonetheless. he suggested that in this case (and in some others) the LASS error boxes should be considered as =50% contours.," Nonetheless, he suggested that in this case (and in some others) the LASS error boxes should be considered as $\simeq50\%$ contours."
+ However. it is clear that these identifications are much less robust than for other LMA sources.," However, it is clear that these identifications are much less robust than for other LMA sources."
+ If this is the case. we should also look for contamination by nearby sources. which could contribute to the total flux measured by HEAO A-1.," If this is the case, we should also look for contamination by nearby sources, which could contribute to the total flux measured by HEAO A-1."
+ The LASS error boxes shown in Fig., The LASS error boxes shown in Fig.
+ 6 are overplotted on the RASS images of the same fields., \ref{heaofields} are overplotted on the RASS images of the same fields.
+ o bright source Hes inside any of the HEAO A-1 boxes., No bright source lies inside any of the HEAO A-1 boxes.
+ Moreover. we measured fluxes of the brightest objects present n the EPIC fields of each target source.," Moreover, we measured fluxes of the brightest objects present in the EPIC fields of each target source."
+ The results are shown n Table 6:: none of these sources have X-ray fluxes comparable to the one measured by the LMA and are all (with the exception of NVSS J1324314+053254 and SDSS J132442.444-052438.9 in the otherwise empty field of IRAS 13218+0552). much dimmer than the AGN.," The results are shown in Table \ref{xsources}: none of these sources have X-ray fluxes comparable to the one measured by the LMA and are all (with the exception of NVSS J132431+053254 and SDSS J132442.44+052438.9 in the otherwise empty field of IRAS 13218+0552), much dimmer than the AGN."
+ Even if some of them can be highly, Even if some of them can be highly
+di not confirm the8.15 60jun detection. ancl did: not detect the object in any other band.,"did not confirm the $60\mu$ m detection, and did not detect the object in any other band."
+ As we did not detect this source in the summ. its reality remains to be determined.," As we did not detect this source in the sub-mm, its reality remains to be determined."
+ Wit1 onlv one LR lux. it is impossible to draw. meaningful constraints on either luminosities or mocel parameters. we therefore only present a compilation of It data for this object.," With only one IR flux, it is impossible to draw meaningful constraints on either luminosities or model parameters, we therefore only present a compilation of IR data for this object."
+ Worth noting however is that all the model fits are consistent with a total Lt luminosity of 1077L. or less., Worth noting however is that all the model fits are consistent with a total IR luminosity of $10^{13.5}L_{\odot}$ or less.
+ The contributions from starburst anc AGN activity to t total Hi emission in the mos luminousLRAS galaxies has been discussed. by several authors., The contributions from starburst and AGN activity to the total IR emission in the most luminous galaxies has been discussed by several authors.
+ Alicl-LR spectroscopy (Genzeletal.1998:Bügopoulouοἱ1999) found that the Hi emission in most ULItis (o SO%) was powered Dy starbursts. although at least half of their samples showed evidence for both starburst anc ACN activity.," Mid-IR spectroscopy \cite{gen,rig} found that the IR emission in most ULIRGs $\sim80\%$ ) was powered by starbursts, although at least half of their samples showed evidence for both starburst and AGN activity."
+ Phere was no deteced trend for the AC-like svstems to reside in the more compact (and hence more advanced merger) systems., There was no detected trend for the AGN-like systems to reside in the more compact (and hence more advanced merger) systems.
+ The fracion of sources powered by an AGN has been shown ὃν several authors to increase with increasing total LH uniinosiv (Veilleuxetαἱ.1995:Shier.Ricke&Rieke1996:Veilleux.Sanders&Wim 1999).," The fraction of sources powered by an AGN has been shown by several authors to increase with increasing total IR luminosity \cite{vei,shi,vei2}. ."
+. Recent observations of ULIRGs Clranetal.2001) found that. at LEX. luminosities low. 10124L.. most ULLIBCs were starburst) ciominated. with the starburst component contributing around δρ to he LR emission.," Recent observations of ULIRGs \cite{tra2} found that, at IR luminosities below $10^{12.4}L_{\odot}$, most ULIRGs were starburst dominated, with the starburst component contributing around $85\%$ to the IR emission."
+ At LR luminositics above 101775. the AGN contribution was much higher. contributing at least 50% of he LR emission.," At IR luminosities above $10^{12.4}L_{\odot}$ the AGN contribution was much higher, contributing at least $50\%$ of the IR emission."
+ Studies of samples of LILERGs to look for contributions rom starbursts and AGN to the total LR. emission. have vicleled interesting results., Studies of samples of HLIRGs to look for contributions from starbursts and AGN to the total IR emission have yielded interesting results.
+ Studies. of small samples of LILIRGs produced. contlieting results. with some authors sugeesting that the LR emission in LILIRCs arises predominantly from a starburst wih star formation rates of the order 10Αν + and other authors suggesting that LILIRGs are powered by a cdusty AGN (Granato.Danese&Franceschini1996:Evanset 1998).," Studies of small samples of HLIRGs produced conflicting results, with some authors suggesting that the IR emission in HLIRGs arises predominantly from a starburst with star formation rates of the order $10^{3}M_{\odot}$ $^{-1}$ , and other authors suggesting that HLIRGs are powered by a dusty AGN \cite{gra2,eva,yun}."
+. Phere is also evidence from spectropolarimetry for a buried. αν in several Ετος (linesetal.1995:Goodrich1996). Previous SED modelling o ΠΕανν (Rowan-Robinson2000:Vermactal.2002) founc that most required an ACN ane Lastarburst to power the total It emission.," There is also evidence from spectropolarimetry for a buried AGN in several HLIRGs \cite{hin0,goo}
+ Previous SED modelling of HLIRGs \cite{rr2,ver} found that most required an AGN and a starburst to power the total IR emission."
+ Overall. about wll were AG dominated.," Overall, about half were AGN dominated."
+ These results were however based. on samples either biased. toward AGN or lacking the tight constraints on starburst Iuminosities given by sub-mum cata., These results were however based on samples either biased toward AGN or lacking the tight constraints on starburst luminosities given by sub-mm data.
+ The total Hi. luminosities of the 10 reiably detected jects in our sample. together wi hthesarburst and AG [ractions. are presented in Table 3..," The total IR luminosities of the 10 reliably detected objects in our sample, together with the starburst and AGN fractions, are presented in Table \ref{hlirgparams}."
+" In all ten objects both a starburst and an AGN are required. to explain the total IR emission: in all cases a ""pure! starburs or AGN model is unable to explain the emission over 1.1 ΠΠ The mean starburst fraction in our sample is 354.", In all ten objects both a starburst and an AGN are required to explain the total IR emission; in all cases a `pure' starburst or AGN model is unable to explain the emission over $1-1000\mu$ m. The mean starburst fraction in our sample is $\sim35\%$.
+ In all ten objects both the AGN ancl starburst components supply at least ‘Of the totalER. luminosity. any ofthe objects would," In all ten objects both the AGN and starburst components supply at least $20\%$ of the totalIR luminosity, any ofthe objects would"
+"(Thy, &CIT, (πηραοςetal.","$_{eff}$ $\lesssim$$_4$ \citep{kir99,me02,geb02}."
+" CIT, IIO II (Tsuji.Oliuaka. jun. (Nakajima (Cubyetal.1999:Liu (Abt&Levy1976:AbtLOST:DuqueunoyMayor1991).. (Isocrueretal1999:"," \citep{kir99,mrt99} $_4$ $_2$ $_2$ \citep{tsu99,bur00,lie00a}, $\micron$ \citep{nak95,opp95}, \citep[hereafter SDSS; York et al.\
+2000]{str99,tsv00,leg00,geb02}, \citep{cub99,liu02,zap02}, \citep{nak95,me00a,els01}. \citep{abt76,abt87,duq91}. \citep{hen90,fis92,rei97}. \citep{koe99,rei01a,clo02}."
+ILL objects and II» 2.12 eemission features. are presented in Table 3..,"HH objects and $_2$ 2.12 emission features, are presented in Table \ref{tab3}."
+ In this paper we designate all the detected IRAC mid-infrared outflows as EGOs., In this paper we designate all the detected IRAC mid-infrared outflows as EGOs.
+ The EGOs are numbered by order in right. ascension., The EGOs are numbered by order in right ascension.
+ HW the overall morphology. of a eroup of knots suggests physical association among them. the knots are considered as parts ol one EGO and the individual knots are distinguished with additional letters to the EGO numbers.," If the overall morphology of a group of knots suggests physical association among them, the knots are considered as parts of one EGO and the individual knots are distinguished with additional letters to the EGO numbers."
+ On the other haud. if the knots in a region are relatively separated [rom each other and the relationship among them is not clear. each knot is (treated as an individual EGO.," On the other hand, if the knots in a region are relatively separated from each other and the relationship among them is not clear, each knot is treated as an individual EGO."
+ Apparent sub-structures in a knot are labelled with additional numerals. such as EGO 2181.," Apparent sub-structures in a knot are labelled with additional numerals, such as EGO 21b1."
+ The locations of all the mid-infrared. outflows identified in the p Ophiuchi cloud are shown in Fig. 2.., The locations of all the mid-infrared outflows identified in the $\rho$ Ophiuchi cloud are shown in Fig. \ref{fig2}.
+ The grevscale image in Fig., The greyscale image in Fig.
+ 2. is the [RAC channel 2 (4.5 san)) image., \ref{fig2} is the IRAC channel 2 (4.5 ) image.
+ The outflows newly discovered in (his work are marked with circles and the counterparts of known flows are labelled with pluses., The outflows newly discovered in this work are marked with circles and the counterparts of known flows are labelled with pluses.
+ Two YSO aggregates in the region. LIGS9- and L1709- (Padgettetal.2008).. are marked with open squares.," Two YSO aggregates in the region, L1689- and L1709-aggregate \citep{pad08}, are marked with open squares."
+ The identilied mid-inlrared outflows are located mainlv in three regions. ie.. the L1688 core. the IRAS 16293-2422 region. and the LIT0O YSO aggregate region.," The identified mid-infrared outflows are located mainly in three regions, i.e., the L1688 core, the IRAS 16293-2422 region, and the L1709 YSO aggregate region."
+ The images of the newly discovered IRAC! mid-inlrared oulllows are shown in Figs. 3--17..," The images of the newly discovered IRAC mid-infrared outflows are shown in Figs. \ref{fig3}- \ref{fig17},"
+ and (he images of mid-inlraved counterparts of HII objects and LI emission features are presented in Figs. 18--40.., and the images of mid-infrared counterparts of HH objects and $_2$ emission features are presented in Figs. \ref{fig18}- \ref{fig40}. .
+ Thirteen new mic-intrarecl outflows have been discovered in our survey., Thirteen new mid-infrared outflows have been discovered in our survey.
+ For each outflow. we present the channel 2 (4.5 jm)) minus 1 (3.6 jm)) image and the three-color image.," For each mid-infrared outflow, we present the channel 2 (4.5 ) minus 1 (3.6 ) image and the three-color image."
+ Figures 32-4 show the region of objects EGO 04 and EGO 12., Figures 3-4 show the region of objects EGO 04 and EGO 12.
+ The distance between the {wo objects is about3°., The distance between the two objects is about.
+ Each object consists of two knots., Each object consists of two knots.
+ EGO 04 and EGO 12 are only detected in TRAC channels 1 (3.6 jam)) and 2 (4.5 μι)., EGO 04 and EGO 12 are only detected in IRAC channels 1 (3.6 ) and 2 (4.5 ).
+ It is brighter in channel 2 (4.5 therefore. appears as an EGO in the three-color image (Fig. 4)).," It is brighter in channel 2 (4.5 ), therefore, appears as an EGO in the three-color image (Fig. \ref{fig4}) )."
+ There are three YSOs to the southeast of EGO 04 and EGO 12. among which we suggest DINLT. J162624-241616 (Class ll in Evansetal. 2009)) as the exciting source of the EGO 04 and EGO 12 outflow. as this source is also detected in the MIPS 24 image.," There are three YSOs to the southeast of EGO 04 and EGO 12, among which we suggest BKLT J162624-241616 (Class II in \citealt{eva09}) ) as the exciting source of the EGO 04 and EGO 12 outflow, as this source is also detected in the MIPS 24 image."
+ The distance from DINLT J162624-241616 to EGO 04 is 4°77 and the distance to EGO 12 is 1/88., The distance from BKLT J162624-241616 to EGO 04 is 7 and the distance to EGO 12 is 8.
+ Outflow EGO 33 (Figs., Outflow EGO 33 (Figs.
+ 5 and 6) is a diffuse nebula and is visible in IRAC channels 2 (4.5 pm)). 3 (5.8 pom)). and 4 (8.0 ," 5 and 6) is a diffuse nebula and is visible in IRAC channels 2 (4.5 ), 3 (5.8 ), and 4 (8.0 )."
+ΕΕΠΠ counterparts of the known near- outflow in this region. three IH» emission features in Khanzadvanetal. (2004)... are also detected in (he IRAC images.," Mid-infrared counterparts of the known near-infrared outflow in this region, three $_{2}$ emission features in \citet{kha04}, , are also detected in the IRAC images."
+ We designated these mid-infraredcounterparts as EGO, We designated these mid-infraredcounterparts as EGO
+AIC map.,MC map.
+ However. the POLL map shows clear signs of nonnormality.," However, the TOH map shows clear signs of non–normality."
+ The nonnormality is clisplaved from larger scales than the univariate counterparts., The non–normality is displayed from larger scales than the univariate counterparts.
+" That is to. say. he distribution of the @p,s appears unusual at scales ereater than (=300."," That is to say, the distribution of the $a_{lm}$ s appears unusual at scales greater than $\ell=300$."
+ The results from the multivariate power transformations using harmonic space data are shown in Figure 23.., The results from the multivariate power transformations using harmonic space data are shown in Figure \ref{fig:multi_transSPH}.
+ Looking at the results from the ILC ane Gaussian MC maps. it appears that we are failing to find he true global minimum.," Looking at the results from the ILC and Gaussian MC maps, it appears that we are failing to find the true global minimum."
+ As discussed earlier. finding a local minimum will result in an underestiniate of the value of X5.," As discussed earlier, finding a local minimum will result in an underestimate of the value of $\chi^2_2$."
+ Therefore. we are not too concerned. about this as failure o [lind global minima since it [aibure docs not result in false claims of non.Craussianity.," Therefore, we are not too concerned about this as failure to find global minima since it failure does not result in false claims of non–Gaussianity."
+ Intriguingly. the shape of the line or the POLL map does not match those of the other two maps.," Intriguingly, the shape of the line for the TOH map does not match those of the other two maps."
+ Lowe are failing to find the global minima then the we would expect a larger number of points to be bevond the confidence region if we corrected. this failure., If we are failing to find the global minima then the we would expect a larger number of points to be beyond the confidence region if we corrected this failure.
+ Equally. the result may reflect that the cistributions extracted from the TOLL maps make finding the elobal minimum easier because they are. to some extent. smoother.," Equally, the result may reflect that the distributions extracted from the TOH maps make finding the global minimum easier because they are, to some extent, smoother."
+ FUV spectroscopy of the five dwarf novae was obtained with MS7/STIS duringUST Cvcle Ll., FUV spectroscopy of the five dwarf novae was obtained with /STIS during Cycle 11.
+" The data were obtained using the G140L. erating and the 52""x0.2"" aperture. providing a spectral resolution of R~1000 over the wavelength range 1140-1720À."," The data were obtained using the G140L grating and the $52^{\prime\prime} \times 0.2^{\prime\prime}$ aperture, providing a spectral resolution of $\sim 1000$ over the wavelength range 1140-1720."
+. Since the total time involved in each snapshot observation was short (<35min). the observations were made in the ACCUM mode in order to minimize the instrument overheads.," Since the total time involved in each snapshot observation was short $< 35$ min), the observations were made in the ACCUM mode in order to minimize the instrument overheads."
+ This resulted. in exposure times of 600 to 900 seconds., This resulted in exposure times of 600 to 900 seconds.
+ Each snapshot observation resulted in a single time averaged spectrum of each dwarf nova., Each snapshot observation resulted in a single time averaged spectrum of each dwarf nova.
+ All of the spectral data were processed wilh IRAF using CALSTIS V2.13b., All of the spectral data were processed with IRAF using CALSTIS V2.13b.
+"During target acquisition./Z/5T points at the nominal target coordinates and takes a5”x5"" CCD image with an exposure time of a few seconds.","During target acquisition, points at the nominal target coordinates and takes a $5^{\prime\prime} \times 5^{\prime\prime}$ CCD image with an exposure time of a few seconds."
+ Subsequently. a small slew is performed that centers (he target in the acquisition box. and a second CCD image is taken.," Subsequently, a small slew is performed that centers the target in the acquisition box, and a second CCD image is taken."
+ The acquisition images for these observations were obtained using the F23x50LP lone-pass filter. which has a central wavelength of 7228.5 and a at half maximum (EWIIM) of citepara(5..," The acquisition images for these observations were obtained using the F28x50LP long-pass filter, which has a central wavelength of 7228.5 and a full-width at half maximum (FWHM) of \\citep{ara05}. ."
+ We caleulated the F28x50LP magnitudes as in Aranjo-Betancorοἱal.(2005)., We calculated the F28x50LP magnitudes as in \citet{ara05}.
+. The instrumental setup aud exposure details of theLST STIS spectra of TU Men. BD Pav. SS Aur. TT Crt. and V442 Cen are provided in the observing log given in Table 2. the entries are bv column: (1) the target. (2) Data ID. (3) Instrument Config/Moce. (4) Grating. (5) Date of Observation (vvyyv-amnnmne-dd). (6) Time of observation. (7) Exposure time (s).aud (3) F28x50LD magnitude.," The instrumental setup and exposure details of the STIS spectra of TU Men, BD Pav, SS Aur, TT Crt, and V442 Cen are provided in the observing log given in Table 2, the entries are by column: (1) the target, (2) Data ID, (3) Instrument Config/Mode, (4) Grating, (5) Date of Observation (yyyy-mm-dd), (6) Time of observation, (7) Exposure time (s),and (8) F28x50LP magnitude."
+cross-identified with the OGLE catalogue of RRL stars in the LAIC (Soszvisski et al.,cross-identified with the OGLE catalogue of RRL stars in the LMC (Soszyńsski et al.
+ 2003)., 2003).
+ The position of all identified stars in the A. J—A color-magnitude diagram is shown on Fie. 2..," The position of all identified stars in the $K$ , $J-K$ color-magnitude diagram is shown on Fig. \ref{figcmd}."
+ Most of the stus have only one randonm-phase measurement. bul some have (wo random-phase observations taken on the two different nights.," Most of the stars have only one random-phase measurement, but some have two random-phase observations taken on the two different nights."
+ A few of the RRL stars were cross-identified in the overlapping fields. which provided three measurements for these objects.," A few of the RRL stars were cross-identified in the overlapping fields, which provided three measurements for these objects."
+ The period-Iuminositw (PL) relations for the J and A bands derived [rom our data are shown in Fig. 3.., The period-luminosity (PL) relations for the $J$ and $K$ bands derived from our data are shown in Fig. \ref{figplav}.
+ For the RRL stars with more than one observation. we took a straight average of (he random-phase magnitudes. which should lead to a better approximation of (heir mean magnitudes.," For the RRL stars with more than one observation, we took a straight average of the random-phase magnitudes, which should lead to a better approximation of their mean magnitudes."
+ In Fig. 3..," In Fig. \ref{figplav},"
+ (wo distinct groups can be distinguished (hat correspond to (he first overtone (Re) and fundamental mode (IItab) pulsators. respectively.," two distinct groups can be distinguished that correspond to the first overtone (RRc) and fundamental mode (RRab) pulsators, respectively."
+ The relatively large scatters seen in both figures 2 and 3 is mostly caused by three factors: 1) the random V.ingle-phase nature of our It measurements. which represents ihe mean magnitude of an RRL variable only to ~0.15 mag (e.g. Del Principe et al.," The relatively large scatters seen in both figures \ref{figcmd} and \ref{figplav} is mostly caused by three factors: 1) the random single-phase nature of our IR measurements, which represents the mean magnitude of an RRL variable only to $\sim 0.15$ mag (e.g. Del Principe et al."
+ 2006). 2) the metallicity spread among RRL stars in the LAIC. and 3) the accuracy of our single measurements. which is 0.03-0.19 mag lor stars of brightnesses of 16.6-13.6 mag in the A band.," 2006), 2) the metallicity spread among RRL stars in the LMC, and 3) the accuracy of our single measurements, which is 0.03-0.19 mag for stars of brightnesses of 16.6-18.6 mag in the $K$ band."
+ since the most important contributor to the scatter in Figure 3. is the replacement of the mean magnitudes by single-phiase or by averaged [ew-phase measurements. we used accurate optical (DVZ7) photometry from (the OGLE archive (Soszvisski et al.," Since the most important contributor to the scatter in Figure \ref{figplav} is the replacement of the mean magnitudes by single-phase or by averaged few-phase measurements, we used accurate optical $BVI$ ) photometry from the OGLE archive (Soszyńsski et al."
+ 2003) of the the studied sample of RRL stars. taken close in time to the NIR data.," 2003) of the the studied sample of RRL stars, taken close in time to the NIR data."
+ With the additional oplical data. we can caleulate improved JA phase points in order to use (he Jones et al. (," With the additional optical data, we can calculate improved $JK$ phase points in order to use the Jones et al. ("
+1996) template light curve method. (available only for the A band).,1996) template light curve method (available only for the $K$ band).
+ The J. N. and (A) magnitudes are given in Table 5..," The $J$, $K$, and $\langle K \rangle$ magnitudes are given in Table \ref{tabavobs}."
+ We compare (he PL relations derived for the RRab stars (averaged observations) and for all of the RRL stars (RRab and RRe) against the existing theoretical (Bono et al., We compare the PL relations derived for the RRab stars (averaged observations) and for all of the RRL stars (RRab and RRc) against the existing theoretical (Bono et al.
+ 2003h: Catelan et al., 2003b; Catelan et al.
+ 2004) and empirical (Sollima et al., 2004) and empirical (Sollima et al.
+ 2006: 2008) relations., 2006; 2008) relations.
+ Table ο lists our results for the slope and zero point values of the PL relations as well as those from the theoretical and empirical ones., Table \ref{tabplav} lists our results for the slope and zero point values of the PL relations as well as those from the theoretical and empirical ones.
+ Furthermore. in Table 6.. we present the K-band zero point and slope values obtained from (A) magnitudes derived from the Jones et al. (," Furthermore, in Table \ref{tabplmean}, we present the K-band zero point and slope values obtained from $\langle K \rangle$ magnitudes derived from the Jones et al. ("
+1996) light curve template method.,1996) light curve template method.
+ In order to derive the apparent distance modulito LAIC from our data. we used the Following calibrations of the near-infrared PL relations of mixed population RRL stars:," In order to derive the apparent distance modulito LMC from our data, we used the following calibrations of the near-infrared PL relations of mixed population RRL stars:"
+where 0 djs the polar angle from thle axis of Μποταν. as seen by the star at the coordinate origin.,"where $\theta$ is the polar angle from the axis of symmetry, as seen by the star at the coordinate origin."
+" Ry is the so-called standoff distance obtained by balancing the ram pressures of the stellar wind and ambicut medium at 0—0 anc ix elven by: Tere ai, is the stellar mass-loss rate. c, the texiuinal velocity of the stellar wind. p, the ambient medi mass density. and ο is the velocity at which the star moves through the media (ie. the relative velocity of the ziubieut iiediu and the star velocity ii case that the züubieut medium is not stationary)."," $R_0$ is the so-called standoff distance obtained by balancing the ram pressures of the stellar wind and ambient medium at $\theta$ =0 and is given by: Here $\dot{m}_w$ is the stellar mass-loss rate, $v_{w}$ the terminal velocity of the stellar wind, $\rho_a$ the ambient medium mass density, and $v_{a}$ is the velocity at which the star moves through the medium (i.e. the relative velocity of the ambient medium and the star velocity in case that the ambient medium is not stationary)."
+ We asstune that the shell has a thickness as shown in Fie 3b in ?..," We assume that the shell has a thickness as shown in $\,$ 3b in \citet{maclow91}."
+ Then we rotate the two-dimeusional shape around its axis of sviuuetrv in order to obtaiu the dimensional shell., Then we rotate the two-dimensional shape around its axis of symmetry in order to obtain the three-dimensional shell.
+ The iodel bv ? incorporates theapproXnnatiou:: the interaction between the stellar wind aud the ambicut medimm takes place in an infinitely thin laver. the two flaws are fully mixed aud inuaneciately cooled.," The model by \citet{wilkin96PhD} incorporates the: the interaction between the stellar wind and the ambient medium takes place in an infinitely thin layer, the two flows are fully mixed and immediately cooled."
+ Iu this case Ry directly gives the distance of the star to the apex of the bow shock., In this case $R_0$ directly gives the distance of the star to the apex of the bow shock.
+ However. ? note that this might not be true if cooling of the shocked stellar wind is Inefiicicut.," However, \citet{comeron98} note that this might not be true if cooling of the shocked stellar wind is inefficient."
+ In this case it is expected that the bow shock would be located at a distance somewhat ercater than Ry (222).," In this case it is expected that the bow shock would be located at a distance somewhat greater than $R_0$ \citep{comeron98, raga97, povich08}."
+ For some GC sources. 2 indeed repored that a better fit can be obtained if the apex of the vow shock. is shifted away from the star.," For some GC sources, \citet{tanner05} indeed reported that a better fit can be obtained if the apex of the bow shock is shifted away from the star."
+ We lave observed the same behavior for several other Galactic Ceuter bow shocks., We have observed the same behavior for several other Galactic Center bow shocks.
+ ? have investigated the case where the wind ejection froun the star is not necessarily isotropic: it is rather confined in a cone of solid angle Q=27(1cos09). where 204 is the openings augle in which the matter is cjected.," \citet{zhang&zheng97} have investigated the case where the wind ejection from the star is not necessarily isotropic; it is rather confined in a cone of solid angle $\Omega=2\pi(1-\cos\theta_0)$, where $\theta_0$ is the opening angle in which the matter is ejected."
+" The staudotff distance is then eiveuby: As can he seen in Fie.3.. with 60, setting sialler. he bow shock profile will have a narrower shape."," The standoff distance is then givenby: As can be seen in $\,$ \ref{modelsketch}, with $\theta_0$ getting smaller, the bow shock profile will have a narrower shape."
+ For MW>3afl the shape is very similar to the oue eiven x 1..," For $\theta_0\rightarrow 3\pi/4$ the shape is very similar to the one given by $\,$ \ref{Req}."
+ Although this model has been developed for the case of a collimated stellar wind. we do not assune such a collimation. siuce it would inply that the collimation of he wind is in the direction of motion of the star. relative o the ambicut medium.," Although this model has been developed for the case of a collimated stellar wind, we do not assume such a collimation, since it would imply that the collimation of the wind is in the direction of motion of the star, relative to the ambient medium."
+ There is no a priori reason for he stellar velocity and wind collimation direction to he inked., There is no a priori reason for the stellar velocity and wind collimation direction to be linked.
+ We iutroduce Jy as an additional factor that allows one to control the width of the how shock outline. while causing the apex of the bow shock to be slightly displaced away frou the star.," We introduce $\theta_0$ as an additional factor that allows one to control the width of the bow shock outline, while causing the apex of the bow shock to be slightly displaced away from the star."
+ To eoenerate a simulated observation. we start by ilunimatiug the shell by the star placed at the origin of the coordinate system aud then calculate the euissiou along cach ray intersecting the shell.," To generate a simulated observation, we start by illuminating the shell by the star placed at the origin of the coordinate system and then calculate the emission along each ray intersecting the shell."
+ The shell is inclined to the line of sight by the angle 7. aud also rotated by a position anele PA in the plane parallel to the plane ofthe sky.," The shell is inclined to the line of sight by the angle $i$, and also rotated by a position angle PA in the plane parallel to the plane ofthe sky."
+ Tuset in Fie.3 explains the convention used for the inclination anele.," Inset in $\,$ \ref{modelsketch} explains the convention used for the inclination angle."
+ PA is measured Eas of North., PA is measured East of North.
+ Each parcel of the shel is assigned the optical depth τ. calculated as: We asstuue a eraphite|silicate mixture with a power-law erain size distribution n4(0) C2).. where o is the erain size.," Each parcel of the shell is assigned the optical depth $\tau$, calculated as: We assume a graphite+silicate mixture with a power-law grain size distribution $n_d(a)$ \citep{MRN77}, where $a$ is the grain size."
+ Dust extinction cocfiicicut is C4—m(Qus| Quai). where Qurs and (Qu are dust absorption and scattering efficiencies. respectively (7).1.," Dust extinction coefficient is $C_{ext}=\pi a^2(Q_{abs}+Q_{sca})$ , where $Q_{abs}$ and $Q_{sca}$ are dust absorption and scattering efficiencies, respectively \citep{laor&draine93}
+."
+ EL is the leusth of the shell parcel along the line of sight., $L$ is the length of the shell parcel along the line of sight.
+ The sonrce NT is polarized (sco7). sugeesting hat scattering of the stellar enuissiou by dust particles in the bow-shock euvelope is probably important.," The source X7 is polarized \citep[see][]{muzic07}, suggesting that scattering of the stellar emission by dust particles in the bow-shock envelope is probably important."
+ In order fo account for scattering. we proceed iu the ollowiug wav: Γαμος. of each parcel of the shell ji contributions both from scattering and thermal chussion: L—L.41!L;g.," In order to account for scattering, we proceed in the following way: Emission of each parcel of the shell has contributions both from scattering and thermal emission: $L$ $L_{sca}$ $L_{th}$ ."
+" LenxwDOMIevJey, 2d LosXd ενPOgaJe mre. where d ds the cistauce roni the star. Bi) is the black body emission at the dust temperature. iutegrated over grain sizes. and over wavelengths im our observing band."," $L_{th}\propto$$\,B(T_d)$$(1-e^{-\tau_{abs}})\epsilon_{th}$ and $L_{sca}\propto$ $^{-2}$$\epsilon_{sca}P(\theta_{sca})$$e^{-\tau_{sca}}$ , where $d$ is the distance from the star, $B(T_d)$ is the black body emission at the dust temperature, integrated over grain sizes, and over wavelengths in our observing band."
+ ἐν aud εκ are, $\epsilon_{th}$ and $\epsilon_{sca}$ are
+Each spectii has a narrow profile with |S and 9846 lan + for sinele-Caussian fits to tιο CÓ (21) aud CO (32) profiles. respectively.,"Each spectrum has a narrow profile with $\pm$ 8 and $\pm$ 6 km $^{-1}$ for single-Gaussian fits to the CO (2–1) and CO (3–2) profiles, respectively."
+ These xofiles are siguificailv narrower than the CO (21) enission from NGC 6210 (Figure 55) aud muplv a nearlv ACC-OLL orientation fo any disk-like structure iu this svseni., These profiles are significantly narrower than the CO (2–1) emission from NGC 6240 (Figure 5) and imply a nearly face-on orientation to any disk-like structure in this system.
+ CO measureneuts also provide estimates of the :unuouut of 1nolecular eas., 	 	CO measurements also provide estimates of the amount of molecular gas.
+ In Mrk 739. Lov) =109433 and 6902551 Jv khns! for the 21 aud 32 lines.," In Mrk 739, $I_{\mathrm{CO}}$ $\pm$ 33 and $\pm$ 51 Jy km $^{-1}$ for the 2–1 and 3–2 lines."
+ Following ? PagGo710.9«0 and 7.5«1 QUNEks!  NAfor he 21 and 32 lines.," Following \citet{Solomon:1992p11420}, $L'_{\mathrm{CO}}$ $\times10^{8}$ and $\times10^{8}$ K km $^{-1}$ $^{-2}$ for the 2–1 and 3–2 lines."
+ Adopting à—1.5-1 Af. (K lan 7) |! for the couversion frou CO. huuinositv. to molecular hydrogen. we find los AL. (112)=9.2-9.6. similar to the Alilkv Way (?)..," Adopting $\alpha$ =1.5-4 $M_\odot$ (K km $^{-2}$ $^{-1}$ for the conversion from CO luminosity to molecular hydrogen, we find $\log$ $M_\odot$ $_2$ )=9.2-9.6, similar to the Milky Way \citep{Sanders:1984p12189}."
+ Measurements of radial velocities provide important insights about the dvuamics of the merger., 	 	 	Measurements of radial velocities provide important insights about the dynamics of the merger.
+ We use the Na TAA 5890. 5896 (Na D) absorption lines from stars and cold gas since narrow cussion lines in ACN often," We use the Na I $\lambda\lambda$ 5890, 5896 (Na D) absorption lines from stars and cold gas since narrow emission lines in AGN often"
+Even casual examination shows hat most clisk galaxies are not truly symmetric but exhibit a variety of morphological peculiarities of which spiral arnas and bars are the most pronounced.,Even casual examination shows that most disk galaxies are not truly symmetric but exhibit a variety of morphological peculiarities of which spiral arms and bars are the most pronounced.
+ After decades of eTort. we know that these features may be driven by «nyvironmental clisturbance acting directly on the disk. in addition to scllexeitation of a local cüsturbance ον swing amplification. 'l'oomre. 1981. Sellwood Carlerg 1984)).," After decades of effort, we know that these features may be driven by environmental disturbance acting directly on the disk, in addition to self-excitation of a local disturbance by swing amplification, Toomre 1981, Sellwood Carlberg \nocite{Toom:81,SeCa:84}) )."
+ However. all clisks are embedded within halos and therefore are. not dynamically independent and. wi| respond to asymnmetries and distortions in the halo. as we‘ll.," However, all disks are embedded within halos and therefore are not dynamically independent and will respond to asymmetries and distortions in the halo, as well."
+ Until recently. οventional wisdom was that halos acted to stabilize clisks but. otlwise remained relatively inert.," Until recently, conventional wisdom was that halos acted to stabilize disks but otherwise remained relatively inert."
+ TΊο argument behind this assumption is as follows., The argument behind this assumption is as follows.
+ Llalos. s[xheroids and bulges are supported against their own gravity ow the random motion of their starsa. so-called shot” distribution (c.g. Binney Tremaine 1987)).," Halos, spheroids and bulges are supported against their own gravity by the random motion of their stars—a so-called “hot” distribution (e.g. Binney Tremaine \nocite{BiTr:87}) )."
+ On all but the largest. scales. they look like nearly homogeneous thermal xuthis ofstars.," On all but the largest scales, they look like nearly homogeneous thermal baths of stars."
+ Because all self-sustaining patterns or waves in a universe of stars with a Maxwellian velocity distribution are predicted: to damp quickly (c.g. Ieuchi. akamura Takahara 1974)). one expects that any pattern will be strongly caniped in halos and. spheroids as well.," Because all self-sustaining patterns or waves in a universe of stars with a Maxwellian velocity distribution are predicted to damp quickly (e.g. Ikeuchi, Nakamura Takahara \nocite{IkNT:74}) ), one expects that any pattern will be strongly damped in halos and spheroids as well."
+ Lowever. recent work sugeests that halos respond to tidal encounters by companions or cluster members and SUSCCsptible to induction of long-livecl moces due to their inhomogeneity.," However, recent work suggests that halos respond to tidal encounters by companions or cluster members and susceptible to induction of long-lived modes due to their inhomogeneity."
+ These modes arc at the largest scales [or which seIl eravity is most effective., These modes are at the largest scales for which self gravity is most effective.
+ In particular. if halos are largeὃν as many currently estimate. halo-halo interactions in eroups or clusters will be frequent and much more common than clisk-clisk interactions.," In particular, if halos are large as many currently estimate, halo-halo interactions in groups or clusters will be frequent and much more common than disk-disk interactions."
+ Because non-local coupling can, Because non-local coupling can
+which may. be coufused with others effects such as nou-Caussian iitial conditions (Fedelictal.2009).,which may be confused with others effects such as non-Gaussian initial conditions \citep{FMM09}.
+. Models with non-Gaussian initial conditions predict that the actual uuuber of clusters with higher mass cau be lower than the models with Caussian initial conditions (c.g.Figure1iuFedeletal. 2009): our nou-equipartition model predicts an apparent sinaller Y for high-mass clusters and this underestimates the nnuber of Ligh mass clusters if the non-equipartition effect is not taken mto account.," Models with non-Gaussian initial conditions predict that the actual number of clusters with higher mass can be lower than the models with Gaussian initial conditions \citep[e.g., Figure~1 in][]{FMM09}; our non-equipartition model predicts an apparent smaller $Y$ for high-mass clusters and this underestimates the number of high mass clusters if the non-equipartition effect is not taken into account."
+" Note that at such small levels of systematic deviations. otler systematic uncertaimties such as tle use of different mass fictions (6...Press&Schechter197I:Sheth&Tormen1999) may introduce larger svstematic biases,"," Note that at such small levels of systematic deviations, other systematic uncertainties such as the use of different mass functions \citep[e.g.,][]{PS74, ST99} may introduce larger systematic biases."
+ However. Fedelietal.(2009) also used the," However, \citet{FMM09} also used the"
+"3em According to the cubic approximation for S(7). ο(τν) isa linear function of οτι.1). Strap). S""(ryp4) and ο(τν). as shown by eq. (","3em According to the cubic approximation for $S(\tau)$ $S^{\prime}(\tau_{NL})$ is a linear function of $S(\tau_{NL-1})$ , $S(\tau_{NL})$, $S^{\prime\prime}(\tau_{NL-1})$ and $S^{\prime\prime}(\tau_{NL})$, as shown by eq. ("
+9).,9).
+" The “physical” equation (45) and the spline equation (9) lead to a new linear relation among οτι4). S""(yq,4). S(TN4) and ο(τν) that. together with ee.s (39). (40) and (44) lead easily to the explicit. values ol the latter four variables."," The ”physical” equation (45) and the spline equation (9) lead to a new linear relation among $S(\tau_{NL-1})$, $S^{\prime\prime}(\tau_{NL-1})$, $S(\tau_{NL})$ and $S^{\prime\prime}(\tau_{NL})$ that, together with eq.s (39), (40) and (44) lead easily to the explicit values of the latter four variables."
+ Consequently we easily obtain also the values of 5'(7y;) and οτν)., Consequently we easily obtain also the values of $S^{\prime}(\tau_{NL})$ and $S^{\prime\prime\prime}(\tau_{NL})$.
+ The explicit values of these variables αἱ τν allow us to compute those of the sel HEτνqu)) through eq. (," The explicit values of these variables at $\tau_{NL}$ allow us to compute those of the set $\lbrace I^{+}(\tau_{NL},\mu_{J})\rbrace$ through eq. ("
+5).,5).
+" Once the explicit values of (τν4). S(T«4). S""(rug4). ο(τν) as well as those of the sel (E(Tuipty.J=1.ND)} are known. it is straightforward to compute those of the set {L(tx,aqu.J=LEND)} via eq. ("," Once the explicit values of $S(\tau_{NL-1})$, $S(\tau_{NL})$, $S^{\prime\prime}(\tau_{NL-1})$, $S^{\prime\prime}(\tau_{NL})$ as well as those of the set $\lbrace I^{+}(\tau_{NL},\mu_{J}, J=1,ND)\rbrace$ are known, it is straightforward to compute those of the set $\lbrace I^{+}(\tau_{NL-1}),\mu_{J}, J=1,ND)\rbrace$ via eq. ("
+15).,15).
+" Then eqs (27) and (28) will vield the explicit values of S(yj,3) and S""(zy,2). hence those of the set (4.(7x49.40.J=1.ή."," Then eq.s (27) and (28) will yield the explicit values of $S(\tau_{NL-2})$ and $S^{\prime\prime}(\tau_{NL-2})$, hence those of the set $\lbrace I^{+}(\tau_{NL-2},\mu_{J}, J=1,ND)\rbrace$."
+ And so on along the back- substitution., And so on along the back- substitution.
+ Our huplicit Integral Method is based on (hie progressive (treatment of the different lavers iab consitute a model of the stellar atinosphere physical svstem. [rom the outermost laver the surface) to the deepest one (the bottom).," Our Implicit Integral Method is based on the progressive treatment of the different layers that consitute a model of the stellar atmosphere physical system, from the outermost layer (the surface) to the deepest one (the bottom)."
+ The protagonist variables of the method are je upgoimg and downgomg specific intensities /7(7.(i) as well as the corresponding source —unctions that besides (he thermal sources imelude a scattering-like integral into which there Mier (he foregoing specilic intensiües.," The protagonist variables of the method are the upgoing and downgoing specific intensities $I^{\pm}(\tau,\mu)$ as well as the corresponding source functions that besides the thermal sources include a scattering-like integral into which there enter the foregoing specific intensities."
+" Precisely. the study. (ancl the elimination) of each V.ingle laver (75.7,5,1) leads to a relation that links linearly the value οτι) of the source function ad 7; with οτι4). the value al τι4."," Precisely, the study (and the elimination) of each single layer $(\tau_{L},\tau_{L+1})$ leads to a relation that links linearly the value $S(\tau_{L})$ of the source function at $\tau_{L}$ with $S(\tau_{L+1})$, the value at $\tau_{L+1}$."
+" Once obtained via the study of the last laver 1e relation between (he values of the source function at the two last optical depth points. 1 boundary. condition al Ty, given by eq. ("," Once obtained via the study of the last layer the relation between the values of the source function at the two last optical depth points, the boundary condition at $\tau_{NL}$ given by eq. ("
+5) makes it possible to compute the explicit values of (τν) and οτι4). hence all the others.,"5) makes it possible to compute the explicit values of $S(\tau_{NL})$ and $S(\tau_{NL-1})$, hence all the others."
+ In order to design the required elimination scheme it is necessary to employ. a mathematical model for 9(7)., In order to design the required elimination scheme it is necessary to employ a mathematical model for $S(\tau)$.
+ In principle (he simplest ancl easiest model would be a piece-wise linear one. but the discontinuity of the first derivative ο(τ) at each knot τι can imply severe errors and possible numerical instabilities because the above discontinuity is incompatible with the radiative transfer (RT) process itself. where both /7(7.ji) and their first derivatives must be continuous. and therefore also the mean intensity J(7) and its first derivative.," In principle the simplest and easiest model would be a piece-wise linear one, but the discontinuity of the first derivative $S^{\prime}(\tau)$ at each knot $\tau_{L}$ can imply severe errors and possible numerical instabilities because the above discontinuity is incompatible with the radiative transfer (RT) process itself, where both $I^{\pm}(\tau,\mu)$ and their first derivatives must be continuous, and therefore also the mean intensity $J(\tau)$ and its first derivative."
+ Thus the foregoing model cannot be correct.but [or extreme casesof the thermal sources.," Thus the foregoing model cannot be correct,but for extreme casesof the thermal sources."
+ A piece-wise parabolic model warrants the continuity of 5'(7) at all depth points., A piece-wise parabolic model warrants the continuity of $S^{\prime}(\tau)$ at all depth points.
+ Such a model shall include also ο(τ) as a protagonist variable in the process of progressive, Such a model shall include also $S^{\prime}(\tau)$ as a protagonist variable in the process of progressive
+ $22..,\ref{subsec:progenitor}. .
+ ?? focuses on the correlations between the interacting components in an AGN pair., \ref{subsec:correlation} focuses on the correlations between the interacting components in an AGN pair.
+ We then examine in $2? the correlation between BH accretion and host recent star formation in AGN pairs and compare with those of control samples of single AGNs., We then examine in \ref{subsec:sfagn} the correlation between BH accretion and host recent star formation in AGN pairs and compare with those of control samples of single AGNs.
+" In this section, we address how galaxy tidal encounters affect recent star formation in AGN pairs."," In this section, we address how galaxy tidal encounters affect recent star formation in AGN pairs."
+" We cannot directly use the commonly used optical emission lines aand [Orr]) as indicators of current star formation for the galaxies in our sample because they contain contribution from AGN (e.g.,???).."," We cannot directly use the commonly used optical emission lines and ) as indicators of current star formation for the galaxies in our sample because they contain contribution from AGN \citep[e.g.,][]{ho97,kauffmann03,brinchmann04}."
+" Instead we adopt the continuum spectral indices aand aas indicators of recent star formation, following ?.."," Instead we adopt the continuum spectral indices and as indicators of recent star formation, following \citet{kauffmann03}."
+" Strong aabsorption lines arise in galaxies that experienced a burst of star formation that ended —0.1—1 Gyr ago, and large vvalues indicate strong recent starburst activity (e.g.,??).. D,,(4000)"," Strong absorption lines arise in galaxies that experienced a burst of star formation that ended $\sim$ 0.1–1 Gyr ago, and large values indicate strong recent starburst activity \citep[e.g.,][]{worthey97,kauffmann03c}."
+"mmeasures the continuum break around rest-frame 4000A, which arises from a series of metal lines typical of 1-2-Gyr old stars (e.g.,???).."," measures the continuum break around rest-frame 4000, which arises from a series of metal lines typical of 1–2-Gyr old stars \citep[e.g.,][]{bruzual83,balogh99,bc03}."
+" iis small for young stellar populations and large for old, rich galaxies, and can be used to characterize the mean stellar age (?).."," is small for young stellar populations and large for old, metal-rich galaxies, and can be used to characterize the luminosity-weighted mean stellar age \citep{kauffmann03c}."
+ Figure 3 shows aand aas functions of and Av for the pair and tidal samples., Figure \ref{fig:sf} shows and as functions of $r_p$ and $\Delta v$ for the pair and tidal samples.
+" Here and r,the error bars denote |-σ uncertainties in the medians."," Here and throughout, the error bars denote $\sigma$ uncertainties in the medians."
+" For throughout,AGN pairs with tidal features, both aand aare significantly correlated with γρ, despite having a large scatter."," For AGN pairs with tidal features, both and are significantly correlated with $r_p$, despite having a large scatter."
+" Spearman correlation tests show that the probability of null correlation P,,j; between ((D,(4000)) and τρ is 9x10° (7x1079)."," Spearman correlation tests show that the probability of null correlation $P_{null}$ between ) and $r_p$ is $9\times 10^{-5}$ $7\times
+10^{-6}$ )."
+ The correlation probabilities quoted are estimated from the individual data points directly., The correlation probabilities quoted are estimated from the individual data points directly.
+ Pairs with smaller projected on average exhibit stronger recent starburst activityseparations and younger mean stellar ages., Pairs with smaller projected separations on average exhibit stronger recent starburst activity and younger mean stellar ages.
+" Compared to the control AGN sample matched in redshift and stellar mass distributions, AGN pairs with tidal features on average have larger aand ((median values of 1.430.1 and 1.55—0.01 compared to 0.1+0.1 and 1.6740.01)."," Compared to the control AGN sample matched in redshift and stellar mass distributions, AGN pairs with tidal features on average have larger and (median values of $1.4\pm0.1$ and $1.55\pm0.01$ compared to $0.1\pm0.1$ and $1.67\pm0.01$ )."
+" The enhancement becomes significant (>3c above the control) for objects with r,<10- 30 kkpc (i.e., log(rp/h54 kkpc) S 1.0-1.5)."," The enhancement becomes significant $> 3\sigma$ above the control) for objects with $r_p \lesssim $ $30$ kpc (i.e., $(r_p/$ kpc) $\lesssim 1.0$ $1.5$ )."
+" In a sample of ~10? star-forming galaxies from SDSS DR4 (?),, ? find that aanti-correlates with specific star formation rate (SSFR) measured from optical emission lines."," In a sample of $\sim10^5$ star-forming galaxies from SDSS DR4 \citep{SDSSDR4}, , \citet{brinchmann04} find that anti-correlates with specific star formation rate (SSFR) measured from optical emission lines."
+" Objects with the smallest separation in our sample (r,~5 kkpc) have median ~1.40+0.04, corresponding to log(SFR/M,)/(yr!)~-10.1 according to the calibration of ?.."," Objects with the smallest separation in our sample $r_p\sim 5$ kpc) have median $\sim 1.40\pm0.04$, corresponding to $({\rm
+SFR}/M_{\ast})/({\rm yr}^{-1})\sim -10.1$ according to the calibration of \citet{brinchmann04}."
+" This is ~0.9+0.2 dex than the SSFRs inferred for the control sample (with median ~1.67+0.01, corresponding to log(SFR/M.)/(yr!) of —11.0)."," This is $\sim0.9\pm0.2$ dex than the SSFRs inferred for the control sample (with median $\sim 1.67\pm0.01$, corresponding to $({\rm
+SFR}/M_{\ast})/({\rm yr}^{-1})$ of $-11.0$ )."
+" We detect a similar dependence of ((D,(4000)) on r, in the parent population of all AGN pairs (i.e., regardless of the detection of tidal features)."," We detect a similar dependence of ) on $r_p$ in the parent population of all AGN pairs (i.e., regardless of the detection of tidal features)."
+" This suggests that the of ((D,,(4000)) is robust against r,-dependencebiases from the requirement of tidal features.", This suggests that the $r_p$ -dependence of ) is robust against biases from the requirement of tidal features.
+" However, the difference between the median and overall ((D,,(4000)) distributions of the pairand control samples is much smaller than that between thetidal and control samples,"," However, the difference between the median and overall ) distributions of the pairand control samples is much smaller than that between thetidal and control samples,"
+ while others abandon a universal initial mass [function ancl consider a full raciiative treatment of dust 2005).,", while others abandon a universal initial mass function and consider a full radiative treatment of dust ."
+. Several authors have also emphasised the importance of feedback due to powerful radio jets. which may. push back and heat the ionized gas. reducing or even stopping the cooling Hows building up the galaxy 2005).," Several authors have also emphasised the importance of feedback due to powerful radio jets, which may push back and heat the ionized gas, reducing or even stopping the cooling flows building up the galaxy ."
+. On the other hand. radiojets could also induce star formation 2004).. as observed in at least one high redshift radio galaxy 2000).," On the other hand, radio jets could also induce star formation , as observed in at least one high redshift radio galaxy ."
+. lt ds clear that. observations of a large number of high5 redshift 5ewlaxies. in particular the most massive ones. are essential in order to constrain the structure formation models described. above.," It is clear that observations of a large number of high redshift galaxies, in particular the most massive ones, are essential in order to constrain the structure formation models described above."
+" adio galaxies are known to be among the most massive galaxies known at cach recishift 2004)... ancl are ideal laboratories for Itocca-studying the ""radio feedback mechanism. described. above."," Radio galaxies are known to be among the most massive galaxies known at each redshift , and are ideal laboratories for studying the `radio feedback' mechanism described above."
+ Modern: racio surveys cover almost the entire sky. and. allow us to pinpoint the most extreme tthe most massive) galaxies. provided one can isolate them in radio catalogues containing up to 2 million. sources.," Modern radio surveys cover almost the entire sky, and allow us to pinpoint the most extreme the most massive) galaxies, provided one can isolate them in radio catalogues containing up to 2 million sources."
+ We have started such a search for distant. radio galaxies in the southern hemisphere. using the MMIIz Sydney University Molonglo Sky Survey and the NRAQ VLA Sky Survey 1998)., We have started such a search for distant radio galaxies in the southern hemisphere using the MHz Sydney University Molonglo Sky Survey and the NRAO VLA Sky Survey .
+.. In. the first. paper of this series 2004)... we define a sample of 76 high redshift radio galaxy candidates selected. on the ois of their ultra-steep radio spectra (USS: àx—— L3. Ssox EU) which has been almost the sole way to. find ⋅ ∆≻↓⋅⋯⇂⊲↓∪⋏∙≟⋜↧↓⋜∟∖⊲," In the first paper of this series , we define a sample of 76 high redshift radio galaxy candidates selected on the basis of their ultra-steep radio spectra (USS; $\alpha \leq -1.3$ , $S_{\nu} \propto \nu^{\alpha}$ ), which has been almost the sole way to find $z>3$ radio galaxies."
+⊓⊾⊳∖↿∖↙⊽∡↙↙↽∐⋖⇩≱⇂⋏∙≟⋖⋅↓⋰↓⊔⋏∙≟⋖⋅⇂⋜↧↥⋡↓≤⋗≤⋗⊤∶∐∢⋅≓≧↓⋅∢⋅⋯∼↳∢⋅⋜↧↓⊳⇉∪∪↓∶↼∣⋜⊔⋅∖⋰↓⊳∖∢⊾⇂⋜↧↓⊳⇉∪∪⊔↦∖∖⊽⋖⋅⋜↧↓⊳∖∪↓≻↓⋅∢⊾⊳∖⋖⋅⊔↿ ↕⊀↓⋏∙≟↓↥−↓⋅⋖⋅⊳∖∪⋯⊲↓∪⊔↓⋅⋯∐∪⊲↓⊔↓⋜↧⋏∙≟⋠↓⊔⋏∙≟↿⋖≱⋖⋟⋡↿⋜↧↕↓↕⋯∼≼⇍⊔↓⋅⋜⋯⋅↓≻∪≱∖⋠⊔⋠↓∪⊔≱∖ ⋜⋯∠⇂⊔↓∪↓⋅↓≻↓↕∪↓∪⋏∙≟⊀⊔⇍∥↓⊀↓⊔⇂⋅∪↓⋅⊔↓⋜∐⊀↓∪⊔⊳⋜⋯∠⇂⊔⋖⋅⋜⊔⋅−∐⊰↕∠⇂⋖⊾↓↥⇂↕↓↕≼∼⋜↧↿↕∢≱↓↕≻ of the host galaxies.," We also present high-resolution radio imaging to obtain accurate positions and morphological information, and near-IR identifications of the host galaxies."
+ In this paper H. we present the results o date of optical spectroscopy to determine their redshifts.," In this paper II, we present the results to date of optical spectroscopy to determine their redshifts."
+ Alulti-frequeney radio observations and ai discussion. of he physies of ultra-steep-spectrum radio galaxies will be oesented. in a future paper (Ixlamer. et. al., Multi-frequency radio observations and a discussion of the physics of ultra-steep-spectrum radio galaxies will be presented in a future paper (Klamer et al.
+ in prep.)., in prep.).
+" Throughout this paper. we adopt a Dat ACDAL cosmology with Ly—ΤΙ + MpeἘν On,=027 and Q420.73 2003).. We used four dillerent imaging spectrographs for the follow-up observations of the sources in our. USS sample."," Throughout this paper, we adopt a flat $\Lambda$ CDM cosmology with $_0=71$ $^{-1}$ $^{-1}$, $\Omega_{\rm M}=0.27$ and $\Omega_{\Lambda}=0.73$ We used four different imaging spectrographs for the follow-up observations of the sources in our USS sample."
+ For identilications detected on the cligitised sky surveys or in the 2 Micron All Sky Survey2003)... we used the Australian National University’s 2.3m telescope at Siding Spring Observatory. Coonabarabran. NSW with the Double Beam SpectrographLOSS).," For identifications detected on the digitised sky surveys or in the 2 Micron All Sky Survey, we used the Australian National University's 2.3m telescope at Siding Spring Observatory, Coonabarabran, NSW with the Double Beam Spectrograph."
+.. For the remaining sources with A19 identifications. we initially attempted to measure redshifts with the ESO Alulti-Mocde Instrument (EEMMI) on the New Technology Telescope (ΝΕ).," For the remaining sources with $K\simlt 19$ identifications, we initially attempted to measure redshifts with the ESO Multi-Mode Instrument (EMMI) on the New Technology Telescope (NTT)."
+ Hono τοΚΙ could be determined. they were re-observed with the ESO Very Large Telescope (VET). along with the faintest Ao19 identifications. using the FOcal Iteducer. ancl Spectrographs on two of the unit telescopes.," If no redshift could be determined, they were re-observed with the ESO Very Large Telescope (VLT), along with the faintest $K\simgt 19$ identifications, using the FOcal Reducer and Spectrographs on two of the unit telescopes."
+ We selected all our targets from the [ist of 76. USS sources defined in paper |. Our goal was to obtain redshift information for the entire sample. but due to adverse weather conditions (mainly at the NEL). we could. only observe 53 sources.," We selected all our targets from the list of 76 USS sources defined in paper I. Our goal was to obtain redshift information for the entire sample, but due to adverse weather conditions (mainly at the NTT), we could only observe 53 sources."
+ Within the available A range. we gave priority to the 53 sources in the sample with à<1.3.," Within the available RA range, we gave priority to the 53 sources in the sample with $\alpha<-1.3$."
+ In total. we have observed 41 out of 53 sources with a<18 ancl 12 out of 23 sources with a1.," In total, we have observed 41 out of 53 sources with $\alpha<-1.3$ and 12 out of 23 sources with $\alpha>-1.3$."
+ Our spectroscopic sample is thus complete for the most important subset of à<1.3 SOULCOS., Our spectroscopic sample is thus complete for the most important subset of $\alpha<-1.3$ sources.
+ Aefore attempting optical spectroscopy. of the LO faintest dy band sources. we first imaged. these fields in £ banel using FORS2 on the ΑΕ (sec Table. 1)).," Before attempting optical spectroscopy of the 10 faintest $K-$ band sources, we first imaged these fields in $I-$ band using FORS2 on the VLT (see Table \ref{imagingjournal}) )."
+ We split. the observations in exposures of typically 3muminutes. each. while dithering the frames by a lew arcseconds to ensure 16 object. did not fall on a bad. pixel on the detector.," We split the observations in exposures of typically minutes each, while dithering the frames by a few arcseconds to ensure the object did not fall on a bad pixel on the detector."
+ The pixel scale of FORS2 is 07225/pix. and the seeing during 1ο observations varied between 0755. and 1700.," The pixel scale of FORS2 is 25/pix, and the seeing during the observations varied between 5 and 0."
+ We used 10 standard. imaging reduction steps inALL consisting X bias subtraction. Batfielding. and registration of the ithered frames.," We used the standard imaging reduction steps in, consisting of bias subtraction, flatfielding, and registration of the dithered frames."
+ We fine-tuned the astrometry using all non- 2NLASS stars (tvpically 10 per field) present in 1e images. vielding a solution which we estimate to be accurate up to 20733. which is suflicient to identify the host galaxies of the radio sources.," We fine-tuned the astrometry using all non-saturated 2MASS stars (typically $\sim$ 10 per field) present in the images, yielding a solution which we estimate to be accurate up to $\sim$ 3, which is sufficient to identify the host galaxies of the radio sources."
+ Finally. we measured the magnitudes in -ddiametor apertures using theZA.," Finally, we measured the magnitudes in diameter apertures using the task."
+ taskphot. ‘Table 20 gives ajournal of our spectroscopic observations., Table \ref{spectroscopyjournal} gives ajournal of our spectroscopic observations.
+ Ifa candidate was observed at dillerent telescopes. we use only the best quality spectrum.," If a candidate was observed at different telescopes, we use only the best quality spectrum."
+Phe columns are:,The columns are:
+Assuming5 that à<0.3. an inaccuracy introduced by this simplification with respect to the Kerr solution (due to neglecting higher order terms in.) is smaller than LOW.,"Assuming that $\tilde{a}\le0.3$, an inaccuracy introduced by this simplification with respect to the Kerr solution (due to neglecting higher order terms in $j$ ) is smaller than $10\%$."
+ We put explicitly the spin dependence through @ on the left hand side for latter convenience., We put explicitly the spin dependence through $\tilde{a}$ on the left hand side for latter convenience.
+" Ànd therefore the relation between AM, and J, becomes approximately According to our parametric resonance model. for slowly rotating stars. the twin kIIz-QDPOs are given by vp=2p, and ο=31,. Where the superscript s stands forslow."," And therefore the relation between $M_*$ and $I_*$ becomes approximately According to our parametric resonance model, for slowly rotating stars, the twin kHz-QPOs are given by $\nu_1^{\rm s} = 2\,\nu_*$ and $\nu_2^{\rm s} = 3\,\nu_*$, where the superscript s stands for."
+ Because v7.v5 are interpreted as the frequencies of the orbital motion. thev need to be less than that at the ISCO For increasing spin of the neutron star 7. al some point. 75 will approach and eventually overtake Vice.," Because $\nu_1^{\rm s},\nu_2^{\rm s}$ are interpreted as the frequencies of the orbital motion, they need to be less than that at the ISCO For increasing spin of the neutron star $\nu_*$, at some point, $\nu_2^{\rm s}$ will approach and eventually overtake $\nu_{\rm
+ isco}$."
+ Thus £5 will be forbidden as a HF-QPO., Thus $\nu_2^{\rm s}$ will be forbidden as a HF-QPO.
+" As a consequence. the next (wo dominant twin kILz-QDPOs are identified as vf—1.55, and vb=2v,."," As a consequence, the next two dominant twin kHz-QPOs are identified as $\nu_1^{\rm f} = 1.5\,\nu_*$ and $\nu_2^{\rm f} = 2\,\nu_*$."
+" Therefore. the QPO frequency ciflerence Avνι=(νο—14)/v,. jumps suddenly from 1.0 and 0.5."," Therefore, the QPO frequency difference $\Delta\nu / \nu_* =
+(\nu_2-\nu_1) / \nu_*$ jumps suddenly from 1.0 and 0.5."
+" According to the data (aken Irom Méndez&Belloni(2007):vanderIxlis(2008).. this should happen in the neutron star spin range 7,€[863.401] Hz."," According to the data taken from \cite{2007MNRAS.381..790M, 2008AIPC.1068..163V}, this should happen in the neutron star spin range $\nu_*\in[363,401]$ Hz."
+ This is probably the most salient feature in the slow against [ast rotator discrepancies., This is probably the most salient feature in the slow against fast rotator discrepancies.
+ Filling these data requires that the switching from slow to last rotator occurs for neutron star spin between 363 Iz and 401 Iz., Fitting these data requires that the switching from slow to fast rotator occurs for neutron star spin between 363 Hz and 401 Hz.
+" More precisely. [or i4,<363 Iz. Ανν,zz1 which we interpret as no effect on motion in the observable disk from the presence of an ISCO."," More precisely, for $\nu_*\le363$ Hz, $\Delta\nu / \nu_* \approx1$ which we interpret as no effect on motion in the observable disk from the presence of an ISCO."
+ This implies that (οὗHz)>3142=1089 Hz. we put the spin rate into coma to disünguish between different rotators. an essential remark for our constrains.," This implies that $\nu_{\rm isco}(363\textrm{ Hz}) \ge 3 \,
+\nu_*=1089$ Hz, we put the spin rate into coma to distinguish between different rotators, an essential remark for our constrains."
+" Next. [or v,>401 Iz. Av/y,©0.5 which we interpret as a clear signature of the ISCO."," Next, for $\nu_* \ge 401$ Hz, $\Delta\nu / \nu_* \approx 0.5$ which we interpret as a clear signature of the ISCO."
+" This implies that ,.,(401Hz)x31x=1203 Hz."," This implies that $\nu_{\rm
+ isco}(401\textrm{ Hz}) \le 3\,\nu*=1203$ Hz."
+ Express in terms of the ISCO. (he transition from slow to fast rotator should happen when the two conditions below are," Express in terms of the ISCO, the transition from slow to fast rotator should happen when the two conditions below are"
+This inequality shows that the heating occurs for particles below a certain critical size.,This inequality shows that the heating occurs for particles below a certain critical size.
+" The mechanism of the stochastic acceleration is effective for collisions with particles of size =a, so that the relevant dust density to be substituted in Eq. (10))"," The mechanism of the stochastic acceleration is effective for collisions with particles of size $\gtrsim a$, so that the relevant dust density to be substituted in Eq. \ref{2}) )"
+" is determined by the integral ng(a)=f""(dna/da)da, where dng/da is the appropriate size distribution."," is determined by the integral $n_d(a)=\int_a^{\infty}(dn_d/da)da$, where $dn_d/da$ is the appropriate size distribution."
+ Usually dna/da decreases rather steeply and therefore the critical size should have a weak dependence on the parameters of the ISM., Usually $dn_d/da$ decreases rather steeply and therefore the critical size should have a weak dependence on the parameters of the ISM.
+" For the estimate, let us assume the MRN size distribution for dust (Mathis, Rumpl, Nordsieck 1977) valid for the size range between a few dozens of and a few tenths of uum: where Ayrn~10779 cm?® (Draine Lee 1984)."," For the estimate, let us assume the MRN size distribution for dust (Mathis, Rumpl, Nordsieck 1977) valid for the size range between a few dozens of and a few tenths of $\mu$ m: where $A_{\rm MRN}\sim10^{-25}$ $^{2.5}$ (Draine Lee 1984)."
+ This yields na(a)/ng0.3AMRNGa~2°., This yields $n_d(a)/n_{\rm H}\sim0.3A_{\rm MRN}a^{-2.5}$.
+" Taking m;10:34 g, nj,~1 cm?, T;~0.1 eV, and assuming also n, ng, we get the critical diameter of ~0.3 jum, i.e., the heating can already be triggered in the submicron range."," Taking $m_i\sim10^{-24}$ g, $n_i\sim1$ $^{-3}$ , $T_i\sim0.1$ eV, and assuming also $n_n\approx n_{\rm H}$ , we get the critical diameter of $\sim0.3~\mu$ m, i.e., the heating can already be triggered in the submicron range."
+ It is noteworthy that the critical size also has a weak dependence on the stochastic properties of the charge fluctuations., It is noteworthy that the critical size also has a weak dependence on the stochastic properties of the charge fluctuations.
+" For instance, from Eq. (9))"," For instance, from Eq. \ref{0}) )"
+ one can readily deduce that the scaling of the critical size on the charging frequency is ος. (, one can readily deduce that the scaling of the critical size on the charging frequency is $\propto\nu_{\rm ch}^{-0.2}$. (
+"ii) In order to vaodetermine the ultimate temperature T3 that can be reached, we recall that when Τα exceeds the critical value the source term in Eq. (9))","ii) In order to determine the ultimate temperature $T_d^{\infty}$ that can be reached, we recall that when $T_d$ exceeds the critical value $T_d^{\rm cr}$, the source term in Eq. \ref{0}) )"
+" is saturated at a constantT?', value.", is saturated at a constant value.
+" Hence, we have T$?~"," Hence, we have $T_d^{\infty}\sim 0.1(\lambda_{{\rm D}}/a)^2(n_d/n_n)T_d^{\rm cr}$ ."
+" By using the definition of z, the critical 0.1(Ap/a)?(na/n,)TS'.temperature can be rewritten as Tj""= "," By using the definition of $z$, the critical temperature can be rewritten as $T_d^{\rm cr}=z^2(T_i/e)^2a$ ."
+"Then we finally get the following estimate for the z?(T;/e))a.ultimate kinetic temperature of grains: Thus, the increase of the dust temperature relative to the temperature of gas depends very strongly on the grain size."," Then we finally get the following estimate for the ultimate kinetic temperature of grains: Thus, the increase of the dust temperature relative to the temperature of gas depends very strongly on the grain size."
+ Assuming again the MRN distribution for na(a) we get T9/T;οςα3δ., Assuming again the MRN distribution for $n_d(a)$ we get $T_d^{\infty}/T_i\propto a^{-3.5}$.
+" The corresponding RMS velocity of grains, dependencevf= scales as xa~*?°, whereas the on the \/kpTS/ma,plasma parameters is weaker, οςT?/n;1/3."," The corresponding RMS velocity of grains, $v_{T_d}^{\infty}=\sqrt{k_{\rm B}T_d^{\infty}/m_d}$, scales as $\propto a^{-3.25}$, whereas the dependence on the plasma parameters is weaker, $\propto T_i^{3/2}n_i^{-1/2}$."
+" For the parameters used above, a 30 nm dust particle can reach the kinetic temperature which is 5-6 orders of magnitude higher than T;. ("," For the parameters used above, a $30$ nm dust particle can reach the kinetic temperature which is 5–6 orders of magnitude higher than $T_i$. ("
+"iii) The heating timescale Του (ie., when T4? is is determined by the stage of linear temperature growth.","iii) The heating timescale $\tau_{\infty}$ (i.e., when $T_d^{\infty}$ is reached) is determined by the stage of linear temperature growth."
+" reached)This can be straightforwardly obtained from Eq. (9)),"," This can be straightforwardly obtained from Eq. \ref{0}) ),"
+" The heating timescale is independent of the plasma density and has rather weak dependence on other parameters, Tx.οςanj!T,1/2'/."," The heating timescale is independent of the plasma density and has rather weak dependence on other parameters, $\tau_{\infty}\propto an_n^{-1}T_i^{-1/2}$."
+ This estimate yields heating timescale of several Myr for 30 nm grains., This estimate yields heating timescale of several Myr for $30$ nm grains.
+" 'The formalism above is rather general and is applicable to any astrophysical media, including accretion disks, circumstellar regions and the ISM."," The formalism above is rather general and is applicable to any astrophysical media, including accretion disks, circumstellar regions and the ISM."
+ In the next section we focus our attention on the ISM dust acceleration., In the next section we focus our attention on the ISM dust acceleration.
+ The ISM dust acceleration is extremely important process for shattering and coagulation of dust (see Hirashita Yan 2009)., The ISM dust acceleration is extremely important process for shattering and coagulation of dust (see Hirashita Yan 2009).
+ We discuss both subsonic and supersonic motion of dust., We discuss both subsonic and supersonic motion of dust.
+ In Eq. (D6)) (, In Eq. \ref{vd_app}) ) (
+Appendix D) we rewrite Eq. (12)),Appendix D) we rewrite Eq. \ref{2a}) )
+" in terms of the RMS velocity of accelerated grains assuming the MRN dust size distribution, Eq. (11)),"," in terms of the RMS velocity of accelerated grains assuming the MRN dust size distribution, Eq. \ref{MRN}) ),"
+ where where n; is in cm? and the renormalizing factor(13) a for the damping rate in the subsonic and supersonic regimes is given by Eq. (2)), where where $n_i$ is in $^{-3}$ and the renormalizing factor $\alpha$ for the damping rate in the subsonic and supersonic regimes is given by Eq. \ref{alpha1}) )
+" and (3)), respectively."," and \ref{alpha2}) ), respectively."
+ Note that in the latter case a is the function of dust velocity and hence Eq. (13)), Note that in the latter case $\alpha$ is the function of dust velocity and hence Eq. \ref{vd}) )
+ should be resolved for v7*., should be resolved for $v_{T_d}^{\infty}$.
+ In Figure we compare the mechanism presented in the paper 2.with the calculations of the velocities of interstellar grains obtained in YLD04., In Figure \ref{f2} we compare the mechanism presented in the paper with the calculations of the velocities of interstellar grains obtained in YLD04.
+ We see that for sufficiently large grains the acceleration due to MHD fluctuations is stronger., We see that for sufficiently large grains the acceleration due to MHD fluctuations is stronger.
+" At the same time, the new mechanism of acceleration associated with charge fluctuations is more efficient for a<107? cm."," At the same time, the new mechanism of acceleration associated with charge fluctuations is more efficient for $a\lesssim10^{-5}$ cm."
+ The efficiency of the charge-fluctuation-induced acceleration increases for smaller grains., The efficiency of the charge-fluctuation-induced acceleration increases for smaller grains.
+" However, the limitation of our treatment above is that the grain charging was considered as a continuous process [see Eq. (5))]."," However, the limitation of our treatment above is that the grain charging was considered as a continuous process [see Eq. \ref{24}) )]."
+" Formally, this approach is not applicable to the situation when the mean grain charge Qo is so small that each act of losing or acquiring an elementary charge results in substantial discrete fluctuations."," Formally, this approach is not applicable to the situation when the mean grain charge $Q_0$ is so small that each act of losing or acquiring an elementary charge results in substantial discrete fluctuations."
+" However, we believe that even in this case we may get right order of magnitude estimate of grain acceleration."," However, we believe that even in this case we may get right order of magnitude estimate of grain acceleration."
+ Further research with Monte-Carlo simulations instead of using Fokker-Plank approach should test the latter conjecture., Further research with Monte-Carlo simulations instead of using Fokker-Plank approach should test the latter conjecture.
+Comparing the Figures 1 and 2 we see the limited nature of our quantitative calculations.,Comparing the Figures \ref{f1} and \ref{f2} we see the limited nature of our quantitative calculations.
+" For grains larger than ~3x107"" cm the interactions decrease due to the less efficient momentum deposition as grains mostly feel the electric field of grains smaller than their size.", For grains larger than $\sim3\times 10^{-7}$ cm the interactions decrease due to the less efficient momentum deposition as grains mostly feel the electric field of grains smaller than their size.
+Because of diffusion. particles are not restricted to the clownstream part of the flow: there will be a df40 also upstream. so that we need to consider appropriate conditions to mateh of in the two regions.,"Because of diffusion, particles are not restricted to the downstream part of the flow: there will be a $\delta\!f \neq 0$ also upstream, so that we need to consider appropriate conditions to match $\delta\!f$ in the two regions."
+ The first condition is obviously the continuity across the shock. It is also well-known that the spatial gradient of 6f needs (to satisIv a boundary condition al (he shock: this is derived by integrating eq.," The first condition is obviously the continuity across the shock, It is also well-known that the spatial gradient of $\delta\! f$ needs to satisfy a boundary condition at the shock: this is derived by integrating eq."
+ 4 on an infinitesimal interval straddling the shock., \ref{boltzmann} on an infinitesimal interval straddling the shock.
+" The unit vector normal to the surface of the flapping shock is n=(1.75,0) where ¢ is (he shock corrugation amplitude."," The unit vector normal to the surface of the flapping shock is $\hat{n}=(1,\imath k_y \zeta)$ where $\zeta$ is the shock corrugation amplitude."
+ So. we obtain: WritingE its first-order linearization and using5 eq.," So, we obtain: Writing its first-order linearization and using eq."
+" 44 we find: We note that (0fορ)=(Of/Oy),0 because at zero-th order the medium is uniform in coordinates parallel tothe shock surface."," \ref{aux19} we find: We note that ${(\partial f/\partial y)}_2={(\partial f/\partial
+y)}_1=0$ because at zero-th order the medium is uniform in coordinates parallel tothe shock surface."
+ The result Eqs., The result Eqs.
+ 4 and 47 are (he appropriate boundary. conditions for our problem., \ref{aux19} and \ref{aux20} are the appropriate boundary conditions for our problem.
+ We now show how to satisfy the boundary conditions. eqs.," We now show how to satisfy the boundary conditions, eqs."
+ 44 and 47 al the shock.," \ref{aux19}
+ and \ref{aux20} at the shock."
+ Using the notation introduced in Sect. ??.. ," Using the notation introduced in Sect. \ref{subsec:newnotation}, ,"
+"we know that df on the downstream side of the shock (the [factor εἰs"" will be omitted for simplicity in this subsection) satislies Please notice (hat both downstream and upstream (he sununalion is over 2 mocles (a pressure wave and a third mode).", we know that $\delta\! f_+$ on the downstream side of the shock (the factor $e^{\imath\omega t-\imath k_y y}$ will be omitted for simplicity in this subsection) satisfies Please notice that both downstream and upstream the summation is over 2 modes (a pressure wave and a third mode).
+ In fact. we know on the one hand (hat particles perturbations are not coupled to entropy. ancl vorticity perturbations. while. on the other hand. we know that pressure and third modes. for givenvalues of απο we. exist Lor opposite," In fact, we know on the one hand that particles perturbations are not coupled to entropy and vorticity perturbations, while, on the other hand, we know that pressure and third modes, for givenvalues of $k_y$ and $\omega$ , exist for opposite"
+map to be scale-dependent by performing the analysis iu harmonic space. the assumption beime that this allows any spatial variations in the spectral dependence of the foregrounds to be adequately tracked.,"map to be scale-dependent by performing the analysis in harmonic space, the assumption being that this allows any spatial variations in the spectral dependence of the foregrounds to be adequately tracked."
+" Tt is uot clear to what extent real variations project outo the harmonic eiecuinodes of the analysis,", It is not clear to what extent real variations project onto the harmonic eigenmodes of the analysis.
+ As with the ILC inethod. conrplex noise properties result. aud so it is uulikely that this iiethod is suitable for high precision cosmological analyses.," As with the ILC method, complex noise properties result, and so it is unlikely that this method is suitable for high precision cosmological analyses."
+ In what follows we will denote the map as TCM the Teemarl et ccleaned map., In what follows we will denote the map as TCM – the Tegmark et cleaned map.
+ Iu this paper a new look is taken at the ILC method presented by Bennettetal.(2003b).. with the main goal of determining whether a map derived iu this manner can be suitable for cosmological purposes.," In this paper a new look is taken at the ILC method presented by \citet{bennett:2003b}, with the main goal of determining whether a map derived in this manner can be suitable for cosmological purposes."
+ Specifically. we derive a new ILC imap based on Lagrauge multipliers Gu what follows to be referred to as the LILC — Lagerauge Tuternal Linear Combination map). which has lower variance than the scieuce team’s ILC (hereafter referred to as WILC) map.," Specifically, we derive a new ILC map based on Lagrange multipliers (in what follows to be referred to as the LILC – Lagrange Internal Linear Combination – map), which has lower variance than the science team's ILC (hereafter referred to as WILC) map."
+ We then generate Alonte Carlo simulations of this map by adding white noise and foreground templates to CMD realizations. aud process these through our ILC pipeline.," We then generate Monte Carlo simulations of this map by adding white noise and foreground templates to CMB realizations, and process these through our ILC pipeline."
+ This allows us to quantify the efficicney of the ILC method. aud realistic foreground residual estimates may be established.," This allows us to quantify the efficiency of the ILC method, and realistic foreground residual estimates may be established."
+ Iu the final section we repeat the large-scale analysis of deOliveira-Costaetal.(2001) both for our new LILC imap aud for the simulations. to assess the impact of residual forceroundse ou these statistics.," In the final section we repeat the large-scale analysis of \citet{de
+Oliveira-Costa:2004} both for our new LILC map and for the simulations, to assess the impact of residual foregrounds on these statistics."
+ However. we study not only the quadrupole aud the octopole. but also consider the properties of the (=1.5 and 6 modes.," However, we study not only the quadrupole and the octopole, but also consider the properties of the $\ell=4, 5$ and 6 modes."
+ Iu fact. we fud that the properties of the latter two are at least as intriguing as those of the quadrupole aud octopole: the (6=5 mode is highly spherically svuuuctric. and the (6=6 mode is planar.," In fact, we find that the properties of the latter two are at least as intriguing as those of the quadrupole and octopole: the $\ell=5$ mode is highly spherically symmetric, and the $\ell=6$ mode is planar."
+ The ILC method as defined by Bennettetal.(20035). is based on a smiple premise: suppose there are & observed CXMB naps at different frequencies (but with identical beams). aud the aimi is to suppress foregrounds aud uoise as far as possible.," The ILC method as defined by \citet{bennett:2003b} is based on a simple premise: suppose there are $k$ observed CMB maps at different frequencies (but with identical beams), and the aim is to suppress foregrounds and noise as far as possible."
+ Each of the & maps may be written Gu theriodyaiunic temperature) on the form T(r)=Town πμ... statistically independent.," Each of the $k$ maps may be written (in thermodynamic temperature) on the form $T(\nu_k) = T_{\textrm{CMB}} +
+T_{\textrm{residual}}(\nu_k)$ , where $T_{\textrm{CMB}}$ and $T_{\textrm{residual}}(\nu_k)$ are statistically independent."
+ Therefore. if oue now forms the linear combination and requires that the resulting map may be written as Thus. the response to the CAIB signal i always uuity since if is mdependeut of the frequency. aud. the klL free weights iav be chosen to minimize the iurpact of the residuals.," Therefore, if one now forms the linear combination and requires that the resulting map may be written as Thus, the response to the CMB signal is always unity since it is independent of the frequency, and the $k-1$ free weights may be chosen to minimize the impact of the residuals."
+ Asstumine the CMD componeut is statistically independent of the foregrounds aud the noise. one convenient measure for this is simply the variance of T. The internal linear. combination method may now be defined succinctly in terms of Equations 1. aud 2.. where he weights are deteriumed by mnünimuüzius the varince in Equation L..," Assuming the CMB component is statistically independent of the foregrounds and the noise, one convenient measure for this is simply the variance of $T$, The internal linear combination method may now be defined succinctly in terms of Equations \ref{eq:lin_comb} and \ref{eq:lin_const}, where the weights are determined by minimizing the variance in Equation \ref{eq:minimization}."
+ We compute the ILC weights by means of Lagrauge nultiplicrs., We compute the ILC weights by means of Lagrange multipliers.
+ Our Lagrange multiplier procedure is similar o the approach taken by Tegimarketal.(2003)/ for computing the harmonic space weights frou which their nap is constructed., Our Lagrange multiplier procedure is similar to the approach taken by \citet{tegmark:2003} for computing the harmonic space weights from which their map is constructed.
+ A useful review of this method is also given by Teemark(1998)., A useful review of this method is also given by \citet{tegmark:1998}.
+.. The variance of T ids seen to be a quadratic form in the weights w;. and its winiuization uuder the coustraimt eiven iu Equation 2 is therefore most couveuieutlv carried out by memus of Lagrange multiplicrs.," The variance of $T$ is seen to be a quadratic form in the weights $w_i$, and its minimization under the constraint given in Equation \ref{eq:lin_const} is therefore most conveniently carried out by means of Lagrange multipliers."
+ As shown in Appendix A the linear system of equations to be solved can be written on the followine form where A is an arbitrary constant. w=(ay.....wpe are the ILC weights. aud is the mmap-to-map covariance matrix.," As shown in Appendix A the linear system of equations to be solved can be written on the following form where $\lambda$ is an arbitrary constant, $\mathbf{w} = (w_{1},
+\ldots, w_{k})^\mathrm{T}$ are the ILC weights, and is the map-to-map covariance matrix."
+ The solutious to this system are the usual inverse covariance weights. If the foreground properties varv stronglv over the sky as a result of spatially dependent spectral indexes. then the ILC method mav perform rather poorly.," The solutions to this system are the usual inverse covariance weights, If the foreground properties vary strongly over the sky as a result of spatially dependent spectral indexes, then the ILC method may perform rather poorly."
+ To remedy this. oue may subdivide the sky iuto disjoiut patches. aud compute indepeudent set of weights for cach patch.," To remedy this, one may subdivide the sky into disjoint patches, and compute independent set of weights for each patch."
+ Bennettetal.(2003b) cavided the full ska iuf twelve regions. clever covering the nou-muiform regious of Galactic plane. aud the last one covering the kp? region plus the well-behaved parts of the Galactic plane.," \citet{bennett:2003b} divided the full sky into twelve regions, eleven covering the non-uniform regions of Galactic plane, and the last one covering the Kp2 region plus the well-behaved parts of the Galactic plane."
+ We will study this particular partitioning more closely in 83.., We will study this particular partitioning more closely in \ref{sec:simulations}.
+ Using such a partitioning. the minimization of the variance in Equation 1. ds carried out for each region separately. and the final step is therefore to coustruct oue single full-xkv map frou those individual patches.," Using such a partitioning, the minimization of the variance in Equation \ref{eq:minimization} is carried out for each region separately, and the final step is therefore to construct one single full-sky map from those individual patches."
+ Iu order to suppress boundary effects Bennettetal.(20035) eenerated a mask (ic. a fullskv map cousistiug of 0x and Ls) for cach patch. and couvolved these masks by a Caussian beam of 907 FWHAL," In order to suppress boundary effects \citet{bennett:2003b} generated a mask (i.e., a full-sky map consisting of 0's and 1's) for each patch, and convolved these masks by a Gaussian beam of $90\arcmin$ FWHM."
+ This final ILC map was then coustructed by first eeuecratiug one full-skv map from each ILC weight set. as described above. aud theu they co-added these maps pixebbw-pixel with weights eiven by the apodized masks.," This final ILC map was then constructed by first generating one full-sky map from each ILC weight set, as described above, and then they co-added these maps pixel-by-pixel with weights given by the apodized masks."
+ We adopt the same method for suppressing boundary effects without modifications., We adopt the same method for suppressing boundary effects without modifications.
+The study of gravitational lenses has become a powerful tool to address several distinet cosmological and astrophysical questions.,The study of gravitational lenses has become a powerful tool to address several distinct cosmological and astrophysical questions.
+ These include the distribution of dark matter in galaxies (Keeton et citekeeton*98)). studying dust extinction. at redshift ς>>0 (Nadeau et citenadeau*9]:; Jean Surde] 1998)). and determining the Hubble parameter Hy from measurements of light travel time delay between different lines of sight (Refsdal 1964)).," These include the distribution of dark matter in galaxies (Keeton et \\cite{keeton*98}) ), studying dust extinction at redshift $z\gg 0$ (Nadeau et \\cite{nadeau*91}; Jean Surdej \cite{jean*98}) ), and determining the Hubble parameter $H_0$ from measurements of light travel time delay between different lines of sight (Refsdal \cite{refsdal64}) )."
+ The appeal of this method lies in its complete independence of the traditional cosmic distance ladder. vielding distances for redshift objects in a single leap.," The appeal of this method lies in its complete independence of the traditional cosmic distance ladder, yielding distances for high-redshift objects in a single leap."
+ The uncertainties of Hy estimation are mainly limited by ambiguities in modelling the deflector mass distribution. since all relevant errors can. be reduced to insignificance.," The uncertainties of $H_0$ estimation are mainly limited by ambiguities in modelling the deflector mass distribution, since all relevant errors can be reduced to insignificance."
+" While ""simple? lenses such as double QSOs do not strongly constrain the deflector model. multiple systems such as quadruply imaged QSOs provide many additional constraints. allowing tto independently test the assumed mass distribution models SSaha Williams 1997)."," While `simple' lenses such as double QSOs do not strongly constrain the deflector model, multiple systems such as quadruply imaged QSOs provide many additional constraints, allowing to independently test the assumed mass distribution models Saha Williams \cite{saha*97}) )."
+" A problem with quadruple systems. however. can be excessive symmetry leading to time delays which are short compared to intrinsic radio and optical variability timescales. às m the famous ""Einstein. Cross’ | 030 or in the ""Clover Leaf H 13131 117."," A problem with quadruple systems, however, can be excessive symmetry leading to time delays which are short compared to intrinsic radio and optical variability timescales, as in the famous `Einstein Cross' $+$ 030 or in the `Clover Leaf' H $+$ 117."
+ In this paper we present the discovery of the new multiple QSO HE 2130., In this paper we present the discovery of the new multiple QSO HE $-$ 2130.
+ The object was originally identified as à high-probability QSO candidate in the course of the Hamburg/ESO survey (Wisotzki et citewisotzki*96)) and subsequently observed as part of a large imaging search for lensed QSOs., The object was originally identified as a high-probability QSO candidate in the course of the Hamburg/ESO survey (Wisotzki et \\cite{wisotzki*96}) ) and subsequently observed as part of a large imaging search for lensed QSOs.
+" The coordinates of this object are Κ.Α. = 02 32"" 337]. Dec 2. 21° ΙΤ’ 26"" (2000.0). as measured in theSurvev."," The coordinates of this object are R.A. = $^\mathrm{h}$ $^\mathrm{m}$ $\fs 1$, Dec = $-21\degr$ $'$ $''$ (J2000.0), as measured in the."
+. Imaging data were obtained at the CTIO mm on 13 November 1998. inJohnson B and R. and Cousins 7.," Imaging data were obtained at the CTIO m on 13 November 1998, inJohnson $B$ and $R$, and Cousins $I$."
+ In each band. three images were taken of mmin exposure time each.," In each band, three images were taken of min exposure time each."
+" The seeing was somewhat variable: 1"" in B. 171 in R. and 078 in J."," The seeing was somewhat variable: $1''$ in $B$, $1\farcs 1$ in $R$, and $0\farcs 8$ in $I$."
+ Pixel size was 0724., Pixel size was $0\farcs 24$.
+ Small sections of these images are reproduced in reffig:brik.. revealing the multiple structure of the QSO already at first glance.," Small sections of these images are reproduced in \\ref{fig:brik}, revealing the multiple structure of the QSO already at first glance."
+ Two almost merging images (Al and Α2: see reffig:pos)) of nearly equal brightness dominate the total magnitudes., Two almost merging images (A1 and A2; see \\ref{fig:pos}) ) of nearly equal brightness dominate the total magnitudes.
+ Another discrete component. B. is clearly apparent with roughly similar colours.," Another discrete component, B, is clearly apparent with roughly similar colours."
+ Components C and D are much less prominent. and in fact appear as two distinet sources rather than just one only in /.," Components C and D are much less prominent, and in fact appear as two distinct sources rather than just one only in $I$."
+" On the same night. near-infrared (K, band. um) images were taken at the CTIO 4mm telescope with the CIRIM imager."," On the same night, near-infrared $K_s$ band, $\mu$ m) images were taken at the CTIO m telescope with the CIRIM imager."
+ Total exposure time was mmin under 079 seeing., Total exposure time was min under $0\farcs 9$ seeing.
+" The pixel size of the raw data was 0!""42. but the single exposures were combined on a 2« finer grid: the resulting image is also shown in reffig:brik.."," The pixel size of the raw data was $0\farcs 42$ , but the single exposures were combined on a $2\times$ finer grid; the resulting image is also shown in \\ref{fig:brik}. ."
+ The centrally located component D 1s now almost, The centrally located component D is now almost
+standard irradiated models (see e.g. ?;; ?? and references therein).,"standard irradiated models (see e.g. \citealt{Gui08}; ; \citealt{CBL08, BCB09} and references therein)."
+" So far, two objects in the CoRoT sample belong to this category: CoRoT-Exo-1b and -2b."," So far, two objects in the CoRoT sample belong to this category: CoRoT-Exo-1b and -2b."
+ Comparison between models and observations for these two planets are shown in Fig., Comparison between models and observations for these two planets are shown in Fig.
+" 2 and 3,, respectively."," \ref{fig:c1} and \ref{fig:c2}, respectively."
+ The effect of irradiation alone does not allow to reproduce the observed radii of these strongly inflated objects., The effect of irradiation alone does not allow to reproduce the observed radii of these strongly inflated objects.
+" These observations thus confirm the existence of a missing mechanism in the models, which slows down the cooling and the contraction of the planet."," These observations thus confirm the existence of a missing mechanism in the models, which slows down the cooling and the contraction of the planet."
+ Several possibilities have been suggested to explain this puzzling property., Several possibilities have been suggested to explain this puzzling property.
+" Tidal heating due to circularization of the orbit, as originally suggested (7?) and then rejected on the basis of a too short characteristic timescale compared with the age of the systems, might provide or at least participate to the lacking mechanism in some cases, provided tidal effects in the planet and in thestar are properly taken into account (Jackson et al."," Tidal heating due to circularization of the orbit, as originally suggested \citep{BLM01} and then rejected on the basis of a too short characteristic timescale compared with the age of the systems, might provide or at least participate to the lacking mechanism in some cases, provided tidal effects in the planet and in the are properly taken into account (Jackson et al."
+ 2008) (??)..," 2008) \citep{LWC09,IB09}. ."
+" In a recent paper, however, ? show that most transiting systems are not in a state of tidal equilibrium and thus that tidal circularization timescale estimates based on equilibrium tides are not correct."," In a recent paper, however, \cite{LWC09}
+ show that most transiting systems are not in a state of tidal equilibrium and thus that tidal circularization timescale estimates based on equilibrium tides are not correct."
+" À major consequence of ? calculations is that assuming zero-eccentricity, when not directly determined, in transiting light curve analysis, is not necessarily correct and such analysis should be redone carefully."," A major consequence of \cite{LWC09} calculations is that assuming zero-eccentricity, when not directly determined, in transiting light curve analysis, is not necessarily correct and such analysis should be redone carefully."
+ Tidal heating might thus in some cases provide an extra source of energy during a significant fraction of the planet’s evolution., Tidal heating might thus in some cases provide an extra source of energy during a significant fraction of the planet's evolution.
+" In the present models, we account for this source of energy, according to ?.."," In the present models, we account for this source of energy, according to \cite{Hut81}."
+" Note that, since zero eccentricity is assumed for the analysis of all the CoRoT transit light-curves, we assume a small but finite value for e in order to have a quantitative estimate of the effect of tidal heating on the bloated CoRoT planets."," Note that, since zero eccentricity is assumed for the analysis of all the CoRoT transit light-curves, we assume a small but finite value for $e$ in order to have a quantitative estimate of the effect of tidal heating on the bloated CoRoT planets."
+" In the case of CoRoT-Exo-1b, an eccentricity e—0.02 provides enough tidal heating to reproduce the observed radius within the error bars (dash-dotted curve in Fig. 2))."," In the case of CoRoT-Exo-1b, an eccentricity $e=0.02$ provides enough tidal heating to reproduce the observed radius within the error bars (dash-dotted curve in Fig. \ref{fig:c1}) )."
+" For CoRoT-Exo-2b, an eccentricity of only a few percents does not provide enough tidal heating and larger values are required to reproduce the radius."," For CoRoT-Exo-2b, an eccentricity of only a few percents does not provide enough tidal heating and larger values are required to reproduce the radius."
+ We find that a value e=0.15 yields a radius in agreement with observations., We find that a value $e=0.15$ yields a radius in agreement with observations.
+" ? suggested another heating mechanism in the deep interior of planets, originating from the strong winds generated at the planet's surface."," \cite{SG02} suggested another heating mechanism in the deep interior of planets, originating from the strong winds generated at the planet's surface."
+" Their numerical simulations of atmospheric circulation produce a downward kinetic energy flux of about of the absorbed stellar incident flux, which dissipates in the interior and slows down the planet’s contraction (?).."," Their numerical simulations of atmospheric circulation produce a downward kinetic energy flux of about of the absorbed stellar incident flux, which dissipates in the interior and slows down the planet's contraction \citep{GS02}."
+" Although the validity of this scenario is still debated, with various simulations producing different results (see e.g. ?)), it is worth exploring this issue."," Although the validity of this scenario is still debated, with various simulations producing different results (see e.g. \citealt{SMC08}) ), it is worth exploring this issue."
+" To test this effect, we have included an extra source of energy corresponding to of the stellar impinging flux contribution in our evolutionary models, as done in ?.."," To test this effect, we have included an extra source of energy corresponding to of the stellar impinging flux contribution in our evolutionary models, as done in \cite{CBB04}."
+" In the case of CoRoT-Exo-1b (dashed curve in Fig. 2)),"," In the case of CoRoT-Exo-1b (dashed curve in Fig. \ref{fig:c1}) ),"
+" the resulting theoretical radius is too large, by ~8%,, compared to the observed value, while for CoRoT-Exo-2b, the same relative contribution yields too small of a radius."," the resulting theoretical radius is too large, by $\sim$, compared to the observed value, while for CoRoT-Exo-2b, the same relative contribution yields too small of a radius."
+" Note that for CoRoT-Exo-4b, analysed in §??,, models including of the stellar impinging flux as an extra source of energy also overestimate by the radius compared to the observed value, except if the heavy element content of the planet is significantly larger than the inferred ~5% mass fraction."," Note that for CoRoT-Exo-4b, analysed in \ref{sec:sp}, models including of the stellar impinging flux as an extra source of energy also overestimate by the radius compared to the observed value, except if the heavy element content of the planet is significantly larger than the inferred $\sim 5\%$ mass fraction."
+" Enhanced atmospheric opacity in transiting planets (?),, although not excluded, remains so far too much of an ad-hoc suggestion to be examined in detail (see ? for a discussion)."," Enhanced atmospheric opacity in transiting planets \citep{BHB07}, although not excluded, remains so far too much of an ad-hoc suggestion to be examined in detail (see \citealt{BCB09} for a discussion)."
+" At any rate, the presently detected most inflated transiting planets, like Tres-4b and WASP-12, can not be explained even with such enhanced-opacity models (?7).."," At any rate, the presently detected most inflated transiting planets, like Tres-4b and WASP-12, can not be explained even with such enhanced-opacity models \citep{Gui08,BCB09}."
+" An other mechanism, based on (inefficient) layered or oscillatory convection in some planet interiors,has been suggestedby ? and has been shown to possibly explain the"," An other mechanism, based on (inefficient) layered or oscillatory convection in some planet interiors,has been suggestedby \cite{CB07} and has been shown to possibly explain the"
+radius of 20 km) now evaporating CO around ο) nuust be where foe is the molecular weielt. Meo=7«107% ke is the total mass of CO iji the disk (AppeneixA:)). and το=2-10Ms tis the photodissociation rate of CO (Van Dishoeck Black. 1988) destraved by the UV interstellar background (the stellar extreme ultraviolet flux is very low and uceleib|.,"radius of 20 km) now evaporating CO around $\beta\:$ must be where $\mu_{CO}$ is the molecular weight, $M_{CO}=7\times 10^{20}$ kg is the total mass of CO in the disk \ref{appendix}) ), and $\tau_{CO}=2\cdot 10^{-10}$ $^{-1}$ is the photodissociation rate of CO (Van Dishoeck Black, 1988) destroyed by the UV interstellar background (the stellar extreme ultraviolet flux is very low and negligible)."
+ This ununber is extremely arec but unavoidable because CO isobserved., This number is extremely large but unavoidable because CO is.
+ Tt can he larger if the evayoration rate or the typical radii of the bodies are smaller., It can be larger if the evaporation rate or the typical radii of the bodies are smaller.
+ Eveu with the lower limit ou the mass of CO eiven in Appendix As. we still have to accept that the nuuber of bodies now evaporating around ο) aaud producing the observed CO iust be in the order of teus of millions.," Even with the lower limit on the mass of CO given in \ref{appendix}, we still have to accept that the number of bodies now evaporating around $\beta\:$ and producing the observed CO must be in the order of tens of millions."
+ We can also roughly evaluate Vorp. the umber of bodies necessary to produce the outerdust disk. if this disk is produced by evaporation of CO like in the OEB scenario.," We can also roughly evaluate $N_{OEB}$, the number of bodies necessary to produce the outer disk, if this disk is produced by evaporation of CO like in the OEB scenario."
+ This wmmber of OED is where A; is the mass of the dust disk. ty as the dust life-time. yg is the mass ratio of the dust effectively kept in the disk to the dust produced. and ομε 1s the mass ratio of dust to CO.," This number of OEB is where $M_d$ is the mass of the dust disk, $t_d$ is the dust life-time, $\varphi$ is the mass ratio of the dust effectively kept in the disk to the dust produced, and $\phi_{dust/CO}$ is the mass ratio of dust to CO."
+" We assume that the disk is about one hun mass (AL,~7105 κο),", We assume that the disk is about one lunar mass $M_d\sim 7\cdot 10^{22}$ kg).
+ The dust life tine (ty 15r) is taken from Artyinowicz (1997)!5., The dust life time $t_d=10^4$ yr) is taken from Artymowicz (1997).
+. One can evaluate οον0.1., One can evaluate $\varphi\sim 0.1$.
+" ouq445:/0 IS Very uncertain: we use the receut value eiven bv Sekanina (1996) ou the comet ILde-Bopp observed at more than 6 AU from the Sun. aud we assume that 054,555"")210."," $\phi_{dust/CO}$ is very uncertain; we use the recent value given by Sekanina (1996) on the comet Hale-Bopp observed at more than 6 AU from the Sun, and we assume that $\phi_{dust/CO}\ga 10$."
+ Finally we obtain: We see that. within the uucertainties. this number has the same order of 1uaguitudoe as the απνο derived above in Eq. 1..," Finally we obtain: We see that, within the uncertainties, this number has the same order of magnitude as the number derived above in Eq. \ref{Nco}."
+ As in Eq., As in Eq.
+ 12 of LVF. this shows that the mass of dust driven by CO around ο) iis consistent with the total mass of dust in the outer part of the disk.," 12 of LVF, this shows that the mass of dust driven by CO around $\beta\:$ is consistent with the total mass of dust in the outer part of the disk."
+ These ~10* 105 objecs unnist be compared to the 10? 10° objects believed to 10 present between 30 aud 100 AU from the Sun as the sotree of the Jupiter Family Comets (Duncan Levison 199T. Morbidelli 1997).," These $\sim 10^7$ $10^8$ objects must be compared to the $10^8$ $10^9$ objects believed to be present between 30 and 100 AU from the Sun as the source of the Jupiter Family Comets (Duncan Levison 1997, Morbidelli 1997)."
+ If we cousider that these bodies have a radius of about 20 km as eiveu by the mocels preseuted in LVF or giveu by the models of Sect. 5..," If we consider that these bodies have a radius of about 20 km as given by the models presented in LVF or given by the models of Sect. \ref{migrating planet},"
+ this corresponds to about an Earth mass., this corresponds to about an Earth mass.
+ This is well below the mass needed to supply the 9+ daisk ouly bv collision (30 Earth mass is needed according to Backinan et al., This is well below the mass needed to supply the $\beta\:$ disk only by collision (30 Earth mass is needed according to Backman et al.
+ 1995)., 1995).
+ Evaporation requires less mass of parent bodies. provided that some process is able to start the evaporation.," Evaporation requires less mass of parent bodies, provided that some process is able to start the evaporation."
+ As we will sce in Sect 5.. a plienomenonu believed to have occurred in the voung Solar System cau explain this evaporation.," As we will see in Sect \ref{migrating planet}, a phenomenon believed to have occurred in the young Solar System can explain this evaporation."
+ It could be difficult to imagine that ~107? bodies have always been active for —105 wears., It could be difficult to imagine that $\sim 10^8$ bodies have always been active for $\sim 10^8$ years.
+ Oulv cousideriug he observed CO sives that diving LOS vears. Meo«]UjvenmwsTeo——20Mggaqa of CO must have been evaporated!," Only considering the observed CO gives that during $10^8$ years, $M_{CO}\times 10^8 {\rm years}\times \tau_{CO} = 20 M_{\rm Earth}$ of CO must have been evaporated!"
+ It seems unlikely that this large nuniber of yodies have been active since the birth of the svstem., It seems unlikely that this large number of bodies have been active since the birth of the system.
+ This eives evidence that either that 9 dis very voung (but already: on the main sequence. which gives a lower nuit on the age o DE10 years (Crifo et al.," This gives evidence that either that $\beta\:$ is very young (but already on the main sequence, which gives a lower limit on the age of $\ga 10^7$ years (Crifo et al.,"
+ 1997. see also discussion in Vidal-Macdjar et al..," 1997, see also discussion in Vidal-Madjar et al.,"
+ 1998) or that j ust be a transient pleno11οion., 1998) or that $\beta\:$ must be a transient phenomenon.
+ There is in fact no reason to believe hat the ) system was always as disty as observed today., There is in fact no reason to believe that the $\beta\:$ system was always as dusty as observed today.
+ Of COULSC. the idea that this disk is uot trausicut is a consequence of auv inodoel of collisional crosion fro asteroid to dust.," Of course, the idea that this disk is not transient is a consequence of any model of collisional erosion from asteroid to dust."
+ But with other scenarios. we can easilv inuaegime that a particular phenomenon occurred recently. aud that the density of the j ddisk must be sSenificautlv ualcr during the quiescent pliase of simple collisional erosion duriug which the density cau be similar to the characteristic deusitv of the more common Veea--like stars.," But with other scenarios, we can easily imagine that a particular phenomenon occurred recently, and that the density of the $\beta\:$ disk must be significantly smaller during the quiescent phase of simple collisional erosion during which the density can be similar to the characteristic density of the more common -like stars."
+ Since the OEDs are able to explain some inportaut aspects of the observations. possible origin of these OEBs. or more exactly the perturbations necessary to explain thei evaporation. have to be explored.," Since the OEBs are able to explain some important aspects of the observations, possible origin of these OEBs, or more exactly the perturbations necessary to explain their evaporation, have to be explored."
+ Lhnuudeed. evaporation takes place only when a body is formed bevond a vaporization Πιτ of a volatile iux| is then relocated inside this luüt.," Indeed, evaporation takes place only when a body is formed beyond a vaporization limit of a volatile and is then relocated inside this limit."
+" Iu fact. the ""relocalon can have different nieaniug for an object on a keplerian orbit."," In fact, the “relocation” can have different meaning for an object on a keplerian orbit."
+ Tt cai be a simple variation of ifs senuianajor axis or a- Increase of its eccentricity., It can be a simple variation of its semi-major axis or an increase of its eccentricity.
+ In any case. the nost 1nportaut Is CCLtainly the periastron where most of the evaporatio- can take place.," In any case, the most important is certainly the periastron where most of the evaporation can take place."
+ In fact. the beeimuiug of evaporatio-," In fact, the beginning of evaporation"
+WDs. however the field population of metal rich stars ancl single low mass WDs imply that low mass stars with ΡοΗ > +0.3 are likely to lose enough mass on the red giant branch to skip the asvimptotic eiant. branch evolution and directly evolve into single helium-core WDs.,"WDs, however the field population of metal rich stars and single low mass WDs imply that low mass stars with [Fe/H] $\geq$ +0.3 are likely to lose enough mass on the red giant branch to skip the asymptotic giant branch evolution and directly evolve into single helium-core WDs."
+ The confixination of the mass loss - metallicity dependence for both NGC 6791 and the field stars has important implications: Even though several groups have been searching for planets around WDs through infrared (Debes Sigurdsson 2002: Burleigh et al., The confirmation of the mass loss - metallicity dependence for both NGC 6791 and the field stars has important implications: Even though several groups have been searching for planets around WDs through infrared (Debes Sigurdsson 2002; Burleigh et al.
+ 2002) and mid-inlrared excess around WDs (Aluially et al., 2002) and mid-infrared excess around WDs (Mullally et al.
+ 2007). there are no known planetary companions to WDs vet.," 2007), there are no known planetary companions to WDs yet."
+ This nay be eaused by a target selection effect., This may be caused by a target selection effect.
+ As all of these groups have targeted tvpical ~0.6 WWDs. (μον have selectively targeted descendants of lower metallicitw stars.," As all of these groups have targeted typical $\sim$ 0.6 WDs, they have selectively targeted descendants of lower metallicity stars."
+ Fischer Valenti (2005) found that at [Fe/II] > +0.3. of the stars have gas giant 1planets whereas less than of the stars with -0.5 < [Fe/H] < 0.0 have Doppler-detected planets.," Fischer Valenti (2005) found that at [Fe/H] $>$ +0.3, of the stars have gas giant planets whereas less than of the stars with -0.5 $<$ [Fe/H] $<$ 0.0 have Doppler-detected planets."
+ Since (he super-solar metallicity stars are expected to be the progenitors of single low mass WDs. these WDs are prime targets for planet detection.," Since the super-solar metallicity stars are expected to be the progenitors of single low mass WDs, these WDs are prime targets for planet detection."
+ If the progenitors of these single low mass WDs have planets likely). the planets are more likely to survive (he post-main sequence evolution. as these stars do not go through the ACB phase.," If the progenitors of these single low mass WDs have planets likely), the planets are more likely to survive the post-main sequence evolution, as these stars do not go through the AGB phase."
+ Even though the fraction of binary low mass WDs seems to be around [or ~0.4 WWDs. we expect this fraction to be [or the extremely low mass WDs with A/<0.2AL.," Even though the fraction of binary low mass WDs seems to be around for $\sim$ 0.4 WDs, we expect this fraction to be for the extremely low mass WDs with $M\leq0.2$."
+.. The mass loss process is most likely not significant enough to produce 0.2 WWDs even at. higher metallicities., The mass loss process is most likely not significant enough to produce 0.2 WDs even at higher metallicities.
+ The inilial-final mass relation for. WDs (Weidemann 2000: Ixalirai οἱ al., The initial-final mass relation for WDs (Weidemann 2000; Kalirai et al.
+ 20075) predict that a 1 sslar would produce a 0.55 WWD., 2007b) predict that a 1 star would produce a 0.55 WD.
+ On the other hand. at |Fe/II] ~ +04. D'Cruz et al. (," On the other hand, at [Fe/H] $\sim$ +0.4, D'Cruz et al. ("
+1996) predict that the same star would produce a 0.44—0.47 WWD. well above the Iimit.,"1996) predict that the same star would produce a $-$ 0.47 WD, well above the limit."
+ So [ar most of the known extremely low mass WDs are found as companions (to pulsars., So far most of the known extremely low mass WDs are found as companions to pulsars.
+ Ον recently. several 0.2 WWDs are lound in the field population (Liebert οἱ al.," Only recently, several 0.2 WDs are found in the field population (Liebert et al."
+ 2004: Ixasvka οἱ al., 2004; Kawka et al.
+ 2006: Eisenstein, 2006; Eisenstein
+receive the Nobel Prize in Physics in 1970 for fundamental work aud discoveries in MIID with fruitful applicatious in different parts of plaua physics.,receive the Nobel Prize in Physics in 1970 for fundamental work and discoveries in MHD with fruitful applications in different parts of plasma physics.
+ The notion of frozen-in maguetic field is the result of his work iu counectiou with the discovery of ΑΠΟ waves (Aven1912)., The notion of frozen-in magnetic field is the result of his work in connection with the discovery of MHD waves \citep{b1}.
+ The central point of MIID theory ds that conductive fiuids can support the magnetic field., The central point of MHD theory is that conductive fluids can support the magnetic field.
+ The preseuce of maguctic fields leads to forces that in turn act on the fluid. typically aplasia. thereby potentially alteius the eeometry and streneth of the magnetic fields themselves.," The presence of magnetic fields leads to forces that in turn act on the fluid, typically a plasma, thereby potentially altering the geometry and strength of the magnetic fields themselves."
+ Based ou the frame of MIID. a lot of theories come out. including ATID turbulence. ATID waves. uagnueto-couvection. MID reconnection. aud ΑΠΟ dyuiuao.," Based on the frame of MHD, a lot of theories come out, including MHD turbulence, MHD waves, magneto-convection, MHD reconnection, and MHD dynamo."
+ Iu fact. the ANID theory mace little scieutific WOLTCSS Hi space Science aik astroplivsics in the vast decades.," In fact, the MHD theory made little scientific progress in space science and astrophysics in the past decades."
+ Problems are still in puzzle., Problems are still in puzzle.
+ The sev problem of ΑΠΟ theory comes from the ecucralized Olii's law. which is believed to eive he electric current m cosmic plasima.," The key problem of MHD theory comes from the generalized Ohm's law, which is believed to give the electric current in cosmic plasma."
+ Iu the standard nou-relativistic form. the MIID equations consist of the basic conservation Laws of lass. nmnioinentuni and enerev together with the induction equation for the maeuetic field.," In the standard non-relativistic form, the MHD equations consist of the basic conservation laws of mass, momentum and energy together with the induction equation for the magnetic field."
+ The equations are. written im SI units: The equation of motion: Where pis the mass density aud u the fluid bulls velocity. pis the eas pressure. B the magnetic field. j the current density. aud σ is the viscous stress tensor.," The equations are, written in SI units: The equation of motion: Where $\rho$ is the mass density and $\bf{u}$ the fluid bulk velocity, $p$ is the gas pressure, $\bf{B}$ the magnetic field, $\bf{j}$ the current density, and $\sigma$ is the viscous stress tensor."
+The equation for the internal energy. which is usually written as an equation for the pressure p Where Q comprises the effects of heating aud cooling as well as thermal conduction aud 5 is the adiabaticitv coefficieut.,"The equation for the internal energy, which is usually written as an equation for the pressure $p$: Where $Q$ comprises the effects of heating and cooling as well as thermal conduction and $\gamma$ is the adiabaticity coefficient."
+ The above equation nuplies the equation of state of the ideal ionized eas: It is satisfied for most dilute plasinas., The above equation implies the equation of state of the ideal ionized gas: It is satisfied for most dilute plasmas.
+ The induction equation. orFaradavs law: Which is derived by inserting Oli's law: Iere. g ds the electrical resistivity.," The induction equation, orFaraday's law: Which is derived by inserting Ohm's law: Here, $\eta$ is the electrical resistivity."
+ In total. the MIID equation consist of two vector aud two scalar partial differential equation (or eight scalar equation) that are to be solved sinmltaucously.," In total, the MHD equation consist of two vector and two scalar partial differential equation (or eight scalar equation) that are to be solved simultaneously."
+ An carly theoretical paper by Lighthill(1960) onu magnetized plasma xoperties iu the MIID description contaius m its plivsies sections eritieisns about the applicability of ideal or resistive MTL— heory for plasmas., An early theoretical paper by \citet{b2} on magnetized plasma properties in the MHD description contains in its physics sections criticisms about the applicability of ideal or resistive MHD theory for plasmas.
+" Specifically, iguoriug the Πα ern in the generalized Ον law concerned. a shupli&cation still made almost routinely iu naguctic fusion theory."," Specifically, ignoring the Hall term in the generalized Ohm's law concerned, a simplification still made almost routinely in magnetic fusion theory."
+ The objections to classical AMID theory aud their consequeuces give rise to he development of Tall MIID aud its application o laboratory aud cosmic plasma (Witalis1986)., The objections to classical MHD theory and their consequences give rise to the development of Hall MHD and its application to laboratory and cosmic plasma \citep{b3}.
+. According to he MIID equation. we cau find hat the key poiut is the decision of the gcucralized Olhuu's law.," According to the MHD equation, we can find that the key point is the decision of the generalized Ohm's law."
+ Witalis(1986). stressed that to retain he Tall tei. using a two-fluid plasma description Is necessary.," \citet{b3} stressed that to retain the Hall term, using a two-fluid plasma description is necessary."
+ Then he generalized Oluu's law for tally ionized plasima can be derived., Then the generalized Ohm's law for fully ionized plasma can be derived.
+ For the macroscopic behavior of plasima. Spitzer(1962) eave the detailed discussion.," For the macroscopic behavior of plasma, \citet{b4} gave the detailed discussion."
+" The basic equations are the equations of ious and electrons together with the Maxwoell's equations. the equation of continuity, and in the condition of 4= 0. the eeneralized Oluu’s law is expressed as: Where."," The basic equations are the equations of ions and electrons together with the Maxwell's equations, the equation of continuity, and in the condition of $\frac{\partial}{\partial t}=0$ , the generalized Ohm's law is expressed as: Where,"
+The discovery of extrasolar planets around solar analogs. initiated by ?.. has radically modified. our understanding of planetary formation.,"The discovery of extrasolar planets around solar analogs, initiated by \citet{Mayor.1995}, has radically modified our understanding of planetary formation."
+ The new models that strive to describe planetary formation mechanisms need observational constraints on planetary companions on a large range of parameters. such as stellar mass. metallicity. planet mass and orbital parameters.," The new models that strive to describe planetary formation mechanisms need observational constraints on planetary companions on a large range of parameters, such as stellar mass, metallicity, planet mass and orbital parameters."
+ Notably. the most prolific methods for exoplanet detection. radial velocities and transits. only cover separations up to a few AU and barely reaches the large separation range where all the Solar System reside.," Notably, the most prolific methods for exoplanet detection, radial velocities and transits, only cover separations up to a few AU and barely reaches the large separation range where all the Solar System reside."
+ This range can be probed by direct resolution imaging., This range can be probed by direct high-resolution imaging.
+ Only a few unambiguously planetary- companions to stars (2???) were discovered by direct imaging. but each one of these detections proved extremely valuable to constrain formation models.," Only a few unambiguously planetary-mass companions to stars \citep{Chauvin.2004,MaroisHR8799.2008,Kalas.2008,Lagrange.2010} were discovered by direct imaging, but each one of these detections proved extremely valuable to constrain formation models."
+ Until recently. core accretion (CA) was the preferred model to explain both our Solar System planets and the close-in giant planets found by radial velocity (RV) around solar-type stars.," Until recently, core accretion (CA) was the preferred model to explain both our Solar System planets and the close-in giant planets found by radial velocity (RV) around solar-type stars."
+ The recently imaged at large separations from their parent A-type stars (??) challenge however this simple view as they are very difficult to form by CA. though ?. shows that under specific condition CA can form Jupiter-mass planets as far as 40-SOAU.," The recently imaged at large separations from their parent A-type stars \citep{MaroisHR8799.2008,Kalas.2008}
+ challenge however this simple view as they are very difficult to form by CA, though \citet{Rafikov.2011} shows that under specific condition CA can form Jupiter-mass planets as far as 40-50AU."
+ In the case of HR8799. this could account for the formation of the 3 innermost planets. but ? claim that HR8799bed formation by CA has to happen at closer separation and be followed by planet scattering at the observed separation. which would not create a stable system.," In the case of HR8799, this could account for the formation of the 3 innermost planets, but \citet{Dodson-Robinson.2009} claim that HR8799bcd formation by CA has to happen at closer separation and be followed by planet scattering at the observed separation, which would not create a stable system."
+ ? claim that. instead. gravitational instability (GI) in a dise could produce the bed planets while ? argue that HR8799bed can be formed by GI. but only at larger separation and show that under certain conditions they can migrate to the separations where they are currently observed.," \citet{Dodson-Robinson.2009} claim that, instead, gravitational instability (GI) in a disc could produce the $bcd$ planets while \citet{Kratter.2010} argue that HR8799bcd can be formed by GI, but only at larger separation and show that under certain conditions they can migrate to the separations where they are currently observed."
+ However the innermost planet of this system. HR8799e (?) appears at least as difficult to form by GI as the outermost planet HR8799b would be to form by CA. hinting that planetary formation models still need improvement to be able to account for the first multi-exoplanetary system directly imaged and badly need further observational constraints. especially in different mass range. before drawing any strong conclusion on the respective part they actually play in planetary formation.," However the innermost planet of this system, HR8799e \citep{Marois.2010HR8799e} appears at least as difficult to form by GI as the outermost planet HR8799b would be to form by CA, hinting that planetary formation models still need improvement to be able to account for the first multi-exoplanetary system directly imaged and badly need further observational constraints, especially in different mass range, before drawing any strong conclusion on the respective part they actually play in planetary formation."
+ Planetary mass objects could also form in the same way as stars. with ? showing that a few objects as light as ~5 can form by core fragmentation. possibly ending in very high mass-ratio binary system. very look-alike to a planetary system.," Planetary mass objects could also form in the same way as stars, with \citet{Bate.2009} showing that a few objects as light as $\sim$ can form by core fragmentation, possibly ending in very high mass-ratio binary system, very look-alike to a planetary system."
+ Such planetary-like binary system could also be formed by disc instabilities in à very massive primordial disc (2).. which could produce systems quite similar to 2M1207AB (?).. together with numerous brown dwarfs and very low-mass stars.," Such planetary-like binary system could also be formed by disc instabilities in a very massive primordial disc \citep{Stamatellos.2011}, which could produce systems quite similar to 2M1207AB \citep{Chauvin.2004}, together with numerous brown dwarfs and very low-mass stars."
+ The dependency of planetary formation modes to the stellar mass is an open field of investigation. with some observational clues that planetary masses. as well as disc masses. scale with stellar masses (e.g.??)..," The dependency of planetary formation modes to the stellar mass is an open field of investigation, with some observational clues that planetary masses, as well as disc masses, scale with stellar masses \citep[e.g.][]{Forveille.2011,Scholz.2006}."
+ Debris and transitional dises tracing planetary formation have been detected even around M dwarts and brown dwarfs. with frequencies similar (?) or even larger (?) than around solar-type stars.," Debris and transitional discs tracing planetary formation have been detected even around M dwarfs and brown dwarfs, with frequencies similar \citep{Plavchan.2009} or even larger \citep{Currie.2011} than around solar-type stars."
+ ? identifies some indications of bi-modal formation. but many more detections are needed to confirm this scenario.," \citet[][]{Boley.2009} identifies some indications of bi-modal formation, but many more detections are needed to confirm this scenario."
+ In that frame. looking for gas giant planets around M-type stars offers a distinct opportunity to test planet formation mechanisms. as models predict that CA is very inefficient ΚΚΟδ)..," In that frame, looking for gas giant planets around M-type stars offers a distinct opportunity to test planet formation mechanisms, as models predict that CA is very inefficient \citep[][KK08]{Kennedy.2008}."
+" KKO8 predict a rate of 1% for CA formed around M stars if the rate 1s of 6% for G type stars) or ""all but impossible"" (DRO9) in producing around those stars."," KK08 predict a rate of $\%$ for CA formed around M stars if the rate is of $\%$ for G type stars) or “all but impossible"" (DR09) in producing around those stars."
+ The main reason being that there is not enough time to form cores with I0Ma before the gas in the protoplanetary disk is dissipated., The main reason being that there is not enough time to form cores with $_{\oplus}$ before the gas in the protoplanetary disk is dissipated.
+" On the contrary. GI can be quite efficient in forming around M stars. at separations larger than typically AAU with ? claiming that ""even late M dwarf stars might be able to form gas giants on wide orbits”."," On the contrary, GI can be quite efficient in forming around M stars, at separations larger than typically AU with \citet{Boss.2011} claiming that “even late M dwarf stars might be able to form gas giants on wide orbits""."
+ This hints at the possibility that there are more at large separations around late-type stars than there are close-in ones., This hints at the possibility that there are more at large separations around late-type stars than there are close-in ones.
+"Fig. 4,,","Fig. \ref{Fig:Obscure},"
+ we plot the deviation from the nominal correlation with X-ray luminosity (as given in Figs., we plot the deviation from the nominal correlation with X-ray luminosity (as given in Figs.
+ 1. 2)) of each line against the depth of the silicate feature., \ref{Fig:LineComb} \ref{Fig:NeII}) ) of each line against the depth of the silicate feature.
+" Circinus shows the strongest deviations from all correlations and, at the same time, has the deepest silicate feature."," Circinus shows the strongest deviations from all correlations and, at the same time, has the deepest silicate feature."
+ It is deeper than in any other type 2 AGN in our sample., It is deeper than in any other type 2 AGN in our sample.
+" In case of foreground extinction (host or our galaxy), we expect that the flux within the lines is reduced in the same way without a change in W;."," In case of foreground extinction (host or our galaxy), we expect that the flux within the lines is reduced in the same way without a change in $\EW$."
+" On the other hand, if part or all of the line emission is originating from inward of the torus, the extinction in the lines could be much stronger in a more or less edge-on torus geometry."," On the other hand, if part or all of the line emission is originating from inward of the torus, the extinction in the lines could be much stronger in a more or less edge-on torus geometry."
+" Given the fact that the bulk of line emission in all other AGN is coming from the unresolved point source, such a scale and inclination effect appears possible."," Given the fact that the bulk of line emission in all other AGN is coming from the unresolved point source, such a scale and inclination effect appears possible."
+ Strong extinction towards the NLR based on mid-infrared spectra of Circinus was also noted by Rocheetal. (2006)., Strong extinction towards the NLR based on mid-infrared spectra of Circinus was also noted by \citet{Roc06}.
+". The negative correlations in the [Arm], iv], and [Neu] lines as observed with VLT/VISIR is much stronger than any known Baldwin effect in optical/UV broad or narrow lines."," The negative correlations in the ], ], and ] lines as observed with VLT/VISIR is much stronger than any known Baldwin effect in optical/UV broad or narrow lines."
+ This is remarkable since the spatial resolution in the optical of the most recent studies should be similar or even better than what is achieved in the mid-infrared with the VLT., This is remarkable since the spatial resolution in the optical of the most recent studies should be similar or even better than what is achieved in the mid-infrared with the VLT.
+" In UV/optical lines, a trend for larger correlation slopes with increasing ionization potential is observed (e.g.Dietrich 2002)."," In UV/optical lines, a trend for larger correlation slopes with increasing ionization potential is observed \citep[e.g.][]{Die02}."
+". The potential of the presented lines is in the range of eeV and, thus, comparable to some of the optical/UV lines."," The potential of the presented lines is in the range of eV and, thus, comparable to some of the optical/UV lines."
+" From the optical/UV lines, slopes of the order of —0.1 to —0.2 would be expected."," From the optical/UV lines, slopes of the order of $-0.1$ to $-0.2$ would be expected."
+ The measured slope of the mid-infrared lines deviates by at least 4-σ from that expectation (except i] w/o NGC 253: 2-σ)., The measured slope of the mid-infrared lines deviates by at least $\sigma$ from that expectation (except ] w/o NGC 253: $\sigma$ ).
+" In Sect ??,, we found some negative trend of W;([Arm]) with /gqa, but possibly not with Mgu."," In Sect \ref{Sec:MBH}, we found some negative trend of $\EW$ ]) with $\ledd$, but possibly not with $\MBH$."
+" If confirmed, this is in contrast to the observed and theoretically-expected behavior of the broad lines (e.g.Wandel1999;Warneretal.2004)."," If confirmed, this is in contrast to the observed and theoretically-expected behavior of the broad lines \citep[e.g.][]{Wan99,War04}."
+". Thus, the unsettled correlation with Mgy might imply that the iNLBE has a different physical origin than the BLBE, in particular questioning the relevance of the Mgy-dependence of the ionising continuum on the mid-infrared narrow-line emission."," Thus, the unsettled correlation with $\MBH$ might imply that the iNLBE has a different physical origin than the BLBE, in particular questioning the relevance of the $\MBH$ -dependence of the ionising continuum on the mid-infrared narrow-line emission."
+" One suggestion for the origin of the oNLBE relates to the size scaling of the narrow-line region (NLR), for which a luminosity dependence RygοςL9? has been found (Bennertetal.2002)."," One suggestion for the origin of the oNLBE relates to the size scaling of the narrow-line region (NLR), for which a luminosity dependence $\RNLR \propto L^{0.5}$ has been found \citep{Ben02}."
+". As a consequence, the NLR of luminous sources may reach galactic size scales and, thus, lose its gas."," As a consequence, the NLR of luminous sources may reach galactic size scales and, thus, lose its gas."
+ This scenario is often referred to as the “disappearing NLR” (Croometal.2002;Net-zeretal. 2006).," This scenario is often referred to as the “disappearing NLR” \citep{Cro02,Net06}."
+". For the mid-infrared lines, we find that the main portion of their emission is not coming from a spatially-extended region, e.g. as for πτ] in the optical, but from scales smaller than ppc."," For the mid-infrared lines, we find that the main portion of their emission is not coming from a spatially-extended region, e.g. as for ] in the optical, but from scales smaller than pc."
+" If the lack of line emission in Circinus is really an inclination-dependent obscuration effect, then the involved scales are significantly smaller."," If the lack of line emission in Circinus is really an inclination-dependent obscuration effect, then the involved scales are significantly smaller."
+ Since it is difficult, Since it is difficult
+The polarimetric study of weakly lensed galaxies can in principle provide information about the lensing gravitational potential jw otherwise would be extremely dillicult or impossible to obtain.,The polarimetric study of weakly lensed galaxies can in principle provide information about the lensing gravitational potential that otherwise would be extremely difficult or impossible to obtain.
+ The method relies on the property that the plane of »olarization of light is not alfected by the gravitational potential., The method relies on the property that the plane of polarization of light is not affected by the gravitational potential.
+ Phus the direction of polarization provides a marker for the xientation of the source galaxy. and 1ence information about the distortion in the image induced by the lens.," Thus the direction of polarization provides a marker for the orientation of the source galaxy, and hence information about the distortion in the image induced by the lens."
+ This reduces 10 uncertainty in the estimation of t1e lens parameters. and can even. in certain situations when a few source galaxies can vctually be measured polarimetσαν. break the cegeneraey of the equations hat determine the lens parameters.," This reduces the uncertainty in the estimation of the lens parameters, and can even, in certain situations when a few source galaxies can actually be measured polarimetrically, break the degeneracy of the equations that determine the lens parameters."
+ Alhough to a large extent we have assumed tiat Thomson or Ravleigh scattering is responsible for the polarization Lux. 10 our conclusions are not greatly afected by the presence of such small scde mechanisms as dichroism provided that on 1e large scale axial svmunetry holds.," Although to a large extent we have assumed that Thomson or Rayleigh scattering is responsible for the polarization flux, the our conclusions are not greatly affected by the presence of such small scale mechanisms as dichroism provided that on the large scale axial symmetry holds."
+ Only in exceptional circumstances woud one expect the direction of the polarization =rot to be perpendicular to the major axis of the eliptical isophotes., Only in exceptional circumstances would one expect the direction of the polarization not to be perpendicular to the major axis of the elliptical isophotes.
+ Since weak lensing would induce a rotation of much less mn 90 degrees. even if the integratec polarization were aligned parallel to the| major axis of the source. this should not Lead o confusion and hence measurement of the direcion of polarization would sill provide rigorous constraints on the lensing )nranmeters.," Since weak lensing would induce a rotation of much less than 90 degrees, even if the integrated polarization were aligned parallel to the major axis of the source, this should not lead to confusion and hence measurement of the direction of polarization would still provide rigorous constraints on the lensing parameters."
+ Aleasurement of the optical polarization of faint. galaxies would. require large telescopes and integration times. but it is important to realise that the main advantage of the proposed method relies on the use of the direction of polarization. ancl not the degree of polarization and the former is considerably. easier to establish.," Measurement of the optical polarization of faint galaxies would require large telescopes and integration times, but it is important to realise that the main advantage of the proposed method relies on the use of the direction of polarization, and not the degree of polarization and the former is considerably easier to establish."
+Power-lawlike X-ray spectra extending up to soft 5-ravys (hereafter Xs) are a common feature of the emission from accreting black holes in active galactic nuclei (Ας) and rvlack-hole binarics (BLIBs).,Power-law–like X-ray spectra extending up to soft $\gamma$ -rays (hereafter $\gamma$ ) are a common feature of the emission from accreting black holes in active galactic nuclei (AGNs) and black-hole binaries (BHBs).
+ ]t is generally accepted. that lOS Svectra are. produced by the Comptonization process aking place in optically thin accreting plasmas with tvpical electron. energies of ~100 keV. Similar spectra are also observed. from weakly-magnetized. accreting neutron. stars Barret οἱ 22000). where they can originate either in the accretion Dow or in a boundary [aver near the neutron star surface.," It is generally accepted that those spectra are produced by the Comptonization process taking place in optically thin accreting plasmas with typical electron energies of $\sim 100$ keV. Similar spectra are also observed from weakly-magnetized accreting neutron stars Barret et 2000), where they can originate either in the accretion flow or in a boundary layer near the neutron star surface."
+ The X5 spectra of Sevfert Ls and X-ray binaries in their hard. (low) states can be mocelled by Comptonization on thermal electrons. see ZZelziarski (1999).," The $\gamma$ spectra of Seyfert 1s and X-ray binaries in their hard (low) states can be modelled by Comptonization on thermal electrons, see Zdziarski (1999)."
+ However. both observations and. theoretical considerations suggest that a small non-thermal component may be present in the electron distribution. giving rise to a weak non-thermal high-energv tail in the observed Comptonization spectrum.," However, both observations and theoretical considerations suggest that a small non-thermal component may be present in the electron distribution, giving rise to a weak non-thermal high-energy tail in the observed Comptonization spectrum."
+ Such a tail above 1 MeV has been detected in the spectrum of (νο Χα in the hard state by the COMPTEL detector aboarel (MeConnell et 11994. 20002).," Such a tail above $\sim 1$ MeV has been detected in the spectrum of Cyg X-1 in the hard state by the COMPTEL detector aboard (McConnell et 1994, 2000a)."
+ In other BIIDs ancl in Sevíferts. the existing upper limits are compatible with the presence of weak non-thermal tails ZZdziarski et 11998: Gonedek et al.," In other BHBs and in Seyferts, the existing upper limits are compatible with the presence of weak non-thermal tails Zdziarski et 1998; Gondek et al."
+ 1996: Johnson et al., 1996; Johnson et al.
+ 1997)., 1997).
+ A non-thermal tail in the electron. distribution in compact N sources can arise due to many cillerent physical mechanisms., A non-thermal tail in the electron distribution in compact $\gamma$ sources can arise due to many different physical mechanisms.
+ One such mechanism is acceleration in the process of dissipation of magnetic field. (similarly to the case of the solar corona) in an optically-thin accretion Low or in active coronal regions above an, One such mechanism is acceleration in the process of dissipation of magnetic field (similarly to the case of the solar corona) in an optically-thin accretion flow or in active coronal regions above an
+photometry in the L-band. where the objects were often rather faint anc (ο) from uncertainties in the telluric extinction corrections at £ ancl to a lesser extent at JJ. (,"photometry in the $L$ -band, where the objects were often rather faint and (c) from uncertainties in the telluric extinction corrections at $L$ and to a lesser extent at $J$ . ("
+Lhe ff anc dv extinctions of the earth's atmosphere track each other very well. so that the JfA colour is the most reliable: Glass Carter. 1989).,"The $H$ and $K$ extinctions of the earth's atmosphere track each other very well, so that the $H-K$ colour is the most reliable; Glass Carter, 1989)."
+ The colours from Table. 3 have been plotted. in a istogram (fig 3)., The colours from Table 3 have been plotted in a histogram (fig 3).
+ Those with errors in excess of 0.4 mag rave been omitted., Those with errors in excess of 0.4 mag have been omitted.
+ There is a strong tendency for the colour temperatures o clump at particular values., There is a strong tendency for the colour temperatures to clump at particular values.
+ The best-determined. colours are d£[NL whose histogram peak at. 1.06 corresponds to a temperature around. 1600Ix.. believed. to be the highest emperature at which cust can exist without sublimation.," The best-determined colours are $H-K$, whose histogram peak at 1.06 corresponds to a temperature around 1600K, believed to be the highest temperature at which dust can exist without sublimation."
+ These wavelengths may. be the least contaminated by the al of the nuclear ultraviolet continuum or an extra dust component., These wavelengths may be the least contaminated by the tail of the nuclear ultraviolet continuum or an extra dust component.
+ Taking only galaxies with errors in dffy < 0.2. the average colour in the 0.95. 1.05 and 1.15 columns of the histogram is {4dy = 1.056 + 0.016.," Taking only galaxies with errors in $H-K$ $<$ 0.2, the average colour in the 0.95, 1.05 and 1.15 columns of the histogram is $H-K$ = 1.056 $\pm$ 0.016."
+ In Jdf. the typical colour temperature is arounc 200018. Lt is likely that the tail of the ultraviolet componen contributes to the J banc (1.255) more than to the ff (1.65pm). leading to an apparently hotter colour.," In $J-H$, the typical colour temperature is around 2000K. It is likely that the tail of the ultraviolet component contributes to the $J$ band $\mu$ m) more than to the $H$ $\mu$ m), leading to an apparently hotter colour."
+ This possibility is considered in detail for 33783. for which UDVRE data are available. by Glass (1992).," This possibility is considered in detail for 3783, for which $UBVRI$ data are available, by Glass (1992)."
+" In 33783 the contribution to the J band from the extrapolated (BV"" Hux is nearly equal to the flux from the short-wavelength tail of the blackbody raciation from cust at 190019. IndyL. the colour temperature is lower. about 13001400Ilx. and may represent dust slightly further [rom the nucleus that that which. predominates in ff ancl AN."," In 3783 the contribution to the $J$ band from the extrapolated $UBV$ flux is nearly equal to the flux from the short-wavelength tail of the blackbody radiation from dust at 1500K. In $K-L$, the colour temperature is lower, about 1300--1400K, and may represent dust slightly further from the nucleus that that which predominates in $H$ and $K$."
+ Lt is »ossible that there is a series of zones. proceeding outwarc rom the nucleus. with decreasing dust temperatures.," It is possible that there is a series of zones, proceeding outward from the nucleus, with decreasing dust temperatures."
+ Lt is interesting to note that an optically thick spherical dust shel waving a radius of 470d (as measured for Fairall 9) and a emperature of 14001 would have an L-band luminosity of 31074 +., It is interesting to note that an optically thick spherical dust shell having a radius of 470d (as measured for Fairall 9) and a temperature of 1400K would have an $L$ -band luminosity of 3 $^{24}$ W $^{-1}$.
+" ""Phis is about 8 times the observed Luminosity eiven in table 1.", This is about 8 times the observed luminosity given in table 1.
+ Fairall 9 has €.D ~ 0:50. indicating tha here is almost no circummuclear absorption along the Line of sight.," Fairall 9 has $U-B$ $\sim$ –0.85, indicating that there is almost no circumnuclear absorption along the line of sight."
+ Thus part of the factor of S may result from a dus distribution which is not spherical but. possibly toroical., Thus part of the factor of 8 may result from a dust distribution which is not spherical but possibly toroidal.
+ ]t is noticeable that the variable part of the Εαν from the Sevlert 2 galaxy 11068 is much redcder in 44dy and Aoo£ than the average for Sevfertsas a whole., It is noticeable that the variable part of the flux from the Seyfert 2 galaxy 1068 is much redder in $H-K$ and $K-L$ than the average for Seyfertsas a whole.
+ This is an indication that its nucleus is highly obscured. with fyoy —247 1.056 2 1414 £0387 or Egv =THOH2Z0 CA =," This is an indication that its nucleus is highly obscured, with $E_{H-K}$ = 2.47 – 1.056 = 1.414 $\pm$ 0.37, or $E_{B-V}$ = 7.6 $\pm$ 2.0 $A_V$ ="
+siuilar result for the starburst galaxies they studied. in which they detected an auti-correlation of «µη with μιι) san).,"similar result for the starburst galaxies they studied, in which they detected an anti-correlation of $_{\rm FIR}$ with $\mu$ $\mu$ m)."
+" ILowever. performing the sale test for the qo, ratios on our dwart galaxy suuple which spaus a range of fux ratios -0.1 «log(F(60 /FOLO0)<0.2 we find no such clear trend."," However, performing the same test for the $_{24}$ ratios on our dwarf galaxy sample which spans a range of flux ratios -0.4 $<$ $<$ 0.2 we find no such clear trend."
+ It is conceivable pan)that this is partly. because our siuuple is too stall to identify a trenddue to the intrinsic scatter n qo., It is conceivable that this is partly because our sample is too small to identify a trenddue to the intrinsic scatter in $_{24}$.
+ Both the radio aud infrared euiission can be used το estimate the star formation rates (SFRs) in ealaxics., Both the radio and infrared emission can be used to estimate the star formation rates (SFRs) in galaxies.
+ However. these correlations depend on a nunber of paralucters inchiding dust content. optical depth aud ietallicitv.," However, these correlations depend on a number of parameters including dust content, optical depth and metallicity."
+ This topic aud poteutial caveats have beeu discussed. oxtensively in the literature (seeCondo1992:IKeunicutt1998.—audreferences therein).," This topic and potential caveats have been discussed extensively in the literature \citep[see][ and references
+therein]{Condon92,Kennicutt98}."
+.. For dwart galaxics. the topic has been addressed iu detail bv IHopkiusetal.(2002).," For dwarf galaxies, the topic has been addressed in detail by \citet{Hopkins02}."
+". The recent wealth of data frou,Spitzer has also provided sufficient motivation to establish a calibration for the SFR using the imfrired.", The recent wealth of data from has also provided sufficient motivation to establish a calibration for the SFR using the infrared.
+ Wuetal.(2005). aud more recently Calzettietal.(2007) have explored this topic using a Lhuge sample of star forming regions aud ucarby galaxies.," \citet{Wu05}
+ and more recently \citet{Calzetti07} have explored this topic using a large sample of star forming regions and nearby galaxies."
+ Iu Fig., In Fig.
+ |o plot the IR estimated SER for our sample using the calibrations proposed by Wuetal.(2005). and Calzettietal.(2007) respectively: as a function of the well-known radio to SFR formula of Condon(1992): ↽∕∏∐∖∖↖⊽∏↸∖↑⋜↧↕∙⊔∩∩⋅↱⊐⋟↖↖↽∪∏↘↽↕↴∖↴⋝⋜↧↴∖↴↸∖≺↧∪∐⋜↧∶↴⋁↕∪⋝⋜↧↕ correlation without separating the low metallicity sources from metal ricli galaxies., 4 we plot the IR estimated SFR for our sample using the calibrations proposed by \citet{Wu05} and \citet{Calzetti07} respectively: as a function of the well-known radio to SFR formula of \citet{Condon92}: The \citet{Wu05} work is based on a global correlation without separating the low metallicity sources from metal rich galaxies.
+ As can be seen iu Fig., As can be seen in Fig.
+ ba eood aerectuent exists between the radio and IR estimated SFRs (indicated by diamonds).," 4, a good agreement exists between the radio and IR estimated SFRs (indicated by diamonds)."
+ If we use the more receut calibration by Calzettietal.(2007) ou SFRs from 214a huninosities. we fud that most of the mid-IR estimated SFRs (Guarked as filled circles) are located below the L:l proportionality line. aud. thus they are cousisteutlv lower than both the SERs estimated from the racio or fron Wuoetal(2005).," If we use the more recent calibration by \citet{Calzetti07} on SFRs from $\mu$ m luminosities, we find that most of the mid-IR estimated SFRs (marked as filled circles) are located below the 1:1 proportionality line, and thus they are consistently lower than both the SFRs estimated from the radio or from \citet{Wu05}."
+ This is not unexpected since eq., This is not unexpected since eq.
+ 2 is calibrated. based on the high metallicity sources that have significant 2102 emission. while most of our sources are metal-poor galaxies.," 2 is calibrated based on the high metallicity sources that have significant $\mu$ m emission, while most of our sources are metal-poor galaxies."
+" We should also note that Calzettietal.(2007) measure the SFR iu individual regious iu apertures within disks aud subtract a ""disk background. which could partly explain the deviations we see from our analysis."," We should also note that \citet{Calzetti07} measure the SFR in individual regions in apertures within disks and subtract a “disk background”, which could partly explain the deviations we see from our analysis."
+ These authors have also mentioned that the SFRs of low metallicity ealaxies would be uuderestinated by a factor of 21 depending onu how mnetalpoor the galaxies are., These authors have also mentioned that the SFRs of low metallicity galaxies would be underestimated by a factor of 2–4 depending on how metal-poor the galaxies are.
+ This deviation from a linear correlation is likely due to the lower opacitics for decreasing metal coutent (Walteral. 2007).., This deviation from a linear correlation is likely due to the lower opacities for decreasing metal content \citep{Walter07}. .
+ When we fit the filled circles ou Fig., When we fit the filled circles on Fig.
+ b. we find that «log(SER[1. LII]3.98). which is consistent with the metallicity correction factor sugeested by Calzettietal.(2007).," 4, we find that $\times$ log(SFR[1.4GHz]/3.98), which is consistent with the metallicity correction factor suggested by \citet{Calzetti07}."
+. We should stress ounce more thoush that as Bell(2003) has noted. the eood agrecment between the IR aud radio estimates of SFRs does not necessarily iiem that these are the “true” SERs for the dwarf galaxies we study. but rather the conrpetition of the lower dust content and suppressed svuchrotrou enmissiou balances cach other.," We should stress once more though that as \citet{Bell03} has noted, the good agreement between the IR and radio estimates of SFRs does not necessarily mean that these are the “true” SFRs for the dwarf galaxies we study, but rather the competition of the lower dust content and suppressed synchrotron emission balances each other."
+ Despite this overall agreciuent between the IR/racdio correlation in our dwarf sample aud the corresponding values of normal galaxies. the two lowest metallicity ealaxies iu our sample. [Zwls and SBS0335-052E. deviate markedly from the correlations aud in opposite directions (see Fig. 31," Despite this overall agreement between the IR/radio correlation in our dwarf sample and the corresponding values of normal galaxies, the two lowest metallicity galaxies in our sample, IZw18 and SBS0335-052E, deviate markedly from the correlations and in opposite directions (see Fig. )."
+) As the two most well-studied BCDs. it is interesting to inspect the IR/radio properties of these two galaxies aud a comparison oftheir physical piuuueters can be found in Table 3..," As the two most well-studied BCDs, it is interesting to inspect the IR/radio properties of these two galaxies and a comparison of their physical parameters can be found in Table \ref{tab3}. ."
+ As indicated by theSpitzer IRS spectrum of SBS0335-052E (IIoucketal.2001b).. the SED of the galaxy has an unusually fat mid-IR slope that peaks at —2Njun. IIuuntetal.(2005a) usedDUSTY models (Ivezié.ctal.1999) to fit its SED and found a qgigvalue of 2.0 for this galaxy. lower than the typical qgig of ~2.3 expected for late type systems.," As indicated by the IRS spectrum of SBS0335-052E \citep{Houck04b}, the SED of the galaxy has an unusually flat mid-IR slope that peaks at $\sim$ $\mu$ m. \citet{Hunt05a} used models \citep{Ivezic99} to fit its SED and found a $_{\rm FIR}$value of 2.0 for this galaxy, lower than the typical $_{\rm FIR}$ of $\sim$ 2.3 expected for late type systems."
+ These caleulatious are based ou a 1.1GITz coutinuuu flux of ~O. [61012Jy. which corresponds to the compact part (673) of SBS0335-052E (Dalectal.2001:TIntet 20014)...," These calculations are based on a 1.4GHz continuum flux of $\sim$ mJy, which corresponds to the compact part ) of SBS0335-052E \citep{Dale01, Hunt04}. ."
+ Similarly. if we simply calculate the q ratio independent of auy modelling work and onlv use the observable parameters. we fud a qa= 2140-1 (assmuingthatthe\UPSTOpauflux 2007)...," Similarly, if we simply calculate the q ratio independent of any modelling work and only use the observable parameters, we find a $_{70}=$ $\pm$0.1 \citep[assuming that the MIPS 70\,$\mu$m flux
+density is 51.1$\pm$4.8\,mJy,][]{Engelbracht07}. ."
+ This is in good agreement with the qry= 2.16:0.17 for the, This is in good agreement with the $_{70}=$ $\pm$0.17 for the
+(3)dalsedulscar.br,(3)dalsdf.ufscar.br
+model with my=50 kpe.,model with $r_{\rm th}=50$ kpc.
+" The steeper ej,(I). profile in the particles with rj«50 kpe is due to a large number of more eccentric orbits (i... a racially anisotropic velocity ellipsoid)."," The steeper ${\sigma}_{\rm los}(R)$ profile in the particles with $r_{\rm p}<50$ kpc is due to a large number of more eccentric orbits (i.e., a radially anisotropic velocity ellipsoid)."
+ The dillerence in a.(22) becomes progressively larger with larger #. which rellects the fact that the outer selected. particles need to have more eccentric orbits to have rs<50 kpe.," The difference in ${\sigma}_{\rm los}(R)$ becomes progressively larger with larger $R$, which reflects the fact that the outer selected particles need to have more eccentric orbits to have $r_{\rm p}<50$ kpc."
+" The velocity dispersions within the central 200 kpe for all particles (05,541) and selected ones (σι) are 336 km s.* and 260 kn +. respectively. for this standard model."," The velocity dispersions within the central 200 kpc for all particles ${\sigma}_{\rm m,all}$ ) and selected ones ${\sigma}_{\rm m}$ ) are 336 km $^{-1}$ and 260 km $^{-1}$ , respectively, for this standard model."
+" Figure 2 shows that σι dependson rj«ny (25 S175 kpe) such that it is smaller for smaller rjj in the stancard model: m,σι is 0.54 for rjj=25 kpe and 0.99 [or my=175 kpe."," Figure 2 shows that ${\sigma}_{\rm m}$ dependson $r_{\rm p}<r_{\rm th}$ (25 $\le r_{\rm th} \le$ 175 kpc) such that it is smaller for smaller $r_{\rm th}$ in the standard model: ${\sigma}_{\rm m}/{\sigma}_{\rm m,all}$ is 0.54 for $r_{\rm th}=25$ kpc and 0.99 for $r_{\rm th}=175$ kpc."
+ The physical reason for this dependence is that the model with smaller rij has a larger number of particles with more eccentric orbits in the outer part of the simulated region (50 kpe</? <200 kpe)., The physical reason for this dependence is that the model with smaller $r_{\rm th}$ has a larger number of particles with more eccentric orbits in the outer part of the simulated region (50 $ \le R \le $ 200 kpc).
+ Figure 3 shows that the standard model with my=150 kpe shows a very minor difference in ay.(22) between all and selected particles. which confirms that ry is a key parameter for kinematics in the present models for a given eyo and mu.," Figure 3 shows that the standard model with $r_{\rm th}=150$ kpc shows a very minor difference in ${\sigma}_{\rm los}(R)$ between all and selected particles, which confirms that $r_{\rm th}$ is a key parameter for kinematics in the present models for a given $e_{\rm m,0}$ and ${\sigma}_0$."
+" Fieure 4 shows that a, depends very weakly on ei, for 0.5«equc0.7 in the sense that ey, is smaller for larger eq, (Le. more eccentric orbits)."," Figure 4 shows that ${\sigma}_{\rm m}$ depends very weakly on $e_{\rm th}$ for $0.5 \le e_{\rm th} \le 0.7$ in the sense that ${\sigma}_{\rm m}$ is smaller for larger $e_{\rm th}$ (i.e., more eccentric orbits)."
+" Lt is confirmed that oy,(I) is steeper in the selected particles with e2eq than in all ones for the central region (2«20 kpc).", It is confirmed that ${\sigma}_{\rm los}(R)$ is steeper in the selected particles with $e > e_{\rm th}$ than in all ones for the central region $R<20$ kpc).
+" οσο results strongly sugeest that riy is a much more important. parameter than ei, for velocity dispersions in the standard. model.", These results strongly suggest that $r_{\rm th}$ is a much more important parameter than $e_{\rm th}$ for velocity dispersions in the standard model.
+ These results in Figures 1.d imply that the observed lower velocity dispersions of UCD populations in clusters of galaxies have physical origins: the lower dispersions are unlikely to result from some uncertainties in observational methocs to derive σι lor à smaller number of UCDs., These results in Figures $1-4$ imply that the observed lower velocity dispersions of UCD populations in clusters of galaxies have physical origins: the lower dispersions are unlikely to result from some uncertainties in observational methods to derive ${\sigma}_{\rm m}$ for a smaller number of UCDs.
+" The three important parameter dependencees of σι aud σεις) on initial orbital properties of particles in. cluster models (Le. 604,0 and 05) are briellv summarized as follows."," The three important parameter dependences of ${\sigma}_{\rm m}$ and ${\sigma}_{\rm los}(R)$ on initial orbital properties of particles in cluster models (i.e., $e_{\rm m,0}$ and ${\sigma}_{\rm 0}$ ) are briefly summarized as follows."
+ Firstly. the dependence of ey on rq derived in the standard model is seen in the models with dillerent. ey. ancl συ.," Firstly, the dependence of ${\sigma}_{\rm m}$ on $r_{\rm th}$ derived in the standard model is seen in the models with different $e_{\rm m,0}$ and ${\sigma}_{\rm 0}$."
+" Figure 5 clearly shows that irrespective of eau and συ. σιµ is smaller for smaller ry. which confirms that ri, is a key parameter for kinematical properties of galaxy. populations in clusters."," Figure 5 clearly shows that irrespective of $e_{\rm m,0}$ and ${\sigma}_{\rm 0}$, ${\sigma}_{\rm m}$ is smaller for smaller $r_{\rm th}$, which confirms that $r_{\rm th}$ is a key parameter for kinematical properties of galaxy populations in clusters."
+" Secondly. σι is significantly steeper αμα systematically smaller in. the selected particles with Foμι than in all ones for models with cdillerent ey, ane ay. though the details of the profiles depends slightly. on Cuo and oy."," Secondly, ${\sigma}_{\rm los}(R)$ is significantly steeper and systematically smaller in the selected particles with $r<r_{\rm th}$ than in all ones for models with different $e_{\rm m,0}$ and ${\sigma}_{\rm 0}$, though the details of the profiles depends slightly on $e_{\rm m,0}$ and${\sigma}_{\rm 0}$ ."
+" Thirdly dependences of m, ancl σι,(11) on eiu are much weaker than those on ry, for dillerent models. which suggests that riy is à more important. parameter for kinematics of galaxy populations."," Thirdly dependences of ${\sigma}_{\rm m}$ and ${\sigma}_{\rm los}(R)$ on $e_{\rm th}$ are much weaker than those on $r_{\rm th}$ for different models, which suggests that $r_{\rm th}$ is a more important parameter for kinematics of galaxy populations."
+ Phese results for cillerent models are thus all consistent with those derived. for the standard model., These results for different models are thus all consistent with those derived for the standard model.
+solar abundances within the errors.,solar abundances within the errors.
+ The residuals clearly show the differences between the various instruments. especially at low energies.," The residuals clearly show the differences between the various instruments, especially at low energies."
+" It is obviously hard to separate the spectrum of the ""ring"" from a diffuse background which provides more signal than the source itself.", It is obviously hard to separate the spectrum of the “ring” from a diffuse background which provides more signal than the source itself.
+ Another source of uncertainty is the fact that the background is taken from a blank field accumulated from pointings at higher galactic latitudes. whereas the source i5 seen against a background about 3.0 degrees below the galactic plane.," Another source of uncertainty is the fact that the background is taken from a blank field accumulated from pointings at higher galactic latitudes, whereas the source is seen against a background about 3.0 degrees below the galactic plane."
+ We therefore repeated the fit with a local background selected from just outside radio remnant from the current observation. 1.8. the upper left region in reffig:rgb2..," We therefore repeated the fit with a local background selected from just outside radio remnant from the current observation, i.e. the upper left region in \\ref{fig:rgb2}."
+ This fit is relatively bad. - = 1.38 (479 d.o.f.)," This fit is relatively bad, $\chi^2_{\rm
+red}$ = 1.38 (479 d.o.f.)"
+" but the best fit parameters are roughly consistent with the above values inside their mutual errors: the power law slope is Γ=2.29+0.50. the temperature of the model is kT=0.40+40.02 kkeV. the ionization time n,xf=(2.7+0.3) sscem™ and the fitted absorption Ny=(0.3841.63)x107 em”."," but the best fit parameters are roughly consistent with the above values inside their mutual errors: the power law slope is $\Gamma = 2.29\pm0.50$, the temperature of the model is $\pm$ keV, the ionization time $_e\times t = (2.7\pm0.3)\times10^{10}$ $^{-3}$ and the fitted absorption $_{\rm H} =(0.38\pm1.63)\times10^{22}$ $^{-2}$."
+ The fitted over-abundances are similar as above. however with formally smaller errors which must be regarded with caution considering the unacceptable quality of the fit.," The fitted over-abundances are similar as above, however with formally smaller errors which must be regarded with caution considering the unacceptable quality of the fit."
+ The residuals of that fit are shown as the bottom panel of reffig:pos-1.., The residuals of that fit are shown as the bottom panel of \\ref{fig:pos-1}.
+ A comparison of the two fits indicates that the data quality is insufficient to accurately constrain all of the model parameters: at higher energies (power law slope) as well as in the soft band (absorption) the fitted parameters depend strongly on the background subtraction., A comparison of the two fits indicates that the data quality is insufficient to accurately constrain all of the model parameters: at higher energies (power law slope) as well as in the soft band (absorption) the fitted parameters depend strongly on the background subtraction.
+ We finally tried to determine the spectrum im the very eastern parts of W50. to the left of the bright ring.," We finally tried to determine the spectrum in the very eastern parts of W50, to the left of the bright ring."
+ The signal is rather low and the fit residuals remain noisy even after all of the obvious faint point sources have been removed from the data., The signal is rather low and the fit residuals remain noisy even after all of the obvious faint point sources have been removed from the data.
+ No simple model fits the data but a power law with model plus an extra line at 0.924+0.07 kkeV results in an acceptable fit with Uu = |.14 for 520 d.o.f., No simple model fits the data but a power law with model plus an extra line at $\pm$ keV results in an acceptable fit with $\chi^2_{\rm red}$ = 1.14 for 520 d.o.f.
+ The power law with E21.32£0.13 dominates completely the emission of the thermal component (kT = 0.3120.03 kkeV) - apart from some prominent emission lines.," The power law with $\Gamma =
+1.32\pm0.13$ dominates completely the emission of the thermal component (kT = $\pm$ keV) - apart from some prominent emission lines."
+" The radio remnant W50 has a very distinct shape consisting of a nearly spherical shell centered on SS433 plus two elongations. the ""ears"". in the direction of the Jets."," The radio remnant W50 has a very distinct shape consisting of a nearly spherical shell centered on SS433 plus two elongations, the “ears”, in the direction of the jets."
+ A commonly accepted model for the elongated morphology ts the interaction between the SNR shell and the out-flowing jets., A commonly accepted model for the elongated morphology is the interaction between the SNR shell and the out-flowing jets.
+" It should be mentioned that according to our understanding of the geometrical conditions of the jet causing the eastern ""ear"" is pointing towards us.", It should be mentioned that according to our understanding of the geometrical conditions of the jet causing the eastern “ear” is pointing towards us.
+ Thus the X-ray spectrum emitted from the jet is a tracer of the physical processes leading to the slowing down of the jet by transforming its kinetic energy into radiation., Thus the X-ray spectrum emitted from the jet is a tracer of the physical processes leading to the slowing down of the jet by transforming its kinetic energy into radiation.
+ The X-ray brightness distribution reveals some details which are not in agreement with the simple shell - jet interaction described above., The X-ray brightness distribution reveals some details which are not in agreement with the simple shell - jet interaction described above.
+ While the locus from to the emission centroid of the terminal shock shows good symmetry with respect to the outer X-ray and radio emission. the lenticular structure and its axis of emission towards certainly breaks this symmetry.," While the locus from to the emission centroid of the terminal shock shows good symmetry with respect to the outer X-ray and radio emission, the lenticular structure and its axis of emission towards certainly breaks this symmetry."
+ Further. the angular extent of the X-ray Jet (about 18° as seen from 4433) 1s much smaller than the precession angle of ~40° expected from the kinematic model.," Further, the angular extent of the X-ray jet (about $^\circ$ as seen from 433) is much smaller than the precession angle of $\sim 40^\circ$ expected from the kinematic model."
+ This indicates either some collimation of the precessing jets (Peter Eichler 1993) or is a consequence of the interaction between the jet and the SNR which generates secondary shock waves along the symmetry axis (Velaizzquez Raga 2000)., This indicates either some collimation of the precessing jets (Peter Eichler 1993) or is a consequence of the interaction between the jet and the SNR which generates secondary shock waves along the symmetry axis (Velázzquez Raga 2000).
+ The lenticular emission structure has an angle of 20-30 degrees to the jet propagation direction. thus it cannot be a simple standing shock perpendicular to the jet axis.," The lenticular emission structure has an angle of $-$ 30 degrees to the jet propagation direction, thus it cannot be a simple standing shock perpendicular to the jet axis."
+ There tis steadily increasing emission. starting outside the right boundary of the XMM field of view. about aaremin from (see Fig.66 of paperID. building up to a very sharp and relatively narrow peak about 35-40 aremin from (see reffig:cutfig 1).," There is steadily increasing emission, starting outside the right boundary of the XMM field of view, about arcmin from (see 6 of I), building up to a very sharp and relatively narrow peak about $-$ 40 arcmin from (see \\ref{fig:cutfig1}) )."
+" This indicates that from distances >6x10? em from the compact source. the kinetic energy of the out-flowing matter is transferred to radiating relativistic electrons at an increasing rate until an ""unstable! situation is reached at the emission peak. where suddenly a much larger fraction of the energy is dissipated."," This indicates that from distances $\ma 6\times10^{19}$ cm from the compact source, the kinetic energy of the out-flowing matter is transferred to radiating relativistic electrons at an increasing rate until an 'unstable' situation is reached at the emission peak, where suddenly a much larger fraction of the energy is dissipated."
+" Further away from the Jet bends southwards towards the terminal shock. the emission dropping to a very low level before entering the region of the terminal shock itself,"," Further away from the jet bends southwards towards the terminal shock, the emission dropping to a very low level before entering the region of the terminal shock itself."
+"From 55217 to 55311 MJD, the power law photon index varied between 1.6 and 2.4 (or a cut-off power law between 1.3-2.3, Table 3)).","From 55217 to 55311 MJD, the power law photon index varied between 1.6 and 2.4 (or a cut-off power law between 1.3–2.3, Table \ref{tab:cutoffPL}) )."
+" In the early stages of the outburst, a cut-off is needed in the spectral fits (the Comptonization model improves the X24. whereas it disappears in the softer states."," In the early stages of the outburst, a cut-off is needed in the spectral fits (the Comptonization model improves the ), whereas it disappears in the softer states."
+" Indeed, later data are better fit with a simple power law model (Table 2))."," Indeed, later data are better fit with a simple power law model (Table \ref{tab:para}) )."
+" This indicates a cooling of the corona as the thermal component gets stronger, and then a change of the emission mechanism at the state transition."," This indicates a cooling of the corona as the thermal component gets stronger, and then a change of the emission mechanism at the state transition."
+ This further confirms that the source made a transition from a hard to soft state., This further confirms that the source made a transition from a hard to soft state.
+" With the model of ?,, we obtained temperatures (ΚΤε) and optical depth (7) values before Rev. 916 that yield a Comptonization parameter (yxm.c7) of 0.10—0.22, typical for a BH in the rising phase of the LHS."," With the model of \cite{Titarchuk:1994}, we obtained temperatures $kT_{\rm e}$ ) and optical depth $\tau$ ) values before Rev. 916 that yield a Comptonization parameter $\times$$m_{\rm{e}}c^2$ ) of 0.10–0.22, typical for a BH in the rising phase of the LHS."
+" The bolometric flux extrapolated between 0.01 keV and 10 MeV (note that the bulk of emission is between 0.1 and 100 keV) rose from 0.8 to 2.9x 107? erg cm? s! while the keV flux rose from 1.4 to 14.6 x 107? erg cm""? s!.", The bolometric flux extrapolated between 0.01 keV and 10 MeV (note that the bulk of emission is between 0.1 and 100 keV) rose from 0.8 to $\times$ $^{-8}$ erg $^{-2}$ $^{-1}$ while the 2--20 keV flux rose from 1.4 to 14.6 $\times$ $^{-9}$ erg $^{-2}$ $^{-1}$.
+" After that, the spectral parameters were compatible with the source making transition from the LHS to the HIMS (until ~ 55303.6 MJD), thena to softer states (to the SIMS until 55304.7 MJD, then quickly to the SS, ~ 55306.1 MJD, and finally back to the SIMS and HIMS (see Fig. 3))"," After that, the spectral parameters were compatible with the source making a transition from the LHS to the HIMS (until $\sim$ 55303.6 MJD), then to softer states (to the SIMS until 55304.7 MJD, then quickly to the SS, $\sim$ 55306.1 MJD, and finally back to the SIMS and HIMS (see Fig. \ref{hid}) )"
+" around 55320 MJD (?),, but this is beyond the scope of our paper since it happened afterthe observations listed in our Tables."," around 55320 MJD \citep{Motta10b}, but this is beyond the scope of our paper since it happened afterthe observations listed in our Tables."
+" Observations at other wavelengths of wwere conducted with Swifi/UVOT from ~ 55280 to 55300 MJD (?),, aimed at fast optical/UV variability in the hard state and transition."," Observations at other wavelengths of were conducted with /UVOT from $\sim$ 55280 to 55300 MJD \citep{yu10}, aimed at fast optical/UV variability in the hard state and transition."
+ No rapid optical/UV variability was detected due to the low UVOT count rates., No rapid optical/UV variability was detected due to the low UVOT count rates.
+ Our long-term FTS optical light curves of this outburst are shown in Fig. 2.., Our long-term FTS optical light curves of this outburst are shown in Fig. \ref{lcall}. .
+ The source, The source
+spectral swapping.,spectral swapping.
+ Finally we briefly discuss whether the oscillation parameters taken in this paper are really valid in views of recent work whose focus is on clarifying the still-veiled , Finally we briefly discuss whether the oscillation parameters taken in this paper are really valid in views of recent work whose focus is on clarifying the still-veiled nature of collective neutrino oscillations.
+," Following \cite{duan10}, there are at least two conditions for the onset of collective neutrino oscillations in the case of inverted neutrino mass hierarchy."
+ST fraction, The first criteria should be satisfied in the so-called bipolar regime of the collective oscillation.
+al excess of neutrinos over ? iG the characteristic mass-squared value of ~2.4x1073 eV? is employed ermi coupling constant.," In the regime, the neutrino number density should exceed the critical value, where $\chi$ is the fractional excess of neutrinos over antineutrinos, $\Delta m^2$ is the characteristic mass-squared splitting (a typical value of $\sim 2.4\times 10^{-3}$ $^2$ is employed here), and $G_\mathrm{F}$ is Fermi coupling constant."
+ By using our wd can estimate x which is often eter (typically ~0.01-0.25) so far.," By using our simulation results, we can estimate $\chi$ which is often treated as a parameter (typically $\sim$ 0.01-0.25) so far."
+" The )n is given in Esteban-Preteletal. 1/Εν,— in the case of vanishing the number flux of νι."," The following estimation is given in \cite{este07}, that is $\chi\simeq F_{\nu_e}/F_{\bar\nu_e}-1$ in the case of vanishing $F_{\nu_x}$, where $F_{\nu_i}$ is the number flux of $\nu_i$."
+ From Figure that x~ 0.2-0.3 for 100-400 ms after ," From Figure \ref{fig:chi}, it can be seen that $\chi\sim$ 0.2-0.3 for 100-400 ms after bounce."
+"typical number density in the 00-300 km) can be estimated as, (lit condition issatisfied!!.."," Since the typical number density in the post-shock region $r\sim$ 200-300 km) can be estimated as, therefore, the first condition is."
+ The ond criteria is related to the decoherence of πμ SOR by matter., The second criteria is related to the decoherence of collective oscillations by matter.
+" In order to overwhelm r[kthbe suppression by the decohenrence, the following condition should be satisfied where n, is the number density of electrons where the decoherence takes place."," In order to overwhelm the suppression by the decohenrence, the following condition should be satisfied where $n_e$ is the number density of electrons where the decoherence takes place."
+" This is equivalent to,"," This is equivalent to,"
+In order to see if the radiooptical spectra can be fit by a single power-law which is absorbed by frec-free emission at longer radio wavelengths. the four spectra in Fig 5 have been by the following simple mocel where So is the amplitude of the power law. à is the unabsorbed spectral index and τι ,"In order to see if the radio–optical spectra can be fit by a single power-law which is absorbed by free-free emission at longer radio wavelengths, the four spectra in Fig 5 have been fit by the following simple model where $S_0$ is the amplitude of the power law, $\alpha$ is the unabsorbed spectral index and $\tau_1$ "
+beams) the predicted: fraction of pulsars with interpulses drops to the observed. value.,beams) the predicted fraction of pulsars with interpulses drops to the observed value.
+ Also the predicted. period distribution of the pulsars with interpulses (thin solid. line Fig. 11) , Also the predicted period distribution of the pulsars with interpulses (thin solid line Fig. \ref{Fig_ip_period}) )
+now agrees well with the observations. indicating hat the model successfully. predicts the observed: undoer-abundance of long period pulsars with interpulses.," now agrees well with the observations, indicating that the model successfully predicts the observed under-abundance of long period pulsars with interpulses."
+ This can »e understood because lone period. pulsars tend. to have arger characteristic ages. which means that they tend to be more aligned. which reduces the probability that the beams of both magnetic poles interseet the line of sight.," This can be understood because long period pulsars tend to have larger characteristic ages, which means that they tend to be more aligned, which reduces the probability that the beams of both magnetic poles intersect the line of sight."
+ Finally. »ecause the a distribution depends on P in this case. the oedieted: correlation between A® and J? will deviate from 1/2.," Finally, because the $\alpha$ distribution depends on $P$ in this case, the predicted correlation between $\Delta\phi$ and $P$ will deviate from $-1/2$."
+ The slope of the correlation becomes Latter (the slope js | 0.40). as is observed.," The slope of the correlation becomes flatter (the slope is $-0.40$ ), as is observed."
+ Finally. one can consider a mocel which both includes the ellipticity of the beam and alignment.," Finally, one can consider a model which both includes the ellipticity of the beam and alignment."
+ In that. case the timescale of alignment is much longer (juin=210(0 ). because both cllectsB reduce the predicted.H number of interpulses.," In that case the timescale of alignment is much longer $\tau_\mathrm{align}=2\times10^9$ ), because both effects reduce the predicted number of interpulses."
+ As a consequence. the predicted. period distribution of the pulsars with interpulses (dashed line Fig. 13) ," As a consequence, the predicted period distribution of the pulsars with interpulses (dashed line Fig. \ref{Fig_ip_period}) )"
+does not deviate much from the model without alignment and therefore does not fit the data as well as the model with circular beams., does not deviate much from the model without alignment and therefore does not fit the data as well as the model with circular beams.
+ If pulsar beams align over time. it would imply that the total population of pulsars has an a distribution which is skewed to low values compared with a random clistribution.," If pulsar beams align over time, it would imply that the total population of pulsars has an $\alpha$ distribution which is skewed to low values compared with a random distribution."
+ This allects the average beaming fraction. because it is a function of a (e.g. Eq., This affects the average beaming fraction because it is a function of $\alpha$ (e.g. Eq.
+ 7 of 2))., 7 of \citealt{tm98}) ).
+ The beamine fraction is the fraction of the celestial sphere swept out by the beams of a pulsar. hence it is the probability that a pulsar is observable for a random line of sight.," The beaming fraction is the fraction of the celestial sphere swept out by the beams of a pulsar, hence it is the probability that a pulsar is observable for a random line of sight."
+ Aligned. beams are less likely to intersect the line of sight and are therefore loss likely to be observed., Aligned beams are less likely to intersect the line of sight and are therefore less likely to be observed.
+ The average beaming fraction for the observed population of pulsars. assuming circular beams. is and for the model with and without à evolution respectively.," The average beaming fraction for the observed population of pulsars, assuming circular beams, is and for the model with and without $\alpha$ evolution respectively."
+ For à random o distribution the fraction of. pulsars with beams which continuously intersect the line of sight is tiny (0.010). because of the à dependence of the beaming fraction in combination with the ellect that aligned. pulsars are Less likely to be formed in the first place.," For a random $\alpha$ distribution the fraction of pulsars with beams which continuously intersect the line of sight is tiny ), because of the $\alpha$ dependence of the beaming fraction in combination with the effect that aligned pulsars are less likely to be formed in the first place."
+ This means that basically no pulsars with extremely wide profiles are expected in the known population of pulsars., This means that basically no pulsars with extremely wide profiles are expected in the known population of pulsars.
+ This fraction appears too low. as we have already argued that some of the interpulses in Table 2. originate [rom a single magnetic pole.," This fraction appears too low, as we have already argued that some of the interpulses in Table \ref{TableInterpulses} originate from a single magnetic pole."
+ A more realistic fraction is predicted by the model assuming circular beams including e evolution (1.6%)). while za is too long to make a dillerence in the case of non-circular beams.," A more realistic fraction is predicted by the model assuming circular beams including $\alpha$ evolution ), while $\tau_\mathrm{align}$ is too long to make a difference in the case of non-circular beams."
+ This can be seen as additional evidence for circular beams and a significant elfect of a evolution., This can be seen as additional evidence for circular beams and a significant effect of $\alpha$ evolution.
+ We have shown that the observed fraction of pulsars with interpulses. their period distribution. the fraction of pulsars with extremely wide profiles and the observed. pulse width versus pulse period correlation (as measured in our data and that of 2)) is inconsistent with a random o distribution. even when non-circular beams are considered.," We have shown that the observed fraction of pulsars with interpulses, their period distribution, the fraction of pulsars with extremely wide profiles and the observed pulse width versus pulse period correlation (as measured in our data and that of \citealt{gl98}) ) is inconsistent with a random $\alpha$ distribution, even when non-circular beams are considered."
+ The main problem is that too many pulsars with interpulses are. precieted. especially those with long periods.," The main problem is that too many pulsars with interpulses are predicted, especially those with long periods."
+ A number of explanations are possible., A number of explanations are possible.
+ First. one could assume that long period pulsars are less ikely to have two active poles.," First, one could assume that long period pulsars are less likely to have two active poles."
+ Not only is this assumption ad-hoc. it also does not explain the observed: correlation tween the pulse width and the pulse period.," Not only is this assumption ad-hoc, it also does not explain the observed correlation between the pulse width and the pulse period."
+ Our model is kept simple and intuitive by assuming that both magnetic voles always emit radio emission. the beams are directed in exactly. opposite directions and have equal widths: anc uminosities.," Our model is kept simple and intuitive by assuming that both magnetic poles always emit radio emission, the beams are directed in exactly opposite directions and have equal widths and luminosities."
+ We therefore aim to find a geometrical solution or the problem., We therefore aim to find a geometrical solution for the problem.
+ secondly. one could assume ai different powerlay relation between p and P.," Secondly, one could assume a different powerlaw relation between $\rho$ and $P$."
+ In order to correctly. predic he number of pulsars with interpulses anc their. perioc distribution the observed. powerlaw relation (Iq. 6)), In order to correctly predict the number of pulsars with interpulses and their period distribution the observed powerlaw relation (Eq. \ref{Eqrho}) )
+ shouk of steeper and the constant smaller., should be steeper and the constant smaller.
+ This could. physically mean that the fraction of the polar cap which Is active or the emission height is Z dependent., This could physically mean that the fraction of the polar cap which is active or the emission height is $P$ dependent.
+ Not only is such a mode incompatible with the observations from. which Eq., Not only is such a model incompatible with the observations from which Eq.
+ Gis derived. it also results in a highly unrealistic pulse width distribution (the predicted. pulse widths are too narrow anc the correlation with P? is too steep).," \ref{Eqrho} is derived, it also results in a highly unrealistic pulse width distribution (the predicted pulse widths are too narrow and the correlation with $P$ is too steep)."
+ Thirdly. the pulsar beam may. not be circular. which is expected if the polar cap is bounced by the last open magnetic field. lines.," Thirdly, the pulsar beam may not be circular, which is expected if the polar cap is bounded by the last open magnetic field lines."
+ Although elliptical beams which are compressed in the plain containing the magnetic ancl the rotation axis could help to reduce the number of predicted interpulses. it is not enough.," Although elliptical beams which are compressed in the plain containing the magnetic and the rotation axis could help to reduce the number of predicted interpulses, it is not enough."
+ Moreover. the cllipticity is not expected to be 2 dependent.," Moreover, the ellipticity is not expected to be $P$ dependent."
+ “Phis implies that this moclel. ike that for circular. beams. does not correctly predict he period distribution of the pulsars interpulses and. the ruse width: period correlation.," This implies that this model, like that for circular beams, does not correctly predict the period distribution of the pulsars interpulses and the pulse width period correlation."
+ Ehe fraction of pulsars with extemely wide profiles is also too low., The fraction of pulsars with extemely wide profiles is also too low.
+ The only geometrical effect. left; as far as we can see. o explain the observations is to relax the assumption that o ds random.," The only geometrical effect left, as far as we can see, to explain the observations is to relax the assumption that $\alpha$ is random."
+ In order to explain the observed. fraction of »ulsars with interpulses the a distribution should be skewed o low values., In order to explain the observed fraction of pulsars with interpulses the $\alpha$ distribution should be skewed to low values.
+ The distribution for short period. pulsars should. be skewed less than that of long period pulsars in order to explain the observed: pulse period distribution of he pulsars with interpulses., The distribution for short period pulsars should be skewed less than that of long period pulsars in order to explain the observed pulse period distribution of the pulsars with interpulses.
+ The skewness therefore appears o be caused by an of a rather than by a non-random birth distribution which is fixed in time., The skewness therefore appears to be caused by an of $\alpha$ rather than by a non-random birth distribution which is fixed in time.
+ In the previous section it has been shown that the moce assuming circular beams can explain the observations wel when the magnetic axis is allowed to align from a random. distribution at birth with a timescale of ~7.10 vears., In the previous section it has been shown that the model assuming circular beams can explain the observations well when the magnetic axis is allowed to align from a random distribution at birth with a timescale of $\sim7\times10^7$ years.
+ The fact that the fraction of pulsars with interpulses. their perioc dependence. the fraction of pulsars with extremely wide profiles and the pulse width distribution can be explainec by adding just one extra model parameter is encouraging.," The fact that the fraction of pulsars with interpulses, their period dependence, the fraction of pulsars with extremely wide profiles and the pulse width distribution can be explained by adding just one extra model parameter is encouraging."
+ Moreover. bv allowing à to evolve it is not necessary (o abandon the observed ppP relation (I2q. 6))," Moreover, by allowing $\alpha$ to evolve it is not necessary to abandon the observed $\rho-P$ relation (Eq. \ref{Eqrho}) )"
+ or to resor in ad-hoc assumptions about magnetic poles which are no alwavs active., or to resort in ad-hoc assumptions about magnetic poles which are not always active.
+ The timescale is comparible with the timescale derive fromthe observed à distribution by ?.., The timescale is comparible with the timescale derived fromthe observed $\alpha$ distribution by \cite{tm98}.
+ Measuring à is far from trivial and for most pulsars it is. at best. only poorly constrained.," Measuring $\alpha$ is far from trivial and for most pulsars it is, at best, only poorly constrained."
+ Lt should be stressed that although we ciscuss à evolution. we do not rely on (indirect) measurements of a.," It should be stressed that although we discuss $\alpha$ evolution, we do not rely on (indirect) measurements of $\alpha$ ."
+ TFherefore our results are an important addition to the existing evidence for alignment., Therefore our results are an important addition to the existing evidence for alignment.
+below which it docs not fall further.,below which it does not fall further.
+ This is qualitatively reproduced at all values of σ. with oulv small changes iu the minimum values of BR. for σ in the rauge of es observed for local dSph galaxies.," This is qualitatively reproduced at all values of $\sigma$, with only small changes in the minimum values of $R_{\rm hm}$ , for $\sigma$ in the range of values observed for local dSph galaxies."
+ This is iuteresting. as it offers a natural explanation for local dSph galaxies showing a minimi lower value for their projected halfhelt radi of around 15ypc. With most line around a factor of 2 above this critical luit. as noticed by Ciliuore et al. (," This is interesting, as it offers a natural explanation for local dSph galaxies showing a minimum lower value for their projected half-light radii of around $150\rm{pc}$, with most lying around a factor of 2 above this critical limit, as noticed by Gilmore et al. ("
+2007) and seen from Table (1).,2007) and seen from Table (1).
+ Indeed. all our λοςcls coustrained by he observational values of σ aud Ry for the local «Splis. occur withiu the flat region of the Fyyavspy space.," Indeed, all our models constrained by the observational values of $\sigma$ and $R_{\rm hl}$ for the local dSphs, occur within the flat region of the $R_{\rm hm} \rm{vs.} \rho_{0}$ space."
+ We cau now try to understand the assigned Newtonian values of MEL for local «Sph galaxies aud the scalines thev show with total luminosity or absolute magnitude. e.9.. Mateo (1998).," We can now try to understand the assigned Newtonian values of $M/L$ for local dSph galaxies and the scalings they show with total luminosity or absolute magnitude, e.g., Mateo (1998)."
+ The assigned Newtoniau values of M/£E will never be far from (e.c. Cülimore et al., The assigned Newtonian values of $M/L$ will never be far from (e.g. Gilmore et al.
+" 2007) Given our results in. the previous section. or alternatively taking the observed. Ry, of around 300pc as an empirical fact. we cau evaluate the above equation to first order at Aj=0.3."," 2007) Given our results in the previous section, or alternatively taking the observed $R_{\rm hl}$ of around $300\rm{pc}$ as an empirical fact, we can evaluate the above equation to first order at $R_{\rm hl}=0.3$."
+ Also. for an average old stellar population. we can take a constant 11 value of 5 as a representative value for the local dSphls. as our inferences vield. aud replace 0? in equation (9) for the corresponding value through equation (10) to vield ποσο =the absolute magnitude Asy-=2HloetLiar)|L82 and takine the logarithin of the above equation gives Figure (1) now gives a plot of equatiou (12). superimposed on recent deteruinations of the Newtounau AJ/L values for a larger sample of local dSph galaxies. using reported values of L to calculate the values plotted on the x-axis. and of L. 0. aud 5; to calculate Newtonian AJ/L values through equation (10).," Also, for an average old stellar population, we can take a constant $M/L$ value of 5 as a representative value for the local dSphs, as our inferences yield, and replace $\sigma^{2}$ in equation (9) for the corresponding value through equation (10) to yield Introducing the absolute magnitude $M_{V} =-2.5 \rm{log}(L_{\rm tot}) + 4.83$ and taking the logarithm of the above equation gives Figure (4) now gives a plot of equation (12), superimposed on recent determinations of the Newtonian $M/L$ values for a larger sample of local dSph galaxies, using reported values of L to calculate the values plotted on the x-axis, and of L, $\sigma$, and $R_{hl}$ to calculate Newtonian $M/L$ values through equation (10)."
+ Error bars give extreme lo coufideuce intervals. inchludiug uncertainties in all the parameters used.," Error bars give extreme $1\sigma$ confidence intervals, including uncertainties in all the parameters used."
+ Data as sumnmauiuised in Walker et al (, Data as summarised in Walker et al. (
+2009a) (ising the revised version of the data for the classical dSphs. as given in Walker et.,"2009a) (using the revised version of the data for the classical dSphs, as given in Walker et."
+ al 20094). aud references therein.," al 2009d), and references therein."
+ The good match is evident. extending for over 13 orders of magnitude in My which is even more remarkable because the trend found extends towards the sanallest newly discovered dSpli svsteuis.," The good match is evident, extending for over 13 orders of magnitude in $M_{V}$, which is even more remarkable because the trend found extends towards the smallest newly discovered dSph systems."
+ We see that the expectations of the uodel for the inferred Newtomian AJ/L values of local «Sph galaxies. which are modelled as equilibrium isothermal solutions to equation (8) aud having typical radii of close to 0:3kpc. aeree very well with iude])vndent measurements.," We see that the expectations of the model for the inferred Newtonian $M/L$ values of local dSph galaxies, which are modelled as equilibrium isothermal solutions to equation (8) and having typical radii of close to $0.3 \rm{kpc}$ agree very well with independent measurements."
+ The «ια scatter in Figure (1) bevond observational error bars is compatible with the variious in the actual intrinsic ML values for the indivicual ealaxies. and variations of fj around the average 0.3kpce used in equation (12). which is not a ft to the daa. but a first-order estimate within our prescription.," The small scatter in Figure (4) beyond observational error bars is compatible with the variations in the actual intrinsic $M/L$ values for the individual galaxies, and variations of $R_{\rm hl}$ around the average $0.3\rm{kpc}$ used in equation (12), which is not a fit to the data, but a first-order estimate within our prescription."
+ A very situple explanation for he (QM/L)y values assigned to all well-studied dSpls is thus naturally afforded * our asstuuptions for the uost average intrinsic AL/L aud A values. without the weed of invokiug complex astroplivsical processes (6.8. idal disruption. strong outfiows. etc.)," A very simple explanation for the $(M/L)_{N}$ values assigned to all well-studied dSphs is thus naturally afforded by our assumptions for the most average intrinsic $M/L$ and $R_{\rm hl}$ values, without the need of invoking complex astrophysical processes (e.g. tidal disruption, strong outflows, etc.)"
+ or fine-tuning any xuanueters., or fine-tuning any parameters.
+ To first order. «ne can understand the well-established scalines of fietre (1) as the extensiou of he Tully-Fisher relation. which appears iu the MOND ornalisui. down to the dSp1 reeinie.," To first order, one can understand the well-established scalings of figure (4) as the extension of the Tully-Fisher relation, which appears in the MOND formalism, down to the dSph regime."
+ We sed à MOND interpolation function that cau be SOCIL as the addition of both the Newtoman acceleration and the low acceleration lint of the MOND formatlisi o- all scales (Bekeustein 2001)., We used a MOND interpolation function that can be seen as the addition of both the Newtonian acceleration and the low acceleration limit of the MOND formalism on all scales (Bekenstein 2004).
+ For the local dSph galaxies. where the application of he MOND formalisii has bee- most controversial we show that isothermal equilibria configurations characterised bywell-defined finite total masses and halfanass radii result. ering ML values d- agreement with naked stellaz populations.," For the local dSph galaxies, where the application of the MOND formalism has been most controversial, we show that isothermal equilibrium configurations characterised bywell-defined finite total masses and half-mass radii result, giving $M/L$ values in agreement with naked stellar populations."
+ As a result.no ark matter is now needed.," As a result,no dark matter is now needed."
+ The observed scaliugs in hi and assigned (AL/£)\ , The observed scalings in $R_{\rm hl}$ and assigned $(M/L)_{N}$ 
+he outer binary. aud the eravitational wave radiation nuescale of the inner binary wotld generally be short.,"the outer binary, and the gravitational wave radiation timescale of the inner binary would generally be short."
+ Also. the Nozai mechanisim could play a role in increasing he eccentricity of the πιο binary (Blaesetal.2002).," Also, the Kozai mechanism could play a role in increasing the eccentricity of the inner binary \citep{Blaesetal2002}."
+. So fu. our primary atteutioi has been directed Oo supermassive black holes iu aree ellipticals. since observational evidence for iassive dark objects is stronecst for large ellipticals.," So far, our primary attention has been directed to supermassive black holes in large ellipticals, since observational evidence for massive dark objects is strongest for large ellipticals."
+ If LSD ealaxies had ceutral dlack holes (which would be more ike iunterimeciate-iass dack holes than massive black holes). some of them must wave experienced major mereer events aud therefore are ikelv to have contained iuultiple black holes at some x)uUt.," If LSB galaxies had central black holes (which would be more like intermediate-mass black holes than massive black holes), some of them must have experienced major merger events and therefore are likely to have contained multiple black holes at some point."
+ Tere. again. the crucial question is whether or rot these black holes can merec.," Here, again, the crucial question is whether or not these black holes can merge."
+ Since the black hole nasses are sualler. if the central region of LSB galaxies is dominated by normal stars. loss-cone depletion is less effective. and black holes eau become very tight.," Since the black hole masses are smaller, if the central region of LSB galaxies is dominated by normal stars, loss-cone depletion is less effective, and black holes can become very tight."
+ Indeed. if we veeard the field particles m our simulations as ποια! stars with typical masses around a solar nass; we have performed simulations of DII binarics with masses up to 10!AZ... for which we have seeu that the hardening rate is still pretty high.," Indeed, if we regard the field particles in our simulations as normal stars with typical masses around a solar mass, we have performed simulations of BH binaries with masses up to $10^4 M_{\odot}$, for which we have seen that the hardening rate is still pretty high."
+ However. if the ceutral region is donmünated by CDAL as suggested by both theory and observations. the thermal relaxation time would be practically infinite and loss-cone depletion would occur instantly.," However, if the central region is dominated by CDM, as suggested by both theory and observations, the thermal relaxation time would be practically infinite and loss-cone depletion would occur instantly."
+ Iu the case of LSB galaxies. it is uulikelv. that two black holes meree during triple interactions. because the central potential is shallow.," In the case of LSB galaxies, it is unlikely that two black holes merge during triple interactions, because the central potential is shallow."
+ All black holes would be ejected from the pareut. galaxy before they would eot a chance o imerse., All black holes would be ejected from the parent galaxy before they would get a chance to merge.
+" Thus. unlike large ellipticals. LSB ealaxies are unlikely to display something like a Magorrian relation or Mo, σ relation."," Thus, unlike large ellipticals, LSB galaxies are unlikely to display something like a Magorrian relation or $M_{bh}$ $\sigma$ relation."
+ The authors thank Sverre Aarseth. Rainer Spurzem for discussions related to this work. aud Piet IIut for carefully reading the manuscript.," The authors thank Sverre Aarseth, Rainer Spurzem for discussions related to this work, and Piet Hut for carefully reading the manuscript."
+" This work is supported in part by Cwaut-in-aid in Scicutific Research in Priority Areas (15037203) aud B (13110058) of the Ministry of Education. Culture, Culture. Science aud Teclinologx. Japan."," This work is supported in part by Grant-in-aid in Scientific Research in Priority Areas (15037203) and B (13440058) of the Ministry of Education, Culture, Culture, Science and Technology, Japan."
+thanks to R. Cockcroft (McMaster) for contributing relevant references.,thanks to R. Cockcroft (McMaster) for contributing relevant references.
+" This research has made use of the NASA Astrophysics Data System Bibliographic services the Extragalactic Database (NED), and (ADS),Google."," This research has made use of the NASA Astrophysics Data System Bibliographic services (ADS), the NASA/IPAC Extragalactic Database (NED), and Google."
+" STSDAS NASA/IPACand PyRAF are products of the Space Telescope Science Institute, which is operated by AURA for NASA."," STSDAS and PyRAF are products of the Space Telescope Science Institute, which is operated by AURA for NASA."
+ We thank the anonymous referee for a constructive report and a prompt reply., We thank the anonymous referee for a constructive report and a prompt reply.
+to image the sub-millimetre galaxies. and UFTI on UKIRT to image the radio galaxies.,"to image the sub-millimetre galaxies, and UFTI on UKIRT to image the radio galaxies."
+ We have taken great care to ensure that all images are as well matched as possible in. A'-band. surface-brightness limit. designing integration times to compensate for the difference in telescope aperture between Gemini North and UKIRT. and for cosmological surface-brightness dimming within our samples.," We have taken great care to ensure that all images are as well matched as possible in $K$ -band surface-brightness limit, designing integration times to compensate for the difference in telescope aperture between Gemini North and UKIRT, and for cosmological surface-brightness dimming within our samples."
+ Also crucial to this study has been our ability (due to the queue-based scheduling system in operation at both elescopes) to demand the very best seeing conditions for all of our observations (PM£A.«0.6 aresec)., Also crucial to this study has been our ability (due to the queue-based scheduling system in operation at both telescopes) to demand the very best seeing conditions for all of our observations $FWHM < 0.6$ arcsec).
+ It is only with this powerful combination of consistent high-quality seeing and long exposure imes (typically 90 mins on Gemini) that it has proved possible to determine the sizes (ie. half-light radii) and basic morphologies (i.e. Sérrsic index) of these massive +—2 galaxies via based imaging., It is only with this powerful combination of consistent high-quality seeing and long exposure times (typically 90 mins on Gemini) that it has proved possible to determine the sizes (i.e. half-light radii) and basic morphologies (i.e. Sérrsic index) of these massive $z \simeq 2$ galaxies via ground-based imaging.
+ The paper is structured as follows., The paper is structured as follows.
+ In Section 2 we summarize he radio galaxy and sub-millimetre galaxy samples. and the existing supporting information.," In Section 2 we summarize the radio galaxy and sub-millimetre galaxy samples, and the existing supporting information."
+ In Section 3 we present the UKIRT and Gemini A&-band imaging observations. and in Section j we explain how these data were reduced.," In Section 3 we present the UKIRT and Gemini $K$ -band imaging observations, and in Section 4 we explain how these data were reduced."
+ Then in Section 5 we describe how we modelled these new images to extract the basic morphological properties of the galaxies., Then in Section 5 we describe how we modelled these new images to extract the basic morphological properties of the galaxies.
+ The results of the individual model fits are presented. summarized and discussed in Section 6. with the images. model fits and residual images provided in Appendix A. We also explore the robustness of our results. through simulations. non parametric tests of galaxy size. and the analysis of radio galaxy and sub-millimetre galaxy image stacks.," The results of the individual model fits are presented, summarized and discussed in Section 6, with the images, model fits and residual images provided in Appendix A. We also explore the robustness of our results, through simulations, non parametric tests of galaxy size, and the analysis of radio galaxy and sub-millimetre galaxy image stacks."
+ Finally. in Section 7 we attempt to place our results in the context of studies of other galaxy populations at both high and low redshift.," Finally, in Section 7 we attempt to place our results in the context of studies of other galaxy populations at both high and low redshift."
+ Our main conculsions are summarized in Section 8., Our main conculsions are summarized in Section 8.
+ The initial target sample for this study consisted of 39 sources. and comprised a complete sample of the 14 most radio-Iuminous radio galaxies at ο—2. and the 25 brightest j/m sources in the SCUBA 8-mly survey.," The initial target sample for this study consisted of 39 sources, and comprised a complete sample of the 14 most radio-luminous radio galaxies at $z \simeq 2$, and the 25 brightest ${\rm \mu m}$ sources in the SCUBA 8-mJy survey."
+ Of the 14 objects in the original radio galaxy sample. 8 were drawn from the 3CRR catalogue (Laing. Riley Longair 1983) and 6 from the equatorial sample of Best et al. (," Of the 14 objects in the original radio galaxy sample, 8 were drawn from the 3CRR catalogue (Laing, Riley Longair 1983) and 6 from the equatorial sample of Best et al. ("
+1999).,1999).
+ The 3CRR sample is complete to à 178-MHz flux-density limit of 10.9 Jy over an area of 4.2 steradians. and contains a total of 170 steep-spectrum sources over the redshift range Q.0«z2.5.," The 3CRR sample is complete to a 178-MHz flux-density limit of $S_{\rm 178\,MHz} > 10.9$ Jy over an area of 4.2 steradians, and contains a total of 170 steep-spectrum sources over the redshift range $0.0<z<2.5$."
+ These data are essentially doubled by the addition of the Best et al., These data are essentially doubled by the addition of the Best et al.
+" sample. which was designed to have a roughly equivalent selection criterion of Susx1H, SJJy. using the Molonglo Reference Catalogue."," sample, which was designed to have a roughly equivalent selection criterion of $S_{\rm 408\,MHz} > 5$ Jy, using the Molonglo Reference Catalogue."
+ From this combined survey we have defined a complete Pentericci et al., From this combined survey we have defined a complete Pentericci et al.
+ 2001) sample of the 1-4 most powerful. frequency-selected. radio sources in the redshift range 1.5.« 2.5.," 2001) sample of the 14 most powerful, low-frequency-selected, radio sources in the redshift range $1.5<z<2.5$ ."
+ One of the sources in this high-power sample. 2294. transpired to lie too close on the sky to a nearby star for deep ἐν -band imaging to be successfully secured.," One of the sources in this high-power sample, 294, transpired to lie too close on the sky to a nearby star for deep $K$ -band imaging to be successfully secured."
+ Thus. for the present study. the parent sample was trimmed to the 13 sources listed in Table I.," Thus, for the present study, the parent sample was trimmed to the 13 sources listed in Table 1."
+ The sub-millimetre galaxies were drawn from the SCUBA 8-mJy survey (Scott et al., The sub-millimetre galaxies were drawn from the SCUBA 8-mJy survey (Scott et al.
+ 2002)., 2002).
+ The SCUBA jin imaging in the 8- survey was split between the Lockman Hole East and Elais N2 fields. and covered a total area of 260 aremin? to a typical rms noise level of es)2.5 +.," The SCUBA ${\rm \mu m}$ imaging in the 8-mJy survey was split between the Lockman Hole East and Elais N2 fields, and covered a total area of 260 $^{2}$ to a typical rms noise level of $\sigma_{850}\simeq2.5$ $^{-1}$."
+ From the resulting source catalogue. all 25 sub-millimetre sources with a signal-to-noise S/N73.5 at jon were selected as potential targets for deep AX-band imaging in the present study.," From the resulting source catalogue, all 25 sub-millimetre sources with a signal-to-noise $S/N > 3.5$ at ${\rm \mu m}$ were selected as potential targets for deep $K$ -band imaging in the present study."
+ In Table | we list the 15 sub-millimetre sources which were actually observed. detected with sufficient signal:noise in. A-band to undertake a meaningful morphological analysis.," In Table 1 we list the 15 sub-millimetre sources which were actually observed, detected with sufficient signal:noise in $K$ -band to undertake a meaningful morphological analysis."
+ All of the radio galaxies in the |3-source sample have unambiguous near-infrared galaxy counterparts. as a result of the sub-aresec positional accuracy provided by the VLA detections of their radio cores.," All of the radio galaxies in the 13-source sample have unambiguous near-infrared galaxy counterparts, as a result of the sub-arcsec positional accuracy provided by the VLA detections of their radio cores."
+ However. as is well Known. the identification of the A-band counterparts of the sub-millimetre sources is more complicated. due to the large (FWHM 215 aresec) size of the Airy disc delivered by the 15-m JCMT at A850μα.," However, as is well known, the identification of the $K$ -band counterparts of the sub-millimetre sources is more complicated, due to the large (FWHM $\simeq 15$ arcsec) size of the Airy disc delivered by the 15-m JCMT at $\lambda \simeq 850\,{\rm \mu m}$."
+ However. for 12 of the 15 sub-millimetre sources listed in Table |. robust radio identifications (and hence positions accurate to  | aresec) were successfully deduced by Ivison et al. (," However, for 12 of the 15 sub-millimetre sources listed in Table 1, robust radio identifications (and hence positions accurate to $\simeq$ 1 arcsec) were successfully deduced by Ivison et al. ("
+"2002). via deep GGHz imaging with the Very Large Array (VLA) reaching a typical Sa detection limit of 9,aeg;2255v.","2002), via deep GHz imaging with the Very Large Array (VLA) reaching a typical $\sigma$ detection limit of $S_{1.4GHz} \simeq 25 {\rm \mu Jy}$."
+ As indicated in Table |. identifications for the remaining 3 sub-millimetre sources in our sub-sample were therefore by necessity determined on the basis of optical-infrared information only.," As indicated in Table 1, identifications for the remaining 3 sub-millimetre sources in our sub-sample were therefore by necessity determined on the basis of optical-infrared information only."
+ For the 2 outstanding ELATS N2 sources. the choice of identification was based on the /.—A colour derived from the existing (and now public) /-band HST ACS imaging of the ELAIS 2 field (PI: Almaini. PID: 9761). and the pre-existing UKIRT Áx-band imaging described by Ivison et al. (," For the 2 outstanding ELAIS N2 sources, the choice of identification was based on the $I-K$ colour derived from the existing (and now public) $I$ -band HST ACS imaging of the ELAIS N2 field (PI: Almaini, PID: 9761), and the pre-existing UKIRT $K$ -band imaging described by Ivison et al. ("
+2002).,2002).
+ Specifically. ollowing Pope et al. (," Specifically, following Pope et al. ("
+2005) the criterion that sub-millimetrας galaxies are typically very red in /.—A colour (see also Smail e al.,2005) the criterion that sub-millimetre galaxies are typically very red in $I-K$ colour (see also Smail et al.
+ 2004) was used to select the most likely A -band counterpart in he absence of a secure radio detection., 2004) was used to select the most likely $K$ -band counterpart in the absence of a secure radio detection.
+ In the case of 8850.04. Where no optical imaging was available at the commencemen of this study. the /v-band source closest to the sub-millimetre yosition was selected.," In the case of 850.04, where no optical imaging was available at the commencement of this study, the $K$ -band source closest to the sub-millimetre position was selected."
+ It should be noted that although 8850.04 is therefore potentially the most uninformed identification. no other xossible A -band counterpart is found within the SCUBA positiona error circle for this source.," It should be noted that although 850.04 is therefore potentially the most uninformed identification, no other possible $K$ -band counterpart is found within the SCUBA positional error circle for this source."
+ While it is to some extent desirable to partially remove the biases introduced by insisting on a robust radio identification. it is nonetheless true that these three non-radio selected galaxy identifications could be regarded as less secure than their radio-selected counterparts.," While it is to some extent desirable to partially remove the biases introduced by insisting on a robust radio identification, it is nonetheless true that these three non-radio selected galaxy identifications could be regarded as less secure than their radio-selected counterparts."
+ Because of this. they are explicitly flagged as a specitie subset in much of the subsequent analysis presented in this paper.," Because of this, they are explicitly flagged as a specific subset in much of the subsequent analysis presented in this paper."
+ Secure spectroscopic redshifts exist for all 13 radio galaxies (generally aided by bright emission lines) and are given in Table l., Secure spectroscopic redshifts exist for all 13 radio galaxies (generally aided by bright emission lines) and are given in Table 1.
+ As is well documented in the literature (e.g. Ivison et al., As is well documented in the literature (e.g. Ivison et al.
+ 2005). obtaining secure spectroscopic redshifts for sub-millimetre galaxies has proved much more difficult. not only because the identifications are harder to establish. but also because the correct," 2005), obtaining secure spectroscopic redshifts for sub-millimetre galaxies has proved much more difficult, not only because the identifications are harder to establish, but also because the correct"
+"The in-flight calibration of auyv astronomical mstruneut relics on a combination of previous measurements using lustrmmentation with known or traceable calibration accuracy, accurate models of standard astronomical sources and calibration sources internal to tle imstriuneut or facility.","The in-flight calibration of any astronomical instrument relies on a combination of previous measurements using instrumentation with known or traceable calibration accuracy, accurate models of standard astronomical sources and calibration sources internal to the instrument or facility."
+ Iu this paper we discuss the methods aud source models employed to couvert SPIRE data to physical wits aud discuss the accuracy of the calibration and any caveats that niust be placed on the cata., In this paper we discuss the methods and source models employed to convert SPIRE data to physical units and discuss the accuracy of the calibration and any caveats that must be placed on the data.
+ The SPIRE lustrument and its overall calibration scheme are described iu detail in ήπια et al. (2010))., The SPIRE instrument and its overall calibration scheme are described in detail in Griffin et al. \cite{griffin10}) ).
+" For the purposes of discussion of the calibration aud performance estimation we treat SPIRE as two separate iustruineuts that are spatially separated at the focal plane of the telescope: a three band anaging photometric camera with bands centred nominally on 250 ((PSW). 350 ({PAI). and 500 ((PLW)J with detector errays of £89, SS and 43 respectively feedhorn coupled NTD-bolometers (Turner et al 2001)). and à two band iniaging Fourier: transform. spectrometer (ETS) covering πάση ((SSW) and 305-6740 ((SLWJ with arrays of 37 and 19 detectors,"," For the purposes of discussion of the calibration and performance estimation we treat SPIRE as two separate instruments that are spatially separated at the focal plane of the telescope: a three band imaging photometric camera with bands centred nominally on 250 (PSW), 350 (PMW), and 500 (PLW) with detector arrays of 139, 88 and 43 respectively feedhorn coupled NTD-bolometers (Turner et al \cite{turner01}) ), and a two band imaging Fourier transform spectrometer (FTS) covering $\mu$ (SSW) and $\mu$ (SLW) with arrays of 37 and 19 detectors."
+ SPIRE contains two internal calibration sources which were desiyned to provide a rapidly modulated signal for relative qain calibration (the PCAL source) and a stable therinal source for calibration of the spectrometer and to balance the power fron. the telescope (the SCAL source)., SPIRE contains two internal calibration sources which were designed to provide a rapidly modulated signal for relative gain calibration (the PCAL source) and a stable thermal source for calibration of the spectrometer and to balance the power from the telescope (the SCAL source).
+" PCAL is placed. at an nage of the telescope secondary in the common light path of the two instruments and therefore can be used to stimulate all detectors in the photometer and spectrometer,", PCAL is placed at an image of the telescope secondary in the common light path of the two instruments and therefore can be used to stimulate all detectors in the photometer and spectrometer.
+ The SCAL source is placed ain the second input port of the spectromcter (Swinyard ct al. 2003))., The SCAL source is placed in the second input port of the spectrometer (Swinyard et al. \cite{swinyard03}) ).
+" Tt ds thercforc constantly eiewed during observations and provides a reference against which the spectrum: from. the telescope port is measured,", It is therefore constantly viewed during observations and provides a reference against which the spectrum from the telescope port is measured.
+" The intention before flight was that the SCAL soiree world be heated to a temperature: sufficient. to entirely spn” the spectrin from the telescope,", The intention before flight was that the SCAL source would be heated to a temperature sufficient to entirely “null” the spectrum from the telescope.
+ We discuss in Sect., We discuss in Sect.
+" 6 why this has proved to be unnecessarg for standard observing conditions,", \ref{spectro_pipe} why this has proved to be unnecessary for standard observing conditions.
+ The internal calibrators are only viewed through part of the optics chain and the overall astronomical calibration and performance of the Bstriinent and telescope together can only be cstablished by observation of astronomical sources as we describe. in this paper., The internal calibrators are only viewed through part of the optics chain and the overall astronomical calibration and performance of the instrument and telescope together can only be established by observation of astronomical sources as we describe in this paper.
+ The structure of the paper is as follows: in Sect., The structure of the paper is as follows: in Sect.
+ 2 we deseribe the initial in-flight performance evaluation of the instrament before going on ta cover the basic conversion. of the data from digital encoder: values. to physical units for the detectors. the beam. steering mirror (DSM) and the spectrometer mechanisin (SMEC) in Sect. d.," \ref{initial} we describe the initial in-flight performance evaluation of the instrument before going on to cover the basic conversion of the data from digital encoder values to physical units for the detectors, the beam steering mirror (BSM) and the spectrometer mechanism (SMEC) in Sect. \ref{dataproc}."
+ In Sect., In Sect.
+ 4 wo describe the characterisation of the tune response of the photometer bolomcters using donising radiation hits and in Sects., \ref{time_constant} we describe the characterisation of the time response of the photometer bolometers using ionising radiation hits and in Sects.
+" 5 and 6 we discuss how the photometric calibration was establishedfor the photometer and spectrometer,", \ref{photo_pipe} and \ref{spectro_pipe} we discuss how the photometric calibration was established for the photometer and spectrometer.
+ In Sect., In Sect.
+ Ἱ we present the results of comparing SPIRE observations to solar system objects with known or well constrained model fluc densities., \ref{testing_models} we present the results of comparing SPIRE observations to solar system objects with known or well constrained model flux densities.
+" The beuimesize and astrometric accuracy are discussed in Sect,", The beamsize and astrometric accuracy are discussed in Sect.
+" & and the werclenyth accuracy in Sect, 9.", \ref{beam} and the wavelength accuracy in Sect. \ref{wavelength_calibration}.
+ We sunmanarisc the present calibration status of SPIRE und draw some conchisions for future calibration activities in Sect. L0.., We summarise the present calibration status of SPIRE and draw some conclusions for future calibration activities in Sect. \ref{conclusions}.
+ The mitial performance verification and comparison between in-flight and ground performance was carried out using the PCAL source whilst the eryostat did was closed., The initial performance verification and comparison between in-flight and ground performance was carried out using the PCAL source whilst the cryostat lid was closed.
+ This allowed a direct, This allowed a direct
+the residual between the best-fit curve and the data. shows a clearly double-peaked shape.,"the residual between the best-fit curve and the data, shows a clearly double-peaked shape."
+ This double-peaked residual jugeests the svsteim is better fit bv two Ciassdans., This double-peaked residual suggests the system is better fit by two Gaussians.
+ Based on our analysis of the generalized listogram. we claim that NGC 5166 contains two closely spaced eroups iu CN band streugth. a result of intrinsic ranges iu carbou and nitrogen abundance aud low overall metallicity.," Based on our analysis of the generalized histogram, we claim that NGC 5466 contains two closely spaced groups in CN band strength, a result of intrinsic ranges in carbon and nitrogen abundance and low overall metallicity."
+ The lack of a gap in CN baud streneth between the two eroups of stars in Figure 5. is secu here as a fairly sanall distance between the peaks of the CNoweak aud CN-stroug ecueralized histograms., The lack of a gap in CN band strength between the two groups of stars in Figure \ref{fig5} is seen here as a fairly small distance between the peaks of the CN-weak and CN-strong generalized histograms.
+ Sinibulv metal-poor elobular clusters such as ALLS (Lee 2000)) ancl M53 (Martell et al. 2008a), Similarly metal-poor globular clusters such as M15 (Lee ) and M53 (Martell et al. )
+) also exhibit small separations between their CN-weals and C'N-stroug eroups aud a sanall overall range iu ο2500} (among red eiauts iu the three clusters 6.$(3839) las a standard deviation of 0.038 ip NGC 5166. 0.056 in M52. and 0.037 in MID).," also exhibit small separations between their CN-weak and CN-strong groups and a small overall range in $S(3839)$ (among red giants in the three clusters $\delta S(3839)$ has a standard deviation of $0.038$ in NGC 5466, $0.056$ in M53, and $0.037$ in M15)."
+ This sugeests that the range of carbon aud nitroseu abuudance in low-ietallicity globular clusters may be sanaller than is typical for more moderatc-metallicity clusters., This suggests that the range of carbon and nitrogen abundance in low-metallicity globular clusters may be smaller than is typical for more moderate-metallicity clusters.
+ In addition. since the measurement errors ol 3839) control the widths of the individual Ctaussiaus conrpridues the generalized histogram. they can have «ποιο effects on our interpretation of it. with very large lueasurement errors potentially rendering two closely spaced populations indistinguishable from a single broad distribution.," In addition, since the measurement errors on $S(3839)$ control the widths of the individual Gaussians comprising the generalized histogram, they can have strong effects on our interpretation of it, with very large measurement errors potentially rendering two closely spaced populations indistinguishable from a single broad distribution."
+ To understand the effects of nicasureimieut error on our eoncralized histogram. we performed a Monte Carlo test exploring how the size of our 1ieasureimoent errors adfects he Kolmogorov-Sinirnoy coefficient p. which quantities he probability that the CN-stroug aud CN-avealk eroups were drawn from the same initial population.," To understand the effects of measurement error on our generalized histogram, we performed a Monte Carlo test exploring how the size of our measurement errors affects the Kolmogorov-Smirnov coefficient $p$, which quantifies the probability that the CN-strong and CN-weak groups were drawn from the same initial population."
+ We rescaled our true measurenmeut errors σς by factors ranging from 1.5 to LO. then selected a new value for cach data poiut roni within a Gaussian distribution centered at the neasured value of 85(3839). with a width of the rescaled uecasureineut error.," We rescaled our true measurement errors $\sigma_{S}$ by factors ranging from $0.5$ to $10$, then selected a new value for each data point from within a Gaussian distribution centered at the measured value of $\delta S(3839)$, with a width of the rescaled measurement error."
+ This resclection was done for the CN-normal aud CN-stroug eroups identified in Figure | independently., This reselection was done for the CN-normal and CN-strong groups identified in Figure \ref{fig4} independently.
+ The I&-S p coefficient was then measured as an dudicator of whether the two CN baud strenetl eroups containiug resclected poiuts with rescaled errors ποσο. likely to have been drawn from the same initial population., The K-S $p$ coefficient was then measured as an indicator of whether the two CN band strength groups containing reselected points with rescaled errors seemed likely to have been drawn from the same initial population.
+ Resclectious were doue LO! times for cach rescaling of the measurement error. and we find that p is strongly concentrated uear zero when the rescaling factor is relatively siuall: for a rescaling factor of 1. 99% of realizations have a p less than 0.1 and none have a p lavecr than 0.61. and 825€ of realizations with a rescaling factor of 3 have a p less than 0.5.," Reselections were done $10^{4}$ times for each rescaling of the measurement error, and we find that $p$ is strongly concentrated near zero when the rescaling factor is relatively small: for a rescaling factor of 1, $99\%$ of realizations have a $p$ less than $0.1$ and none have a $p$ larger than $0.61$, and $82\%$ of realizations with a rescaling factor of 3 have a $p$ less than $0.5$."
+ With a rescaling factor of 10. we fud that p values are fairly evenly distributed between 0 and 1. with 51% below p=0.5.," With a rescaling factor of 10, we find that $p$ values are fairly evenly distributed between 0 and 1, with $51\%$ below $p=0.5$."
+ Based on this test. we are fairly confident that our two CN eroups really are distinct. even if our errors in S(3839) are uuderestimated by a factor of 3.," Based on this test, we are fairly confident that our two CN groups really are distinct, even if our errors in $S(3839)$ are underestimated by a factor of 3."
+ Carbon abuudances were determined bw spectral svuthesis of the ας region., Carbon abundances were determined by spectral synthesis of the G-band region.
+ For cach star. a sot of svuthetic spectra were created with a singleTg.og(g).. and |Fe/TII]. aud a range of possible values for C/Foe]. sinoothed to the resolution of the VIRUS-P data.," For each star, a set of synthetic spectra were created with a single, and [Fe/H], and a range of possible values for [C/Fe], smoothed to the resolution of the VIRUS-P data."
+ The svuthetic spectra were ecucrated using a αποποσα version of MOOG 2009 (Sueden 1973)) which ecucrates Huxed spectra instead of normalized spectra., The synthetic spectra were generated using a modified version of MOOG 2009 (Sneden ) which generates fluxed spectra instead of normalized spectra.
+ We used he PLEZ2000 erid of metal-poor model atinospheres (Plez 2000))., We used the PLEZ2000 grid of metal-poor model atmospheres (Plez ).
+ The temperatures and surface eravitics were derived from the photometry in Table 1... using he prescriptions in Ramurez Moléuudez(2005).," The temperatures and surface gravities were derived from the photometry in Table \ref{t1}
+ using the prescriptions in rez Melénndez."
+.. We adopted [a /Fe| =0.3. 2O/C=6 and |Fe/TI| =2.0 for allstars.," We adopted $\alpha$ /Fe] $= 0.3$, $^{12}$ $/^{13}$ $ = 6$ and [Fe/H] $= -2.0$ for all stars."
+ The choice of the carbon isotope ratio does male a sinall difference to the derived total carbon abundance., The choice of the carbon isotope ratio does make a small difference to the derived total carbon abundance.
+ At these imetallicities. where the lines are reasonably weak. a change of C/O from 6 to 50 would male," At these metallicities, where the lines are reasonably weak, a change of $^{12}$ $/^{13}$ C from 6 to 50 would make"
+By using these frequency determinations. we can compute and plot large separations in Fig. 10..,"By using these frequency determinations, we can compute and plot large separations in Fig. \ref{fig:dnu}."
+" There are clear variations around the mean value Av=98.320.1Hz. recomputed as the average of Av;,,."," There are clear variations around the mean value $\Dnu=98.3\pm 0.1\muHz$, recomputed as the average of $\Delta\nu_{l,n}$."
+ The variation of the large separation can also be measured from the autocorrelation of the time series computed as the power spectrum of the power spectrum windowed with a narrow filter. as proposed by ? and ?..," The variation of the large separation can also be measured from the autocorrelation of the time series computed as the power spectrum of the power spectrum windowed with a narrow filter, as proposed by \citet{Roxburgh06} and \citet{Roxburgh09}."
+ ? have defined the envelope autocorrelation funetion (EACF) and explain its use as an automated pipeline., \citet{Mosser09ACF} have defined the envelope autocorrelation function (EACF) and explain its use as an automated pipeline.
+ We applied this pipeline to the spectrum of tto find the frequency variation Av(v) of the large separation., We applied this pipeline to the spectrum of to find the frequency variation $\Delta\nu(\nu)$ of the large separation.
+ We used a narrow cosine filter. with a full width at half maximum equal to 2Av.," We used a narrow cosine filter, with a full width at half maximum equal to ."
+. The result is plotted in Fig. 10.., The result is plotted in Fig. \ref{fig:dnu}..
+" The agreement between Ανν) and the fitted values Av;,, are quite good."," The agreement between $\Delta\nu(\nu)$ and the fitted values $\Delta\nu_{l,n}$ are quite good."
+" We recover the main variations. particularly. the large oscillation. visible as a bump around2100,4/Hz."," We recover the main variations, particularly the large oscillation, visible as a bump around."
+. Such oscillations in the frequencies and seismic variables can be created by sharp features in the stellar structure (e.g.??)..," Such oscillations in the frequencies and seismic variables can be created by sharp features in the stellar structure \citep[e.g.][]{Vorontsov88,Gough90}."
+ The observed oscillations are probably the signature of the helium’s second ionisation zone located below the surface of the star., The observed oscillations are probably the signature of the helium's second ionisation zone located below the surface of the star.
+ Such a signature. predicted by models and observed on the Sun. should help us to put constraints on the helium abundance of the star (e.g.?2)..," Such a signature, predicted by models and observed on the Sun, should help us to put constraints on the helium abundance of the star \citep[e.g.][]{Basu04,Piau05}."
+ We then computed small separations Results are shown in Fig. 11.., We then computed small separations Results are shown in Fig. \ref{fig:d02}. .
+ The error bars take the correlations between the frequency determination of the /20 and 2 modes into account., The error bars take the correlations between the frequency determination of the $l=0$ and 2 modes into account.
+ These correlations are small., These correlations are small.
+ It is worth noticing that these small separations do not decrease with frequency as for the Sun. but remain close to their mean value. 0o»=8.140.2wHz. re-evaluated from fitted frequencies.," It is worth noticing that these small separations do not decrease with frequency as for the Sun, but remain close to their mean value, $\overline{\delta_{02}}=8.1\pm 0.2\muHz$, re-evaluated from fitted frequencies."
+ Thesmall separations between /=0 and | modes. defined by? às are also interesting seismic diagnosis tools (e.g.2?)..," Thesmall separations between $l=0$ and 1 modes, defined by \citet{Roxburgh93} as are also interesting seismic diagnosis tools \citep[e.g.][]{Roxburgh03,Roxburgh05}."
+ Results for aare shown in Fig. 12., Results for are shown in Fig. \ref{fig:d01}.
+. As with the Oo» separations they do not display the decrease in frequency seen in the solar values. and have an average of 3.28+0.09Hz.," As with the $\delta_{02}$ separations they do not display the decrease in frequency seen in the solar values, and have an average of $3.28 \pm 0.09\muHz$."
+ These are additional diagnosties of the interior and may even be showing an oscillation due to the radius of an outer convective zone and the signature of a convective core (e.g.?2)..," These are additional diagnostics of the interior and may even be showing an oscillation due to the radius of an outer convective zone and the signature of a convective core \citep[e.g.][]{Roxburgh09_d01,Silva11}."
+" The results for the parameters v, and / obtained by the different groups are spread but consistent.", The results for the parameters $\nu_\mathrm{s}$ and $i$ obtained by the different groups are spread but consistent.
+ They appear to be sensitive to the prior for the background and the frequency range of the fit., They appear to be sensitive to the prior for the background and the frequency range of the fit.
+" A detailed study dedicated to determining v, and 7 will be presented in a separate paper (Gizon et al.."," A detailed study dedicated to determining $\nu_\mathrm{s}$ and $i$ will be presented in a separate paper (Gizon et al.,"
+ in preparation)., in preparation).
+" The results for the projected splitting v?=v,sinf. which ts easter to determine than v, and / separately (?).. are also consistent among the different groups."," The results for the projected splitting $\nu_\mathrm{s}^*=\nu_\mathrm{s} \sin i$, which is easier to determine than $\nu_\mathrm{s}$ and $i$ separately \citep{Ballot06}, are also consistent among the different groups."
+ For the analysis presented in this work we find v;=0.4540.10Hz for52265.," For the analysis presented in this work we find $\nu_\mathrm{s}^*=0.45 \pm 0.10
+\muHz$ for."
+". Moreover. ? show that the frequency estimates are not correlated with estimates for other parameters. in particular v, and i.so.the frequency estimates are robust."," Moreover, \citet{Ballot08} show that the frequency estimates are not correlated with estimates for other parameters, in particular $\nu_\mathrm{s}$ and $i$ ,so,the frequency estimates are robust."
+ To consider the window effects that affect p modes. we divideall the fitted mode heights by the factor rj... as mentioned in," To consider the window effects that affect p modes, we divideall the fitted mode heights by the factor , as mentioned in"
+where 34. 55 and Q are the Fourier transforms of sí). ο) and QU—P).,"where $\tilde{s}_1$ , $\tilde{s}_2$ and $\tilde{Q}$ are the Fourier transforms of $s_1(t)$, $s_2(t^\prime)$ and $Q(t-t^\prime)$."
+ an(zfTf)/zfisthe finite-time approximation to the Dirac delta funcüon., $\delta_T(f)=\sin(\pi fT)/\pi f$ is the finite-time approximation to the Dirac delta function.
+ If: we knew the signal h(/). we could construct a Q that maximizes the S/N as QUI)=nf)(SCS)SoC).," If we knew the signal $h(t)$, we could construct a $Q$ that maximizes the S/N as $\tilde{Q}(f)=|\tilde{h}(f)|/S_1(|f|)S_2(|f|)$."
+ Sif) and 53(f) are the power spectral densities of the noises of the two detectors., $S_1(f)$ and $S_2(f)$ are the power spectral densities of the noises of the two detectors.
+ Finn et al. (, Finn et al. (
+1999) suggested to adopt Chis filler with ACH? assumed to be unitv in the detector band if a detailed knowledge of h(/) is lacking.,1999) suggested to adopt this filter with $|\tilde{h}(f)|^2$ assumed to be unity in the detector band if a detailed knowledge of $h(t)$ is lacking.
+" Therefore. assuming lor simplicity that the power spectral densities are identical. 5(/)=54(f)Ss(f) in the two detectors. we adopt QUI)—A/S""(|f|) where \ is a normalization constant."," Therefore, assuming for simplicity that the power spectral densities are identical, $S(f)\equiv S_1(f)=S_2(f)$ in the two detectors, we adopt $\tilde{Q}(f)=\lambda/S^2(|f|)$ where $\lambda$ is a normalization constant."
+ The expected value of (he cross-correlation signal (averaged over the source population) ls Using data segments not associated with GRBs (oll-source). we can also evaluate the fluctuation of the cross-correlation of the noise. The factor 7 on the right hand side arises from evaluating ὁ(0).," The expected value of the cross-correlation signal (averaged over the source population) is Using data segments not associated with GRBs (off-source), we can also evaluate the fluctuation of the cross-correlation of the noise, The factor $T$ on the right hand side arises from evaluating $\delta_T(0)$."
+ We assume that the duration of gravitational wave bursts ave shorter than 7=10 sec., We assume that the duration of gravitational wave bursts are shorter than $T=10$ sec.
+ In the previous section. we assunied (hat the merger phase starts al 200 Iz (collapsars). or al /;~Q.41Iz and DIL-He binaries).," In the previous section, we assumed that the merger phase starts at $\sim 200$ Hz (collapsars), or at $f_i\sim 0.4$ Hz (BH-WD and BH-He binaries)."
+ During T=10 sec. the svstem can rotate multiple times. and ib is plausible (ο assume that a significant. fraction of the energy in the system is emitted during this time.," During $T=10$ sec, the system can rotate multiple times, and it is plausible to assume that a significant fraction of the energy in the system is emitted during this time."
+ If the energy spectrum of the gravitational waves dI/df is (lat. as we have assumed. most of gravitational wave energy is emitted in the high frequency band ~J).," If the energy spectrum of the gravitational waves $dE/df$ is flat, as we have assumed, most of gravitational wave energy is emitted in the high frequency band $\sim f_q$."
+ The emission timescale might be comparable to the damping time of the quasi-normal moce T (Flanagan IIughes 1998)., The emission timescale might be comparable to the damping time of the quasi-normal mode $\tau$ (Flanagan Hughes 1998).
+ Although 7 is much shorter than 10 sec. the lag between a GRB signal and the gravitational wave signal are model-dependent and rather uncertain.," Although $\tau$ is much shorter than 10 sec, the lag between a GRB signal and the gravitational wave signal are model-dependent and rather uncertain."
+ We assunie a conservative estimate of 7'=10 sec., We assume a conservative estimate of $T=10$ sec.
+" We deline the S/N of the gravitational wave signal in the merger phase as X,/0,,.", We define the S/N of the gravitational wave signal in the merger phase as $\overline{X}_{on}/\sigma_{off}$ .
+ For the nearest collapsar which occurs in à vear. the S/N ratio calculated numerically fromequations (21))(22)) is where the uncertainty range in p from the uncertainty range in (he rate 2 is 0.2—3.8 (rom Table 1). the blob masses enter throughthe reduced mass po=ΠΟ(ΗΕ ma). ancl we," For the nearest collapsar which occurs in a year, the S/N ratio calculated numerically fromequations \ref{eq:Xon}) \ref{eq:sigmaoff}) ) is where the uncertainty range in $\rho$ from the uncertainty range in the rate $R$ is $0.2-3.8$ (from Table 1), the blob masses enter throughthe reduced mass $\mu=m_1 m_2/(m_1+m_2)$ , and we"
+We thank KX. Stanek for providing an updated. compilation of the data points.,We thank K. Stanek for providing an updated compilation of the data points.
+ This work was supported in part by NASA through a IIubble Fellowship grant from the Space Telescope Science Institute. which is operated bv the Association of Universities for Research in Astronomy. Inc.. under NASA contract NAS5-26555 (for $.G.). bv the Horowitz foundation and US-Israel BSF grant BSF-9800225 (for J.G.) and the US-Israel BSF grant. DSE-9500343. NSF grant. AST-9900877. and NASA erant NAG5-7039 (for A.L.).," This work was supported in part by NASA through a Hubble Fellowship grant from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555 (for S.G.), by the Horowitz foundation and US-Israel BSF grant BSF-9800225 (for J.G.) and the US-Israel BSF grant BSF-9800343, NSF grant AST-9900877, and NASA grant NAG5-7039 (for A.L.)."
+a absorption systems in ITzRCs (Rottecringetal.199001: vanOjiketal.1997: Dev1999)}) we shall interpret the πιοατα asviuuetrv im the pprofile of aas being due to foreground absorption by neutral hydrogen.,$\alpha$ absorption systems in HzRGs \cite{rot95}; ; \cite{oji97}; \cite{dey99}) ) we shall interpret the blue-ward asymmetry in the profile of as being due to foreground absorption by neutral hydrogen.
+ The vestframe equivalent width of iu WUδέν910450AG is twice as lieh as in the wellstudied radio ealaxy LC 11.17 (2=3.80: Devetal. L997)).," The rest–frame equivalent width of in $W_{\lambda}^{rest}=210 \pm 50$, is twice as high as in the well--studied radio galaxy 4C 41.17 $z=3.80$; \cite{dey97}) )."
+" The large huninosity (Lis,~Ls104! ere lo after correction for absorption) makes tthe most luminous celuitting radio galaxy known.", The large luminosity $L_{\rm Ly\alpha} \sim 4\times 10^{44}$ erg $^{-1}$ after correction for absorption) makes the most luminous emitting radio galaxy known.
+" Following Spiuradοal.(1995)... we measure the continui discontinuitv across the line. defined as |CF,(1250.—1350A)/F,(1100. 1200À5| = 1.562:0.21."," Following \cite{spi95}, we measure the continuum discontinuity across the line, defined as $\langle F_{\nu}(1250 -
+1350{\rm \AA}) / F_{\nu}(1100 - 1200{\rm \AA}) \rangle$ ] = $1.56 \pm 0.24$."
+" Siuilavh. for the Lyman linüt at Apes =912À.. we find ο1000ÀA)/F,(850. 910A5| = 2.240.5. hough this value is uucertain because the flux calibration at the edge of the spectrum is poorly determined."," Similarly, for the Lyman limit at $\lambda_{rest}$ =, we find $\langle F_{\nu}(940 - 1000{\rm \AA}) / F_{\nu}(850 - 910{\rm \AA})
+\rangle$ ] = $2.2 \pm 0.5$, though this value is uncertain because the flux calibration at the edge of the spectrum is poorly determined."
+ The presence of these continu cdiscoutinuities rtler confini our nuüeasured redshift., The presence of these continuum discontinuities further confirm our measured redshift.
+ Towever. the redshift of the svstem is difficult to determine accurately vecatise our VLT spectrum does not cover oor1610.," However, the redshift of the system is difficult to determine accurately because our VLT spectrum does not cover or."
+. Furthermore. since the ecnissiou is heavily absorbed. it is likely that the redshift of the peak of the eenissjon (at 6206ΙΑ. 2=1.105+ 0.005) does not exactly coincide with the redshift of the galaxy.," Furthermore, since the emission is heavily absorbed, it is likely that the redshift of the peak of the emission (at $6206
+\pm 4 {\rm \AA}$, $z=4.105\pm 0.005$ ) does not exactly coincide with the redshift of the galaxy."
+ We shall assume 2=LAL., We shall assume $z = 4.11$.
+ sshares several properties in comuuon with other ITZRCis but some of its characteristics deserve special couuneut., shares several properties in common with other HzRGs but some of its characteristics deserve special comment.
+ Tere we shall bricfiv diseuss these., Here we shall briefly discuss these.
+" Assuming photoionization. case B recombination. and a temperature of T=1041 I sve use the observed cenussion to derive a total mass CAZ(IEID) of the eeas MeCurthyetal. 1990)) using ALCL)=10°οS4)? AML... Hero fs is the filline actor in units of 10°. Ly, is the Ihuuinosity in units of 104+ eres st and Tey is he total voluue in units of 10 cni."," Assuming photoionization, case B recombination, and a temperature of $T=10^4$ K we use the observed emission to derive a total mass $M(\HII)$ ) of the gas \cite{mcc90}) ) using $M(\HII)=10^9(f_{-5}L_{44}V_{70})^{1/2}$ $_{\odot}$ Here $f_{-5}$ is the filling factor in units of $^{-5}$, $L_{44}$ is the luminosity in units of $10^{44}$ ergs $^{-1}$, and $V_{70}$ is the total volume in units of $10^{70}$ $^3$."
+ Assuming a filline actor of 7 (MeCartliyetal. 19903). and a cubical volume with aside of 15 kpe. we find ALL)z2.5ς105 AL...," Assuming a filling factor of $^{-5}$ \cite{mcc90}) ), and a cubical volume with a side of 15 kpc, we find $M(\HII) \approx 2.5\times 10^8$ $_{\odot}$ ."
+ This value is ou the high side. but well within the rauge that has been found for ITzRCs vanOjiketal. 1997))).," This value is on the high side, but well within the range that has been found for HzRGs \cite{oji97}) ))."
+ Previous authors have shown hat easone clouds of stich mass ean cause radio jets to beud aud decollimate vanBreneclFilippeuxoIHeckuiau&Miley.1985.Lonsdale&Barthel1986.Miley:1988) ).," Previous authors have shown that gas clouds of such mass can cause radio jets to bend and decollimate \cite{wvb85,lon86,bar88}) )."
+" Likewise. the extreme asviunietrv in the rracdio source could well be the result of strong interaction between the radiocutting plasma aud the ous,"," Likewise, the extreme asymmetry in the radio source could well be the result of strong interaction between the radio–emitting plasma and the gas."
+OO;N Our spectymu also shows evidence for deep blue-wiud absorption of the eenission liue., Our spectrum also shows evidence for deep blue-ward absorption of the emission line.
+ We believe that this is probably due to resonant scattering bv cold eeas dn a halo surrounding the radio galaxy. as seen in αμα other TzRCs Rottecrimeetal.1995.. vanOjiketal. 1997.. Dev 1999)).," We believe that this is probably due to resonant scattering by cold gas in a halo surrounding the radio galaxy, as seen in many other HzRGs \cite{rot95}, \cite{oji97}, \cite{dey99}) )."
+ The spatial extent of the absorption edge as seen in the 2-dimensional spectrum (Fig., The spatial extent of the absorption edge as seen in the 2-dimensional spectrum (Fig.
+ 3) implies that the exteut of the absorbing eas is simular or even lavecr than the L’(30 kpc} ΟΗΕting region., 3) implies that the extent of the absorbing gas is similar or even larger than the (30 kpc) emitting region.
+ To coustrain the absorption parameters we constructed a simple model that describes the pprofile with a Gaussian cussion function aud a single Voigt absorption function., To constrain the absorption parameters we constructed a simple model that describes the profile with a Gaussian emission function and a single Voigt absorption function.
+ As a first step. we fitted the red wing of the cussion ine with a Caussian enuüssion profile.," As a first step, we fitted the red wing of the emission line with a Gaussian emission profile."
+Because the absorption is very broad.and exteuds to the red side of the peak. the parameters of thisGaussian,"Because the absorption is very broad,and extends to the red side of the peak, the parameters of thisGaussian"
+All seven o ‘the candidate BIB stars in tlis moving group had unmeasured metallicitics in SDSS DRG. but metallicities were assigned to these same spectra in the much improved SSPP pipeline of SDSS DRT (Leeetal.2008a.b)..,"All seven of the candidate BHB stars in this moving group had unmeasured metallicities in SDSS DR6, but metallicities were assigned to these same spectra in the much improved SSPP pipeline of SDSS DR7 \citep{SSPP1,SSPP2}."
+ Our metallicities come from the DIU. which came publie in October 20WS.," Our metallicities come from the DR7, which became public in October 2008."
+" Phe SDSS DI tabulates many measures of the metallicities of stars. of which the most commonly. used is FELLA. t1ο ""adopted? metallicity. which is derived. [rom a comparison of all methods. used to measure metallicity."," The SDSS DR7 tabulates many measures of the metallicities of stars, of which the most commonly used is FEHA, the “adopted"" metallicity, which is derived from a comparison of all methods used to measure metallicity."
+ In this paper we use instead the metallicity. of Wilhelm.Beers.&Gray (1999)... hereafter NBG. which is specifically designed. to measure the metallicities of BIB stars.," In this paper we use instead the metallicity of \citet{wbg99}, hereafter WBG, which is specifically designed to measure the metallicities of BHB stars."
+ The WBC metallicities of the seven DIID stars in the moving eroup are given in Table 1 and are shown in Figure 4., The WBG metallicities of the seven BHB stars in the moving group are given in Table 1 and are shown in Figure 4.
+ The mean of the distribution is. Fe/1I]—2.46+40.14. and the width iso=0.4.," The mean of the distribution is $=-2.46\pm0.14$, and the width is $\sigma=0.4$."
+ We now show that the metallicities of the moving group stars are not consistent with being drawn at random from he stellar halo., We now show that the metallicities of the moving group stars are not consistent with being drawn at random from the stellar halo.
+ We selected all stars in DRT with the same shotometric Constraints as the stars in the moving eroup. and with 645° so that the BIB stars are likely to come from tΊο same component (the spheroid). and: not »* confused with dise. populations of DBIIDs.," We selected all stars in DR7 with the same photometric constraints as the stars in the moving group, and with $b>45^\circ$ so that the BHB stars are likely to come from the same component (the spheroid), and not be confused with disc populations of BHBs."
+ To remove BIIB stars in the Sagittarius cwarl tidal stream from the sample. WO 0so eliminated. stars with &< 57.," To remove BHB stars in the Sagittarius dwarf tidal stream from the sample, we also eliminated stars with $\delta<5^\circ$ ."
+ There were 208 stars with 18.65<gy19.15. 0.8«(qru90.0. LS<(ug)o1.5. and that passed the photometric low surface gravity cut.," There were 298 stars with $18.65 < g_0 < 19.15$, $-0.3 < (g-r)_0 < 0.0$, $0.8 < (u-g)_0 < 1.5$, and that passed the photometric low surface gravity cut."
+ Of these. 260 had measured metallicities hat are shown in Figure 4.," Of these, 260 had measured metallicities that are shown in Figure 4."
+ The mean of the distribution is Feflj=169£0.04. and the width is &=0.71.," The mean of the distribution is $=-1.69\pm0.04$, and the width is $\sigma=0.71$."
+ ote hat the error quoted. here is à statistical error., Note that the error quoted here is a statistical error.
+ )ecause very few of the SEGUE calibration stars are as blue as AX stars. the systematic errors at this end are about 0.38.," Because very few of the SEGUE calibration stars are as blue as A stars, the systematic errors at this end are about $\pm 0.3$ ."
+ The comparison stars are about 45 kpe from the Sun. and above 6=45. so they come from the distant halo.," The comparison stars are about 45 kpc from the Sun, and above $b=45^\circ$, so they come from the distant halo."
+ Phe blue strageler contaminants are LO to 28 kpe away. which is also fairlyfar from the disc athigh Galactic LItitudes.," The blue straggler contaminants are 10 to 28 kpc away, which is also fairlyfar from the disc athigh Galactic latitudes."
+"When rj/r,= 0.15-0.20. clusters survive only if the SFE is high and the gas expulsion time-scale slow 3 when η~0.20:][]bau07.","When $r_h/r_t = 0.15$ $0.20$, clusters survive only if the SFE is high and the gas expulsion time-scale slow 3 when $r_h/r_t \simeq 0.20$."
+ This is because high SFE and slow gas expulsion dampen the spatial expansion of an exposed cluster. thereby increasing its volume density and resilience to an external tidal field.," This is because high SFE and slow gas expulsion dampen the spatial expansion of an exposed cluster, thereby increasing its volume density and resilience to an external tidal field."
+ We have emphasized that observationally inferred. cluster-forming core mass-radius diagrams must be handled with caution as these can be imprints of the molecular tracer used to map molecular cores (ee. CHO vs. CS: see Fig. 1).," We have emphasized that observationally inferred cluster-forming core mass-radius diagrams must be handled with caution as these can be imprints of the molecular tracer used to map molecular cores (e.g. $C^{18}O$ vs. $CS$; see Fig. \ref{fig:obs}) ),"
+ or of the method used to measure the core radius and the mass enclosed within that radiusDmainBodyCitationEnd9400]mue02., or of the method used to measure the core radius and the mass enclosed within that radius.
+ Apart from the cluster mass function evolution. how the mass-radius relation of cluster-forming cores couples to an external tidal field to produce a tidal field impact upon expanded gas-free clusters also affects the star formation history of galaxies as reconstructed from the age distribution of their surviving star clusters2007).," Apart from the cluster mass function evolution, how the mass-radius relation of cluster-forming cores couples to an external tidal field to produce a tidal field impact upon expanded gas-free clusters also affects the star formation history of galaxies as reconstructed from the age distribution of their surviving star clusters."
+". Because constant surface density cores lead to the preferential removal of the most massive clusters. the bound fraction of stars integrated over the whole core mass function Εδω) sensitively depends on adopted core mass function parameters when 6=1/2 (spectral index ϱβ and upper mass limit 77,5: see Table 2))."," Because constant surface density cores lead to the preferential removal of the most massive clusters, the bound fraction of stars integrated over the whole core mass function $\overline{F_{bound}}$ ) sensitively depends on adopted core mass function parameters when $\delta=1/2$ (spectral index $-\beta$ and upper mass limit $m_{up}$; see Table \ref{tbl:gal}) )."
+ Independently of the tidal field impact. we note that the cluster-forming core mass-radius relation is also relevant to the duration of cluster violent relaxation and to the slope of the mass-metallicity relation of old globular clusters as predicted by models (see Section 4)).," Independently of the tidal field impact, we note that the cluster-forming core mass-radius relation is also relevant to the duration of cluster violent relaxation and to the slope of the mass-metallicity relation of old globular clusters as predicted by self-enrichment models (see Section \ref{sec:conseq}) )."
+ G.P. is supported by a Research Fellowship of the Alexander von Humboldt Foundation., G.P. is supported by a Research Fellowship of the Alexander von Humboldt Foundation.
+protostellar jets and Chev predicted (hat the integrated II» line Iumninositv is greatest. in IRAC band 2 (4.5jm)). about an order of magnitude higher than those in other IRAC bands (see Table 3 in Smith&Rosen 2005)).,"protostellar jets and they predicted that the integrated $_{2}$ line luminosity is greatest in IRAC band 2 ), about an order of magnitude higher than those in other IRAC bands (see Table 3 in \citealt{sr05}) )."
+ Observations bySpitzer IRAC have in general confirmed this prediction (Noriega-Crespoetal.2004:Jorgensen 2006)...," Observations by IRAC have in general confirmed this prediction \citep{nor04,jor06}. ."
+ According to Noriega-Crespoetal.(2004).. the remaining emission in the$pizer IRAC band 2 image alter the subtraction of the emission in band 1 is mainly from shocked Hà and CO gas.," According to \citet{nor04}, the remaining emission in the IRAC band 2 image after the subtraction of the emission in band 1 is mainly from shocked $_{2}$ and CO gas."
+ Using the three-color images constructed from [RAC bands |. 2. and 4..," Using the three-color images constructed from IRAC bands 1, 2, and 4.,"
+ Cvganowski compiled a catalog of outflow candidates with theSpizer Galactic Legacy. Infrared. λα- Survey Extraordinaire (GLIMPSE) data., \citet{cyg08} compiled a catalog of outflow candidates with the Galactic Legacy Infrared Mid-Plane Survey Extraordinaire (GLIMPSE) data.
+ Thev identified the Extended Green Objects (EGOs) in the three-color images as outflow candidates., They identified the Extended Green Objects (EGOs) in the three-color images as outflow candidates.
+ In this paper. we use (he difference image constructed from [RAC channel 2 (4.5 jm)) minus channel 1 (3.6 jm)) to search for mid-infrared outflows.," In this paper, we use the difference image constructed from IRAC channel 2 (4.5 ) minus channel 1 (3.6 ) to search for mid-infrared outflows."
+ We use the three-color (3.6 iin blue. 4.5 iin green. 8.0 iin red) images to confirm that the identified outflows are EGQOs.," We use the three-color (3.6 in blue, 4.5 in green, 8.0 in red) images to confirm that the identified outflows are EGOs."
+ Images covering the common sky area of the p Ophiuchi cloud in the four ΗνΑς bands were trimmed from the downloaded.e2d archive mosaic images., Images covering the common sky area of the $\rho$ Ophiuchi cloud in the four IRAC bands were trimmed from the downloaded archive mosaic images.
+ The trimmed images were then rotated so that the image north direction is top aud the east is to the left., The trimmed images were then rotated so that the image north direction is top and the east is to the left.
+ For the subtraction of emission in band 1 from that in band 2. we selected 30 bright stars in the images that are roughly evenly distributed. over the sky and. are not saturated in either band 1 or band 2.," For the subtraction of emission in band 1 from that in band 2, we selected 80 bright stars in the images that are roughly evenly distributed over the sky and are not saturated in either band 1 or band 2."
+ The average flux ratio in band 1 to band 2 of these 80 stars is used to adjust the flux level of stars in band 1 to that in band 2., The average flux ratio in band 1 to band 2 of these 80 stars is used to adjust the flux level of stars in band 1 to that in band 2.
+ The resultant difference images were visually inspected to search for mid-infrared outflows., The resultant difference images were visually inspected to search for mid-infrared outflows.
+ We selected extended sources (hat have significant excess at aas mid-nfrared outflows., We selected extended sources that have significant excess at as mid-infrared outflows.
+ ‘Toward the core of L1633. several survevs for IHE objects and Is emission have been carried out (Wuetal.2002:GómezIXhanzadyanοἱal.2004) and 46 HII objects. including components of HII objects. and 119 II» emission features have been reported.," Toward the core of L1688, several surveys for HH objects and $_{2}$ emission have been carried out \citep{wu02,gom03,kha04} and 46 HH objects, including components of HH objects, and 119 $_{2}$ emission features have been reported."
+ We examined whether there are mid-inlrared counterparts of these LIT objects and Ils emission features at their corresponding positions in the IRAC images., We examined whether there are mid-infrared counterparts of these HH objects and $_{2}$ emission features at their corresponding positions in the IRAC images.
+ We discovered 13 [RAC mid-infrared outflows., We discovered 13 IRAC mid-infrared outflows.
+ Their coordinates and figure reference are list in Table 1.., Their coordinates and figure reference are list in Table \ref{tab1}. .
+ The likely exciting sources for these outflows and their Lada classes are eiven in Table 2.., The likely exciting sources for these outflows and their Lada classes are given in Table \ref{tab2}. .
+ Note that IAS. 16293-2422 is classified as a Class 0 source based on its bolometric temperature (Evansetal.2009)., Note that IRAS 16293-2422 is classified as a Class 0 source based on its bolometric temperature \citep{eva09}.
+. We also identilied 31 mid-infrared counterparts of previously known III objects and IH» 2.12 eemission features., We also identified 31 mid-infrared counterparts of previously known HH objects and $_{2}$ 2.12 emission features.
+ The coordinates of these mid-infraredcounterparts of thepreviouslyknown outflows. together with the corresponding," The coordinates of these mid-infraredcounterparts of thepreviouslyknown outflows, together with the corresponding"
+eiant of the M31. bulge.,giant of the M31 bulge.
+ ] have made a preliminary search [or other possible candidate remnants of M21 RY by registering frames (taken in the same bandpasses. and then visually blinking then as well as caleulating difference images.," I have made a preliminary search for other possible candidate remnants of M31 RV by registering frames taken in the same bandpasses, and then visually blinking them as well as calculating difference images."
+ This was done using the live FSI4W frames (1999 to 2010). and separately using the three F555W frames (1999 to 2009). as listed in Table 1.," This was done using the five F814W frames (1999 to 2010), and separately using the three F555W frames (1999 to 2009), as listed in Table 1."
+ Although {here are manv variable stars in these frames. none of them lie within the M31 RV error box.," Although there are many variable stars in these frames, none of them lie within the M31 RV error box."
+ For well-exposed stars. variability in excess of ~0.2 mag would have been seen.," For well-exposed stars, variability in excess of $\sim$ 0.2 mag would have been seen."
+ A separate search for variable objects at. fainter levels will be reported later. but (here are no objects among (he numerous [ainter sources that appear to have faded dramatically or disappeared.," A separate search for variable objects at fainter levels will be reported later, but there are no objects among the numerous fainter sources that appear to have faded dramatically or disappeared."
+ The M31 RV site has been imaged once hy in the near-IR. using the IR channel of WFCS3 and the broad-band FLIOW ancl F160W filters.," The M31 RV site has been imaged once by in the near-IR, using the IR channel of WFC3 and the broad-band F110W and F160W filters."
+ I have compared (hese [rames visually with the optical (F555W and ES14W) images. and found no unusually red objects.," I have compared these frames visually with the optical (F555W and F814W) images, and found no unusually red objects."
+ Thus. if the remnant of M31 RV is now embedded in dust. the obscuration is such that the source is not detectable out to the 77 band.," Thus, if the remnant of M31 RV is now embedded in dust, the obscuration is such that the source is not detectable out to the $H$ band."
+" l have investigated all available images of the site of M31 RV. an ILRT event that occurred in \I31 in 1055,"," I have investigated all available images of the site of M31 RV, an ILRT event that occurred in M31 in 1988."
+ These Irames were taken between 1995 and 2010., These frames were taken between 1995 and 2010.
+ I have been unable to verily (he reports by SZPYN1O of (a) a UV-bright object at the outburst site in 1995 (the apparent source is due (o an unfortunate case of cosmic ravs striking the same pixel in two different images): (b) a hot optical source in 1999 and 2008 (due to a misinterpretation of the SImagnitude system): and (¢) variability of this candidate source., I have been unable to verify the reports by SZPYK10 of (a) a UV-bright object at the outburst site in 1995 (the apparent source is due to an unfortunate case of cosmic rays striking the same pixel in two different images); (b) a hot optical source in 1999 and 2008 (due to a misinterpretation of the STmagnitude system); and (c) variability of this candidate source.
+ Thus the conclusions we reached in Paper I all remain valid., Thus the conclusions we reached in Paper I all remain valid.
+ I quote them verbatim: ~All ol the stars at the outburst site have the magnitudes and colors of ordinary red giants in the AI31 bulge., I quote them verbatim: “All of the stars at the outburst site have the magnitudes and colors of ordinary red giants in the M31 bulge.
+ The absence of any. conspicuous remnant star has three possible explanations:, The absence of any conspicuous remnant star has three possible explanations:
+jet-based variability but also calibration errors. ancl is a strict but important criterion to ensure that poor quality measurements do not overwhehnu the leusimg signal.,"jet-based variability but also calibration errors, and is a strict but important criterion to ensure that poor quality measurements do not overwhelm the lensing signal."
+ These 3878 quasars have. in total. 230.671 QUEST measurement pairs.," These 3878 quasars have, in total, 230,674 Palomar-QUEST measurement pairs."
+ We only use measurement pairs for which: Cut 1 eliminates a long but low-level tail out to ]lueer time lags. aud throws away 19.523 pairs.," We only use measurement pairs for which: Cut 1 eliminates a long but low-level tail out to larger time lags, and throws away 19,523 pairs."
+ Cut 2 js necessary to avoid the calculation of imaginary V values for individual measurement pairs and is very restrictive because the Palomar-QUEST systematic error is of the order of tvpical quasar variability amplitudes over rest frame timescales of weeks., Cut 2 is necessary to avoid the calculation of imaginary $V$ values for individual measurement pairs and is very restrictive because the Palomar-QUEST systematic error is of the order of typical quasar variability amplitudes over rest frame timescales of weeks.
+ Stil as will be seen in figure 2.. the data after this cut continue to follow a variabilityv-luuinosity relation described hy equation 2.. and therefore cau be used to measure iuaesnifcation using equation 6..," Still, as will be seen in figure \ref{v_vs_l2500}, the data after this cut continue to follow a variability-luminosity relation described by equation \ref{v_vs_l_equation}, and therefore can be used to measure magnification using equation \ref{mu_equation}."
+ We also note that removing the lowest signal-to-noise V values reduces error due to imprecise determination of the survey's systematic uncertaiuties. eiven the dependence of V ou the measurement uncertainty as stated in equation 3.," We also note that removing the lowest signal-to-noise $V$ values reduces error due to imprecise determination of the survey's systematic uncertainties, given the dependence of $V$ on the measurement uncertainty as stated in equation \ref{v_equation}."
+ Because the systematic uncertaiutics steni iu part from imuperfect flat field) aud fringe calibration inages. they can be time ancl position dependent aud therefore diffienlt to determine precisely for each measurement.," Because the systematic uncertainties stem in part from imperfect flat field and fringe calibration images, they can be time and position dependent and therefore difficult to determine precisely for each measurement."
+ This cut eliminates 131.715 pairs.," This cut eliminates 134,745 pairs."
+ Cut 3 is a final check against outlicrs that are nucharacteristic of quasar behavior. and only throws away 17 ieasurement pairs.," Cut 3 is a final check against outliers that are uncharacteristic of quasar behavior, and only throws away 17 measurement pairs."
+ Cut d serves to keep α Μποαν of exceptionally wellineasured quasars from dominating the dataset. since the ΠΡΟ of lcasurement pairs rises as the number of measurements squared.," Cut 4 serves to keep a minority of exceptionally well-measured quasars from dominating the dataset, since the number of measurement pairs rises as the number of measurements squared."
+ This cut removes 11.671 pairs.," This cut removes 14,671 pairs."
+ Normalization of the quasars. described iu section 3.. is only done if there are sufficient. quasars with similar properties so that we can determine accurate zero points.," Normalization of the quasars, described in section \ref{mu_meas_section}, is only done if there are sufficient quasars with similar properties so that we can determine accurate zero points."
+ We require 50 measurement pairs per 7 AL 2 EL bin. which removes: 1359 measurement pairs.," We require 50 measurement pairs per $\tau$ $M$ $z$ $L$ bin, which removes 4359 measurement pairs."
+" We are left with 57.359 useable measurement pairs. from, 3573 quasars."," We are left with 57,359 useable measurement pairs, from 3573 quasars."
+" We note that many of our cuts are due to properties of the data such as statistics and measurement errors, rather than properties of the quasars."," We note that many of our cuts are due to properties of the data such as statistics and measurement errors, rather than properties of the quasars."
+ Cuts on future quasar data sets may not necd to be so restrictive., Cuts on future quasar data sets may not need to be so restrictive.
+ After normalizing the data as described iu section 3.. we plot the normalized variability vs. Iuuinosity relation. which is shown in figure 2..," After normalizing the data as described in section \ref{mu_meas_section}, we plot the normalized variability vs. luminosity relation, which is shown in figure \ref{v_vs_l2500}."
+ The y axis error bars on the points are errors on the mean calculated from the spread of the measurements im each bin., The $y$ axis error bars on the points are errors on the mean calculated from the spread of the measurements in each bin.
+ The bius in the figure are smaller than those used in the uormalization. in order to show the relation iu more detail.," The bins in the figure are smaller than those used in the normalization, in order to show the relation in more detail."
+" The trend is indeed wellkdescribed by a power law. with index 0.565£0,007."," The trend is indeed well-described by a power law, with index $\alpha = 0.565 \pm 0.007$ ."
+ The best fit line is shown on the plot., The best fit line is shown on the plot.
+" This slope is steeper than that found in other works (ο,δν, ?.. 7): however. it is not directly comparable since we only consider measurements for which we observe significant variability."," This slope is steeper than that found in other works (e.g., \citealt{vandenberk04}, \citealt{bauer09a}) ); however, it is not directly comparable since we only consider measurements for which we observe significant variability."
+ The precision of our niaguification niecasuremoents relies on a harrow spread iu the variabilitv-Iuninosity relation., The precision of our magnification measurements relies on a narrow spread in the variability-luminosity relation.
+" The residuals iu los(V) around the best-fit line are distributed normally, with a standard deviation 0=0.22 that is coustant across our ranee of luuinosity."," The residuals in $V$ ) around the best-fit line are distributed normally, with a standard deviation $\sigma = 0.22$ that is constant across our range of luminosity."
+ These residuals are larger than the typical imieasurement errors. and therefore reflect an inherent spread iu the quasar variability properties.," These residuals are larger than the typical measurement errors, and therefore reflect an inherent spread in the quasar variability properties."
+ There are several possible sources of the enhanced variance., There are several possible sources of the enhanced variance.
+ The aeasurement error on the luminosity is siuiall. and is based simply on the uncertainties iu the SDSS broadband magnitudes.," The measurement error on the luminosity is small, and is based simply on the uncertainties in the SDSS broadband magnitudes."
+ This error estimate ignores the fact that the quasars vary bv more than these uncertainties: the true mean luuinosity could be tens of percent different from what was measured when the SDSS inade its photometric observation., This error estimate ignores the fact that the quasars vary by more than these uncertainties; the true mean luminosity could be tens of percent different from what was measured when the SDSS made its photometric observation.
+ Since the spectroscopic quasar sunuple is much brighter than the SDSS photometric detection lait. there should be no Malhuquist bias in the quasar sample aud the eror on the measured luminosity should be nubiased.," Since the spectroscopic quasar sample is much brighter than the SDSS photometric detection limit, there should be no Malmquist bias in the quasar sample and the error on the measured luminosity should be unbiased."
+ Therefore. when discussing the variance of the relation. we assiuue that the huninosity is correct and simply absorb its error into the y axis spread in values.," Therefore, when discussing the variance of the relation, we assume that the luminosity is correct and simply absorb its error into the $y$ axis spread in values."
+ Iu future studies this uncertainty can be reduced in cases where there are more photometric measurements per quasar. which can be averaged to produce a more accurate typical Iuninositv. (," In future studies this uncertainty can be reduced in cases where there are more photometric measurements per quasar, which can be averaged to produce a more accurate typical luminosity. ("
+This is not possible with the Palomar-QUEST data. as it is relatively calibrated rather than absolutely.),"This is not possible with the Palomar-QUEST data, as it is relatively calibrated rather than absolutely.)"
+ Or. if any nieasurenments are available. a lighteurve iiay be fit to cach quasars data to determine the object’s baseline," Or, if many measurements are available, a lightcurve may be fit to each quasar's data to determine the object's baseline"
+In this section we use N-bods simulations to find optimal nmiass-estinmation weights for halos binned by mass (rather than. e.g.. angular momentum. axis ratio. formation time or concentration).,"In this section we use N-body simulations to find optimal mass-estimation weights for halos binned by mass (rather than, e.g., angular momentum, axis ratio, formation time or concentration)."
+ We show the stochasticity that results from applving the optimal weights and compare to common optimal choices., We show the stochasticity that results from applying the optimal weights and compare to common sub-optimal choices.
+ For purposes of exploration. we wish to have a wide range of halo masses. while still having σους statistics for large halos and L[arge-scale modes.," For purposes of exploration, we wish to have a wide range of halo masses, while still having good statistics for large halos and large-scale modes."
+ For the first. purpose. we use the Millennium: simulation (?).. which resolves halos down to ~LOMTALL.," For the first purpose, we use the Millennium simulation \citep{Springel05b}, which resolves halos down to $\sim 10^{10} h^{-1}M_{\odot}$."
+ For the second. we use the suite of 49 cosmological dark matter 7NYU simulations described in. 2).. which have a total volume SOO) larger than the Millennium. but the minimum resolved halo mass is 1000. larger.," For the second, we use the suite of 49 cosmological dark matter “NYU” simulations described in \citet{Manera10}, which have a total volume $800\times$ larger than the Millennium, but the minimum resolved halo mass is $1000\times$ larger."
+ Phe simulations assume very similar AC'DAL cosmological models and. were carried. on using the same code (2).., The simulations assume very similar $\Lambda$ CDM cosmological models and were carried on using the same code \citep{Springel05a}.
+ Fable 1 provides details of the basic simulation parameters., Table \ref{tab1} provides details of the basic simulation parameters.
+ The initial power spectrum was generated by (2)..., The initial power spectrum was generated by \citep{Seljak96}.
+ Ehe initial density field— of the Millennium.— simulation was realized. by perturbing a homogeneous. glass-like particle distribution with a Caussian random realization of the initial power spectrum (?).. while the NYU simulations use an algorithm motivated. by Second. Order Lagrangian Perturbation Vheory (2)..," The initial density field of the Millennium simulation was realized by perturbing a homogeneous, glass-like particle distribution with a Gaussian random realization of the initial power spectrum \citep{white96}, while the NYU simulations use an algorithm motivated by Second Order Lagrangian Perturbation Theory \citep{Scoccimarro98}."
+ Dark matter halos with at least 20 particles are identified in both simulations using the friends-of-friends. (POL) group finder. with a linking length of 0.2 times the mean particle separation (?).., Dark matter halos with at least 20 particles are identified in both simulations using the friends-of-friends (FOF) group finder with a linking length of $0.2$ times the mean particle separation \citep{Davis85}.
+ The lowest mass halo of the two simulations that we use are ντ10tal. andLO«10775.TAL. respectively., The lowest mass halo of the two simulations that we use are $1.7\times10^{10}h^{-1}M_{\odot}$ and$1.0\times10^{13}h^{-1}M_{\odot}$ respectively.
+" The Millennium simulation has zLS.10' halos at at —0.220.5 and z=1 while the NYU simulations have ~43105. 3.5e10"" and 2.510° halos in cach of these redshift slices."," The Millennium simulation has $\approx 1.8\times10^7$ halos at at $z=0$, $z=0.5$ and $z=1$ while the NYU simulations have $\sim 4.3\times10^7$, $3.5\times10^7$ and $2.5\times10^7$ halos in each of these redshift slices."
+ For a more detailed analysis of the halo mass function and bias factors in these simulations. see 7). and 2]," For a more detailed analysis of the halo mass function and bias factors in these simulations, see \citet{Springel05b} and \citet{Manera10}."
+ We start our measurements by dividing halos into bins sorted. by mass., We start our measurements by dividing halos into bins sorted by mass.
+ The clustering of halos depends: weakly on the halo mass when halo mass is low but increases rapidly with mass when Al10bM..., The clustering of halos depends weakly on the halo mass when halo mass is low but increases rapidly with mass when $M>10^{13}~h^{-1}M_{\odot}$.
+ Furthermore halo abundances drop sharply at high mass., Furthermore halo abundances drop sharply at high mass.
+ As we do not want a wide range of masses within a single bin. we include fewer halos per bin at high masses.," As we do not want a wide range of masses within a single bin, we include fewer halos per bin at high masses."
+ We choose bins so that the number of halos in each decreases exponentially at high masses., We choose bins so that the number of halos in each decreases exponentially at high masses.
+ The optimal weighting and stochasticity are robust to changes in. the binning. as long as the function (AL) is well sampled.," The optimal weighting and stochasticity are robust to changes in the binning, as long as the function $b(M)$ is well sampled."
+ We divide the halos into 10 bins for the NYU simulations. anc use up to 30 bins for the Millennium simulation.," We divide the halos into 10 bins for the NYU simulations, and use up to 30 bins for the Millennium simulation."
+ We have tested that using more bins does not change our results., We have tested that using more bins does not change our results.
+ Within each halo bin. we weight halos by their masses and assign them to a 7=256° 3D mesh of cubic grid cells using the clouc-in-cell (CIC) assignment scheme (7). Le. we take the Fourier transform of the mass distribution within a halo bin.," Within each halo bin, we weight halos by their masses and assign them to a $N_g^3=256^3$ 3D mesh of cubic grid cells using the cloud-in-cell (CIC) assignment scheme \citep{Hockney81}, i.e. we take the Fourier transform of the mass distribution within a halo bin."
+ This is in anticipation of the result below. that optimal weighting is closer to masseweighting than number-weighting of halos., This is in anticipation of the result below that optimal weighting is closer to mass-weighting than number-weighting of halos.
+ Ifthe bins are narrow in mass. this choice oLintra-bin weighting should have little ellect. which we have verified.," If the bins are narrow in mass, this choice of intra-bin weighting should have little effect, which we have verified."
+ We separately Fourier transform the overdensity. Ποια of each halo bin and the total mass distribution., We separately Fourier transform the overdensity field of each halo bin and the total mass distribution.
+" We correct cach Fourier mode for the convolution with the CIC window function by the operation: where Lr.g.2)=Los/2NNqοο”. and UN, is the number of grid cells in cach dimension."," We correct each Fourier mode for the convolution with the CIC window function by the operation: where $\{x, y, z\}=\{k_xL_{box}/2N_g, k_yL_{box}/2N_g, k_zL_{box}/2N_g\}$, and $N_g$ is the number of grid cells in each dimension."
+" For each bin in & we construct the covariance matrix of Fourier coellicients C5;(&)=(δι)(Αν where i£ and j range over all halo bins. as well as the mass power £=(57,5. and hence the covariance biases oof the halos against the mass."," For each bin in $k$ we construct the covariance matrix of Fourier coefficients $C_{ij}(k)=\langle\delta_i(k)\,\delta_j(k)\rangle$, where $i$ and $j$ range over all halo bins, as well as the mass power $P=\langle \delta_m^2\rangle$, and hence the covariance biases of the halos against the mass."
+ For the NYU simulations we average results from all 49 realizations to produce a mean covariance mnatrix., For the NYU simulations we average results from all 49 realizations to produce a mean covariance matrix.
+ Figures 1. shows the optimal weights we derive from the simulations (black solid. lines). ancl compares them. with various functionsof mass.," Figures \ref{wvsm} shows the optimal weights we derive from the simulations (black solid lines), and compares them with various functionsof mass."
+ Purple-dashed lines show wxm: blue curves show the bias weighting ax6 that would be appropriate if the stanclarel biasec-Poisson mocel were correct. and dotted: curves show the optimal weight in equation (56)) derived from the halo model of Section 2.5..," Purple-dashed lines show $w\propto m$; blue curves show the bias weighting $w\propto b$ that would be appropriate if the standard biased-Poisson model were correct, and dotted curves show the optimal weight in equation \ref{halomodelw}) ) derived from the halo model of Section \ref{halomodel}."
+ Clearly. neither AL nor b are optimal weights.," Clearly, neither $M$ nor $b$ are optimal weights."
+ Indeed.the shape of eCM) depends on the cut-olf mass Maui of the halo catalogue. (," Indeed,the shape of $w_{\rm opt}(M)$ depends on the cut-off mass $M_{\rm min}$ of the halo catalogue. ("
+"This dependence follows that found by 1110. in their study of 0,.)","This dependence follows that found by H10, in their study of $\sigma_w$ .)"
+ When Miiorh ΤΑ. the massive end of the way(M) is close to mass weighting. as illustrated by the right-hand. plot of Figures 1..," When $M_{\rm min}<10^{13}~h^{-1}M_{\odot}$ , the massive end of the $w_{\rm opt}(M)$ is close to mass weighting, as illustrated by the right-hand plot of Figures \ref{wvsm}. ."
+ When AM is close to Adin. however. wap(Al) is Datter than mass weighting.," When $M$ is close to $M_{\rm min}$ , however, $w_{\rm opt}(M)$ is flatter than mass weighting."
+ Moreover. the slope of ayCM) gets shallower as," Moreover, the slope of $w_{\rm opt}(M)$ gets shallower as"
+a limited region in which the disk temperature can actually increase with radius - a racial temperature inversion.,a limited region in which the disk temperature can actually increase with radius - a radial temperature inversion.
+ In (his Letter. we adopt the disk model developed in I1P10 (see Table. 1))," In this Letter, we adopt the disk model developed in HP10 (see Table. \ref{param_disk}) )"
+ and focus our computations on M clwarl svstems. such as the recently discovered Super-Earth (~5.: AL is the Earth's mass) with an orbital radius of 2 AU 2006).," and focus our computations on M dwarf systems, such as the recently discovered Super-Earth $\sim 5 M_{\oplus }$; $M_{\oplus }$ is the Earth's mass) with an orbital radius of 2 AU \citep{bbfw06}."
+. We refer the readers to IIPIO for the detail., We refer the readers to HP10 for the detail.
+ The low mass of disks around M stars allows much more comprehensive Monte Carlo simulations to be performed. but our analvsis in principle applies to any. protoplanetary disk.," The low mass of disks around M stars allows much more comprehensive Monte Carlo simulations to be performed, but our analysis in principle applies to any protoplanetary disk."
+ We adopt the torque formula (Ward1997:Menou&Goodman2004:Jang-CondellSasselov2005) in which only the Lindblad torque is considered.," We adopt the torque formula \citep{ward97,mg04,js05} in which only the Lindblad torque is considered."
+ This is because the corotation torque is readily saturated (i.e. is canceled out) in dead zones in our radiativelv heated disk model., This is because the corotation torque is readily saturated (i.e. is canceled out) in dead zones in our radiatively heated disk model.
+" We compare the libration Uimescale (ie the timescale Lor gas to complete an orbit in the horseshoe region). 75528zr5/2,7.. with the viscous timescale, TrigP3iT /3y. where the half-width of the horseshoe region is rr,=ΓΗ"".3 (Paardekooper&Papaloizou2009).. (he kinematic viscosity is v=aLPOKey Sunvaey 1973).. and the disk scale height is ."," We compare the libration timescale (ie the timescale for gas to complete an orbit in the horseshoe region), $\tau_{lib}\approx 8\pi r_p/3\Omega_p x_s$, with the viscous timescale, $\tau_{vis}\approx x_s^2/3\nu$ , where the half-width of the horseshoe region is $x_s/r_p=1.68(M_p r_p/M_* h_p)^{1/2}$ \citep{pp09}, the kinematic viscosity is $\nu=\alpha h^2 \Omega_{Kep}$ \citep{ss73}, and the disk scale height is $h$."
+" Our numerical results give νο0.05. so that the critical value of turbulence. 0,5. below which the corotation torque is saturated (and therefore negligible). is Since the dead zone has a=10. we can salely neglect the corotation torque for AL,Z0.5. in the dead zone of our disk model."," Our numerical results give $h_p/r_p\simeq 0.05$, so that the critical value of turbulence, $\alpha_{crit}$ below which the corotation torque is saturated (and therefore negligible), is Since the dead zone has $\alpha=10^{-5}$, we can safely neglect the corotation torque for $M_p \gtrsim 0.5 M_{\oplus}$ in the dead zone of our disk model."
+ In addition. we confirmed that. in the aclive region where the corotation torque is generally. unsaturated. both Lindblad. aud corotation torques result in inward migration in our disk model (Paarcekooperetal.2009).," In addition, we confirmed that, in the active region where the corotation torque is generally unsaturated, both Lindblad and corotation torques result in inward migration in our disk model \citep{pbck09}."
+. Thus. exclusion of the corotation torque in the active region does not affect our findings in our disk model.," Thus, exclusion of the corotation torque in the active region does not affect our findings in our disk model."
+ We note that corotation torque max be unsaturated in dead zones [or sulliciently small planetary masses.bul the exactlimit will depends on knowing the disk," We note that corotation torque may be unsaturated in dead zones for sufficiently small planetary masses,but the exactlimit will depends on knowing the disk"
+Garaud model. initially.,"Garaud model, initially."
+ We note that the Garaud expression steepens quickly for rp213 cnm. suggesting that the growth rates of the (wo approaches are more comparable al later limes.," We note that the Garaud expression steepens quickly for $r_L \gtrsim 13$ cm, suggesting that the growth rates of the two approaches are more comparable at later times."
+ However. (he minor “kink” in the long-dashed curve is due to the shift fom Epstein to Stokes flow.," However, the minor “kink” in the long-dashed curve is due to the shift from Epstein to Stokes flow."
+ A much more subcdued kink is visible in the explicit approach which uses the full expressions for the turbulent velocity. whereas in the derivation of Eq. (," A much more subdued kink is visible in the explicit approach which uses the full expressions for the turbulent velocity, whereas in the derivation of Eq. ("
+29). it is assumed that the stopping time in (he Stokes regine is defined by some mean characteristic velocity which leads to à much more noticeable discontinuity.,"29), it is assumed that the stopping time in the Stokes regime is defined by some mean characteristic velocity which leads to a much more noticeable discontinuity."
+ Regardless. the overall more subdued erowth rate elicited by Eq. (," Regardless, the overall more subdued growth rate elicited by Eq. ("
+29) (despite the steepening at later times due (o a shift in flow regimes) is apparent [rom the lower set of curves in Fig.,29) (despite the steepening at later times due to a shift in flow regimes) is apparent from the lower set of curves in Fig.
+" 5 where (he initial conditions were chosen with r,(0)=1 cm.", 5 where the initial conditions were chosen with $r_L(0) = 1$ cm.
+ For this case. growth occurs only in the Epstein regime. but still. the curves for (he explicit approach aud (hat caleulated from Eq. (," For this case, growth occurs only in the Epstein regime, but still, the curves for the explicit approach and that calculated from Eq. ("
+29) begin to diverge.,29) begin to diverge.
+ Thus. the Garaud expression apparently underestimates the growth rate relative to our approach.," Thus, the Garaud expression apparently underestimates the growth rate relative to our approach."
+" We empliasize that the treatment of particle growth by Garaud(2007) requires several approximations in order to derive a purely analvlical expressions lor dr,/dl.", We emphasize that the treatment of particle growth by \citet{gar07} requires several approximations in order to derive a purely analytical expressions for $dr_L/dt$.
+ Besicles the alorementioned restriction of q=11/6. Garaucl approximates (he full expressions for the turbulent velocities that we use here by partitioning ey into seperate cases dependent on the particles’ stopping time relative to the turnover times of the smallest and largest scale eddies.," Besides the aforementioned restriction of $q=11/6$, Garaud approximates the full expressions for the turbulent velocities that we use here by partitioning $v_T$ into seperate cases dependent on the particles' stopping time relative to the turnover times of the smallest and largest scale eddies."
+ Furthermore. some question may be raised as (o the comparability of the moment equations used here to derive Eq. (," Furthermore, some question may be raised as to the comparability of the moment equations used here to derive Eq. ("
+23). versus the particle growth equation used by 29)..,"28), versus the particle growth equation used by \citet[her Eq. 29]{gar07}."
+ The equation used bx Garaud(2007) is more akin to a “sweep-up” equation. with no sources or sinks. (han to a formal coagulation equation.," The equation used by \citet{gar07}
+ is more akin to a “sweep-up” equation, with no sources or sinks, than to a formal coagulation equation."
+ Because it bears some resemblance (to the equation for M» (wilh some algebra. to [actors of order unity). we concluded that our Eq. (," Because it bears some resemblance to the equation for $M_2$ (with some algebra, to factors of order unity), we concluded that our Eq. ("
+28) is Che appropriate analog.,28) is the appropriate analog.
+ Despite the differences in growth rate. we find the agreement in the general trend of growth of the (vo approaches reassuring.," Despite the differences in growth rate, we find the agreement in the general trend of growth of the two approaches reassuring."
+ Evolutionary models of protoplanetary nebulae. giant planet atmospheres. etc.," Evolutionary models of protoplanetary nebulae, giant planet atmospheres, etc."
+ must somehow (real the escape of thermal radiation (Pollackοἱal.1996:IIubickv]Durisenetal. 2007).," must somehow treat the escape of thermal radiation \citep{pol96,hub05,dur07}."
+. In most cases. particles provide the primary opacity lor these moclels.," In most cases, particles provide the primary opacity for these models."
+ Observations of these ancl similar objects often rely on Spectral Energy. Distributions (SEDs) which can be compared (ο a model once the models internal temperature distribution is known: clear evidence is seen for grain growth in many cases (seereviewbyNatta 2007)., Observations of these and similar objects often rely on Spectral Energy Distributions (SEDs) which can be compared to a model once the model's internal temperature distribution is known; clear evidence is seen for grain growth in many cases \citep[see review by][]{nat07}.
+. Because of the nearly instwmountable computational burden involved with »erforminege a fully sell-consistent calculation of particle egrowth by coagulationc» alonge with an already difficult Hid. dvnamical calculation. most modelers simply assume some invariant," Because of the nearly insurmountable computational burden involved with performing a fully self-consistent calculation of particle growth by coagulation along with an already difficult fluid dynamical calculation, most modelers simply assume some invariant"
+lowering the aye parameter. which controls the effective strength’ of the artificial viscosity and is proportional to the excess pressure assigned to each eas particle in the equation of motion.,"lowering the $\alpha_{\rm visc}$ parameter, which controls the effective `strength' of the artificial viscosity and is proportional to the excess pressure assigned to each gas particle in the equation of motion."
+ Lowering the value of avi; from the default value of 0.8 has no significant consequences for the resulting bound mass of gas., Lowering the value of $\alpha_{\rm visc}$ from the default value of 0.8 has no significant consequences for the resulting bound mass of gas.
+ This is the case even when the artificial viscosity is seb τοzero?., This is the case even when the artificial viscosity is set to.
+. We conclude that our ram pressure results are robust to our choice of resolution and artificial viscosity strength., We conclude that our ram pressure results are robust to our choice of resolution and artificial viscosity strength.
+ Lt should be noted. however. that ram pressure Is not the only mechanism bv which gas can be stripped from ealaxies as hey orbit about groups and clusters.," It should be noted, however, that ram pressure is not the only mechanism by which gas can be stripped from galaxies as they orbit about groups and clusters."
+ In particular. Ixelvin-Ieblmholtz (sll) and Bavleigh-Tavlor (ICE) instabilities can potentially develop at the interface between the hot halo of he galaxy ancl the ICM and eventually completely disrupt or destroy the gaseous halo of the galaxy.," In particular, Kelvin-Helmholtz (KH) and Rayleigh-Taylor (RT) instabilities can potentially develop at the interface between the hot halo of the galaxy and the ICM and eventually completely disrupt or destroy the gaseous halo of the galaxy."
+ ἵς known that SPILL simulations tend to suppress such instabilities in the orescence of large density gradients across the interface., It is known that SPH simulations tend to suppress such instabilities in the presence of large density gradients across the interface.
+ This. in turn. will make the hot halo of à galaxy more resilient o stripping than it otherwise would. have been.," This, in turn, will make the hot halo of a galaxy more resilient to stripping than it otherwise would have been."
+ | good example of this can be found in Agertz ct ((2007). where a comparison between several Eulerian grid-based: codes (which accurately. follow the growth of these instabilities) and several Lagrangian SPII codes is performed. for an idealised case where a blob. of gas moves through a uniformi density medium.," A good example of this can be found in Agertz et (2007), where a comparison between several Eulerian grid-based codes (which accurately follow the growth of these instabilities) and several Lagrangian SPH codes is performed for an idealised case where a `blob' of gas moves through a uniform density medium."
+ For example. their 44 shows that. for one particular case. the grid-based codes all predict complete disruption of the blob at {TkH (where rig ds the Wl timescale. Le. the time it takes WIL instabilities to fully erow). whereas the SPLL codes predict that the blob shoulc remain intact.," For example, their 4 shows that, for one particular case, the grid-based codes all predict complete disruption of the blob at $t \ga
+\tau_{\rm KH}$ (where $\tau_{\rm KH}$ is the KH timescale, i.e., the time it takes KH instabilities to fully grow), whereas the SPH codes predict that the blob should remain intact."
+ With this in mind. one might conclude that SPL simulations such as ours will overestimate the survivability ol the hot halo of a galaxy.," With this in mind, one might conclude that SPH simulations such as ours will overestimate the survivability of the hot halo of a galaxy."
+ However. it is important to no that Agertz et [find that the grid-based and SPL-basec codes agree with each other rather well for /Tg (see also Appendix A of the present studs).," However, it is important to note that Agertz et find that the grid-based and SPH-based codes agree with each other rather well for $t \la \tau_{\rm KH}$ (see also Appendix A of the present study)."
+" Following the approach of Mori Burkert (2000) (sec also Nulsen 1982: Murray e 11993: and. Mayer et 22006). the Ixelvin-Lelmholtz timescale (including the stabilising effects of gravitv) can be estimated as: where f£ is the barvon fraction of the galaxy. Mo. is the total mass of the ealaxy within the radius down to which the galaxy has been stripped bv ram. pressure. micat is the number density of hydrogen atoms in the ICM. ane 6,4, is he velocity of the galaxy with respect to the IC'M."," Following the approach of Mori Burkert (2000) (see also Nulsen 1982; Murray et 1993; and Mayer et 2006), the Kelvin-Helmholtz timescale (including the stabilising effects of gravity) can be estimated as: where $F$ is the baryon fraction of the galaxy, $M_0$ is the total mass of the galaxy within the radius down to which the galaxy has been stripped by ram pressure, $n_{\rm 
+ICM}$ is the number density of hydrogen atoms in the ICM, and $v_{\rm orb}$ is the velocity of the galaxy with respect to the ICM."
+ For our default 2-ystem run (see 83.3.1). for example. we estimate from ((4) that the Ixelvin-Llelmholtz imescale at pericentre is approximately 4.5 Civr Coe. which is comparable to the curation of our simulations).," For our default 2-system run (see 3.3.1), for example, we estimate from (4) that the Kelvin-Helmholtz timescale at pericentre is approximately 4.5 Gyr (i.e., which is comparable to the duration of our simulations)."
+ Since most of the orbital period of the galaxy is spent [ar [rom »icentre. the value of mou will be substantially longer han this.," Since most of the orbital period of the galaxy is spent far from pericentre, the value of $\tau_{\rm KH}$ will be substantially longer than this."
+ Note also that the timescale associated with he erowth of ICE instabilities is comparable to or exceeds Tsu., Note also that the timescale associated with the growth of RT instabilities is comparable to or exceeds $\tau_{\rm KH}$.
+ Fherefore. we do not expect KLE or ICE. instability stripping to have important consequences for the results or conclusions of this study.," Therefore, we do not expect KH or RT instability stripping to have important consequences for the results or conclusions of this study."
+ We also point out that (C(4) neglects the possibly important. elfects of radiative cooling. physical viscositv. magnetic fields. ete.," We also point out that (4) neglects the possibly important effects of radiative cooling, physical viscosity, magnetic fields, etc.,"
+ all of which will tend to damp (and possibly halt) the growth of such instabilities in real cluster galaxies., all of which will tend to damp (and possibly halt) the growth of such instabilities in real cluster galaxies.
+ Finally. in order to dispel any lingering doubts that our adopted SPILL approach is unable to treat ram pressure stripping accurately. we have mace a cireet comparison of the predictions of the Lagrangian SPILL code CLADCGITI-2 and the Eulerian AMI code FLAS for one of our uniform medium runs.," Finally, in order to dispel any lingering doubts that our adopted SPH approach is unable to treat ram pressure stripping accurately, we have made a direct comparison of the predictions of the Lagrangian SPH code GADGET-2 and the Eulerian AMR code FLASH for one of our uniform medium runs."
+ Phis comparison is presented in Appenclix A and shows that there is excellent. quantitative agreement between the results of the two codes., This comparison is presented in Appendix A and shows that there is excellent quantitative agreement between the results of the two codes.
+ The study of ram pressure stripping of galaxies as they fall into groups and clusters dates back to the seminal paper of Gunn Gott (1972)., The study of ram pressure stripping of galaxies as they fall into groups and clusters dates back to the seminal paper of Gunn Gott (1972).
+ Using a static force argument. these authors derived a simple. physicallv-motivated condition for the instantaneous ram. pressure stripping of a gaseous clisk moving face-on through the ICM.," Using a static force argument, these authors derived a simple, physically-motivated condition for the instantaneous ram pressure stripping of a gaseous disk moving face-on through the ICM."
+" The gas will be stripped ifthe ram pressure. Lai. defined as pesteorl2 (where pic is the density of the ICM. and ei is the speed. of the galaxy with respect to the LOCAL). exeeeds the gravitational restoring force per unit area on the disk. which they derive as 24€CYX,Nus (where X, and Xu are the stellar and gaseous surface densities of the disk. respectivelv)."," The gas will be stripped if the ram pressure, $P_{\rm ram}$, defined as $\rho_{\rm ICM} v_{\rm orb}^2$ (where $\rho_{\rm ICM}$ is the density of the ICM and $v_{\rm orb}$ is the speed of the galaxy with respect to the ICM), exceeds the gravitational restoring force per unit area on the disk, which they derive as $2 \pi 
+G \Sigma_* \Sigma_{\rm gas}$ (where $\Sigma_*$ and $\Sigma_{\rm 
+gas}$ are the stellar and gaseous surface densities of the disk, respectively)."
+ We now seek to derive an analogous mocel for the ram pressure stripping of a sphericallv-symmetric gas cistribution., We now seek to derive an analogous model for the ram pressure stripping of a spherically-symmetric gas distribution.
+ Since it is the least bound material. gas at the outer projected. οσο of the system will be stripped first (see the schematic diagram in Figure 3).," Since it is the least bound material, gas at the outer projected edge of the system will be stripped first (see the schematic diagram in Figure 3)."
+ Consider gas in a projected annulus between radii 2 and. 2| dH., Consider gas in a projected annulus between radii $R$ and $R+dR$ .
+ The projected area of this annulus. d. is 2xIU.," The projected area of this annulus, $dA$, is $2 \pi R dR$."
+" Therefore. the force due to ram pressure on this annulus is simply P,=Paci."," Therefore, the force due to ram pressure on this annulus is simply $F_{\rm ram} 
+= P_{\rm ram} dA$."
+ The annulus of gas will be displaced in the direction opposite to eon Ovhieh we will call the 2 direction) ancl will be stripped. if the force due to the rani pressure exceeds. the maximum gravitational restoring force in this direction., The annulus of gas will be displaced in the direction opposite to $v_{\rm orb}$ (which we will call the $z$ direction) and will be stripped if the force due to the ram pressure exceeds the maximum gravitational restoring force in this direction.
+ The maximum gravitational restoring force. frig. can be written approximately as quisGT)NuGd. where quis(4?) is the maximum restoring acceleration in the 2 direction and Nas) is the projected surface density of the gas in the annulus.," The maximum gravitational restoring force, $F_{\rm grav}$, can be written approximately as $g_{\rm 
+max}(R) \Sigma_{\rm gas}(R) dA$, where $g_{\rm max}(R)$ is the maximum restoring acceleration in the $z$ direction and $\Sigma_{\rm 
+gas}(R)$ is the projected surface density of the gas in the annulus."
+ Therefore.the ram pressure stripping concdition can be written as:," Therefore,the ram pressure stripping condition can be written as:"
+to log(Nei/Niot)=—15 dex.,to $\log(N_\mathrm{el}/N_\mathrm{tot})=-15$ dex.
+ The abundance in the upper atmosphere was determined by the code by fitting observed profiles for a predefined set of abundance step positions between logtsooo=--2 to —5., The abundance in the upper atmosphere was determined by the code by fitting observed profiles for a predefined set of abundance step positions between $\log \tau_\mathrm{5000}=-2$ to $-5$.
+ Then the fit corresponding to each of the stratification models was compared to the one obtained with a homogeneous vertical distribution., Then the fit corresponding to each of the stratification models was compared to the one obtained with a homogeneous vertical distribution.
+ In Fig., In Fig.
+ 13 we illustrate the line profiles fits obtained foru., \ref{Fig10} we illustrate the line profiles fits obtained for.
+ The fit provided by an extreme chemically stratified model yields worse agreement with the observations than does the model with a homogeneous Y distribution., The fit provided by an extreme chemically stratified model yields worse agreement with the observations than does the model with a homogeneous Y distribution.
+ Standard deviation corresponding to different stratification profiles is reported in Table 6.., Standard deviation corresponding to different stratification profiles is reported in Table \ref{tab6}.
+ The first row in this table corresponds to a homogeneous distribution., The first row in this table corresponds to a homogeneous distribution.
+ There is a marginal indication of Y stratification starting at log159099~—2., There is a marginal indication of Y stratification starting at $\log \tau_\mathrm{5000}\approx-2$.
+" This is suggested by a slightly better fit of the observed line profiles, shown in Fig. 14.."," This is suggested by a slightly better fit of the observed line profiles, shown in Fig. \ref{Fig11}."
+ Ti line profiles were examined in the same way as Y. Figure 15 illustrates the individual line profile fit in the format similar to Fig. 13.., Ti line profiles were examined in the same way as Y. Figure \ref{Fig12} illustrates the individual line profile fit in the format similar to Fig. \ref{Fig10}.
+" The assumption that all Ti is concentrated in the upper layers of the stellar atmosphere gives significantly worse results, compared to the homogeneous model."," The assumption that all Ti is concentrated in the upper layers of the stellar atmosphere gives significantly worse results, compared to the homogeneous model."
+" Similar to the case of Y, we found a marginal indication of Ti stratification with a step position at log15099=--2."," Similar to the case of Y, we found a marginal indication of Ti stratification with a step position at $\log \tau_\mathrm{5000}=-2$."
+ Our study shows that ddoes not possess a strong magnetic field., Our study shows that does not possess a strong magnetic field.
+" Owing to the quality of our spectropolarimetric data, we were able to set a sensitive upper limit of 4 G on the mean longitudinal magnetic field."," Owing to the quality of our spectropolarimetric data, we were able to set a sensitive upper limit of 4 G on the mean longitudinal magnetic field."
+ Our results also do not support the hypothesis that magnetic field is associated with spots of chemical elements., Our results also do not support the hypothesis that magnetic field is associated with spots of chemical elements.
+" From the spectral lines of inhomogeneously distributed chemical elements, we inferred the upper limit of 115 G for(B;)."," From the spectral lines of inhomogeneously distributed chemical elements, we inferred the upper limit of 15 G for."
+". As in the case of other HgMn stars studied with high-precision spectropolarimetry (??), the results of this work clearly demonstrate the absence of global magnetic field structures."," As in the case of other HgMn stars studied with high-precision spectropolarimetry \citep{Auriere:2010, Makaganiuk:2011a}, the results of this work clearly demonstrate the absence of global magnetic field structures."
+" Analysis of LSD profiles also rules out the presence ofcomplex magnetic fields, similar to those found in active late-type stars (2).."," Analysis of LSD profiles also rules out the presence ofcomplex magnetic fields, similar to those found in active late-type stars \citep{Donati:2009}. ."
+Rotation in stars has a number of profound effects on their evolution.,Rotation in stars has a number of profound effects on their evolution.
+ Not only are there significant changes in the hydrostatic structure (Endal&Sofia1976) but this causes a thermal imbalance that can lead to a strong meridional circulation current (Sweet1950)., Not only are there significant changes in the hydrostatic structure \citep{Endal76} but this causes a thermal imbalance that can lead to a strong meridional circulation current \citep{Sweet50}.
+. Meridional circulation leads to additional shear which induces a number of instabilities., Meridional circulation leads to additional shear which induces a number of instabilities.
+ The resulting turbulence leads to strong mixing of both angular momentum and chemical composition., The resulting turbulence leads to strong mixing of both angular momentum and chemical composition.
+" Although most modellers include all of these effects, the exact implementation of stellar rotation can vary dramatically."," Although most modellers include all of these effects, the exact implementation of stellar rotation can vary dramatically."
+" For example Heger,Langer&Woosley(2000) use a model where meridional circulation is treated diffusively.", For example \citet{Heger00} use a model where meridional circulation is treated diffusively.
+ The diffusion owing to shear is the linear combination of a number of coefficients based on different possible instabilities., The diffusion owing to shear is the linear combination of a number of coefficients based on different possible instabilities.
+" On the other hand, models such as those of Zahn(1992),, Talonetal.(1997) and Maeder(2003) treat the circulation as advective and use a single diffusion coefficient based on the magnitude of Kelvin-Helmholtz instabilities induced by shear."," On the other hand, models such as those of \citet{Zahn92}, \citet{Talon97} and \citet{Maeder03} treat the circulation as advective and use a single diffusion coefficient based on the magnitude of Kelvin-Helmholtz instabilities induced by shear."
+ In these models it is also necessary to define the magnitude of diffusion along isobars and this too varies between different models., In these models it is also necessary to define the magnitude of diffusion along isobars and this too varies between different models.
+" We have shown that different models generally give rise to similar qualitative conclusions but there are significant differences in the results based on mass, rotation rate and metallicity."," We have shown that different models generally give rise to similar qualitative conclusions but there are significant differences in the results based on mass, rotation rate and metallicity."
+ There are also open questions about how angular momentum transport occurs in convective zones., There are also open questions about how angular momentum transport occurs in convective zones.
+" Comparing the models based on Talonetal. and Maeder(2003),, case 1, and that based on Heger,Langer&Woosley (2000),, case 2, we find that case 1 gives higher luminosity stars for masses less massive than 10 andmoreluminous, hotter starsathighermassesthancase 2."," Comparing the models based on \citet{Talon97} and \citet{Maeder03}, case 1, and that based on \citet{Heger00}, case 2, we find that case 1 gives higher luminosity stars for masses less massive than $10$ $\,$ and more luminous, hotter stars at higher masses than case 2."
+ massstarsgivesimilarlevelsofnitrogenenrichmentineachcasebutcase 2prod massandintermediate—massstarsthancase 1., High-mass stars give similar levels of nitrogen enrichment in each case but case 2 produces far less enrichment for low-mass and intermediate-mass stars than case 1.
+Thesituationissimilar forthe Senrichment., The situation is similar for their helium-3 enrichment.
+ 'The predicted effects of rotation appear to be highly, The predicted effects of rotation appear to be highly
+as single particles of mass Alc=210 MAL. with a softening of LOppe.,"as single particles of mass $M_{\rm GC} = 2
+\times 10^5$ $_{\odot}$ with a softening of pc."
+ We do not expect our conclusions to hange if we used a particle model for cach globular since rev are stable against tidal disruption within Fornax., We do not expect our conclusions to change if we used a particle model for each globular since they are stable against tidal disruption within Fornax.
+ We start the globulars outside the core. mostly on circular orbits and let them orbit. expecting them to spiral in to the centre of their respective host. haloes due to dynamical friction.," We start the globulars outside the core, mostly on circular orbits and let them orbit, expecting them to spiral in to the centre of their respective host haloes due to dynamical friction."
+ The distance from the centre of the host halo as a function of time. r(/). can be computed using Chancdrasckhar’s dynamical friction. formula.1987).. which is given by: which gives the force acting on the massive particle crossing the halo.," The distance from the centre of the host halo as a function of time, $r(t)$, can be computed using Chandrasekhar's dynamical friction formula, which is given by: which gives the force acting on the massive particle crossing the halo."
+" The density profile p(r) is given by our equation 1.. anc we assume that the velocity distribution is isotropic and Alaxwellian at all radii,"," The density profile $\rho(r)$ is given by our equation \ref{eq:cusp}, and we assume that the velocity distribution is isotropic and Maxwellian at all radii."
+ Of. course this assumption does not hold. but is good enough for our purposes2004).," Of course this assumption does not hold, but is good enough for our purposes."
+. We can then easily caleulate the velocity dispersion using the Jeans equatIon: We [ind similar sinking times for eccentric orbits. therefore for brevity we show only the circular orbits in this paper and leave the detailed parameter space study for a future paper which explores the technical aspects of dynamical friction in structures with cdillerent. density profiles.," We can then easily calculate the velocity dispersion using the Jeans equation: We find similar sinking times for eccentric orbits, therefore for brevity we show only the circular orbits in this paper and leave the detailed parameter space study for a future paper which explores the technical aspects of dynamical friction in structures with different density profiles."
+ Additionally we assume. that Afc« nya.," Additionally we assume, that $M_{\rm GC} \gg$ $m_{\rm par}$."
+ This is a little problematic in the case of the cuspy and the big cored potential because the particles in the outermost shell have my=310 MM. and Hay=1210 MM... respectively.," This is a little problematic in the case of the cuspy and the big cored potential because the particles in the outermost shell have $m_{\rm par} = 3 \times 10^4$ $_\odot$ and $m_{\rm par}
+= 1.2 \times 10^5$ $_\odot$, respectively."
+ However. these particles rarely. penetrate the innermost kkpe of the halo.," However, these particles rarely penetrate the innermost kpc of the halo."
+ In equation 2.. In(à). is the Coulomb {οσαha: in this definition bias is the largest impact parameter o be considered.," In equation \ref{eq:dff}, ${\rm
+ln} \Lambda (r)$ , is the Coulomb logarithm: in this definition $b_{\rm max}$ is the largest impact parameter to be considered."
+ This parameter is defined by one of the assumptions Chandrasekhar made while deriving the above dvnanmical friction formula: The intrucer must be moving hrough a medium with constant density. therefore bya is he greatest distance for which this is still valid.," This parameter is defined by one of the assumptions Chandrasekhar made while deriving the above dynamical friction formula: The intruder must be moving through a medium with constant density, therefore $b_{\rm max}$ is the greatest distance for which this is still valid."
+ We keep buns as à free parameter when fitting our analytic formulae o the simulations., We keep $b_{\rm max}$ as a free parameter when fitting our analytic formulae to the simulations.
+ We find for the cuspy haloes bans=25 kkpe and for the cored ones busLOkkpe., We find for the cuspy haloes $b_{\rm max} = 0.25$ kpc and for the cored ones $b_{\rm max} = 1.0$ kpc.
+ In this equation ος(1) is the cireular velocity at radius r., In this equation $v_{\rm c}(r)$ is the circular velocity at radius $r$.
+ The force exerted by dynamical frietion on the perturber is tangential with respect to its movement and thus causes the cluster to lose angular momentum per unit mass at a rate Since the cluster continues to orbit at a speed ei(r) as it spirals to the centre. its angular momentum per unit mass at radius rds at all times £=ορ)," The force exerted by dynamical friction on the perturber is tangential with respect to its movement and thus causes the cluster to lose angular momentum per unit mass at a rate Since the cluster continues to orbit at a speed $v_{\rm c}(r)$ as it spirals to the centre, its angular momentum per unit mass at radius $r$ is at all times $L = rv_{\rm c}(r)$."
+ Substituting the time derivative of this expression into equation 5 we obtain Substituting values for the initial radii we obtain the analytical curves drawn in figure 2. plotted on top of the results from the numerical simulations., Substituting the time derivative of this expression into equation \ref{eq:L1} we obtain Substituting values for the initial radii we obtain the analytical curves drawn in figure \ref{figcpco} plotted on top of the results from the numerical simulations.
+ For the cuspy potential the analvtie calculation agrees very well with the numerical simulation., For the cuspy potential the analytic calculation agrees very well with the numerical simulation.
+ Haloes with a core give a poorer agreement., Haloes with a core give a poorer agreement.
+ After an initial sinking rate that agrees well with the analytic expectation. the elobulars sink faster as they approach twice the core radius. and then stop sinking at the core radius.," After an initial sinking rate that agrees well with the analytic expectation, the globulars sink faster as they approach twice the core radius, and then stop sinking at the core radius."
+ The analytic formula predicts a continued. but slow infall to the centre.," The analytic formula predicts a continued, but slow infall to the centre."
+ “Phis resonancescattering οσοι will be investigated in à more detailed paper2006)., This resonance/scattering effect will be investigated in a more detailed paper.
+. We note. however. that it is not trivially due to the facet that the globular is of comparable mass to that enclosed by its orbit: the radius at which AZ(r)=Mec is ~ three times smaller than the core radius (see Figure 4)).," We note, however, that it is not trivially due to the fact that the globular is of comparable mass to that enclosed by its orbit; the radius at which $M(r)=M_{GC}$ is $\sim$ three times smaller than the core radius (see Figure \ref{figcore}) )."
+ The stalling results are apparent in both ofthe cored halo simulations (small core and big core)., The stalling results are apparent in both of the cored halo simulations (small core and big core).
+ We conclude that the presence of a central density core leads to the infall of the elusters stopping at the core radius: the problem is in this sense scalable., We conclude that the presence of a central density core leads to the infall of the clusters stopping at the core radius; the problem is in this sense scalable.
+merge (Conselice|2009) so we can expect that the effect of galaxy mergers might make a difference between the low and high redshift range.,merge \citep{2009MNRAS.399L..16C} so we can expect that the effect of galaxy mergers might make a difference between the low and high redshift range.
+" As expected from different magnitude cutoffs, the values of N and b are lower for the brighter 1b sample than the fainter la sample."," As expected from different magnitude cutoffs, the values of $\overline{N}$ and $b$ are lower for the brighter 1b sample than the fainter 1a sample."
+ We can compare the value of b for the brighter sample with the expected value of b if the brighter sample is a randomly selected sample from a fainter parent sample using (Lahav&Saslaw|1992):: This allows us to compare the clustering of the brighter la sample with the fainter 1b sample., We can compare the value of $b$ for the brighter sample with the expected value of $b$ if the brighter sample is a randomly selected sample from a fainter parent sample using \citep{1992ApJ...396..430L}: This allows us to compare the clustering of the brighter 1a sample with the fainter 1b sample.
+" We find that for all samples, the brighter galaxies from the 1b samples have a higher value of b than would be expected if they were a random subsample of the 1a sample."," We find that for all samples, the brighter galaxies from the 1b samples have a higher value of $b$ than would be expected if they were a random subsample of the 1a sample."
+" This means that bright galaxies are more strongly clustered than fainter galaxies which may be result of brighter galaxies being concentrated around clustera centers, or in dark matter haloes."," This means that bright galaxies are more strongly clustered than fainter galaxies which may be a result of brighter galaxies being concentrated around cluster centers, or in dark matter haloes."
+" If so, it suggests an upper limit of about 10! Mpc for the size of these haloes."," If so, it suggests an upper limit of about $10 h^{-1}$ Mpc for the size of these haloes."
+ We summarize the comparisons in table [7] and plot the comparison of b for 3D cells over a range of cell sizes in figure [7]., We summarize the comparisons in table \ref{tres-bscomp} and plot the comparison of $b$ for 3D cells over a range of cell sizes in figure \ref{fres-bscomp}.
+ In this paper we have compared the counts-in-cells distribution of galaxies fy(N) with two theoretical models and found that the observed distribution has large field to field variations which may be as great as 20% across quadrants., In this paper we have compared the counts-in-cells distribution of galaxies $f_V(N)$ with two theoretical models and found that the observed distribution has large field to field variations which may be as great as $20\%$ across quadrants.
+" These large variations essentially mean that there is a considerable amount of cosmic variance in the data, and that the galaxies in different quadrants are not identically distributed."," These large variations essentially mean that there is a considerable amount of cosmic variance in the data, and that the galaxies in different quadrants are not identically distributed."
+ This also means that errors determined using the jackknife procedure will underestimate the true range of variation in the data because different subsets of the data are not identically distributed., This also means that errors determined using the jackknife procedure will underestimate the true range of variation in the data because different subsets of the data are not identically distributed.
+ We see the existence of these subregions of different local density from the bumps in the counts-in-cells distribution for large cells., We see the existence of these subregions of different local density from the bumps in the counts-in-cells distribution for large cells.
+" We note that as shown by|Coleman&Saslaw|(1990)),, these bumps may be the result of regions of different local density."," We note that as shown by \citet{1990ApJ...353..354C}, these bumps may be the result of regions of different local density."
+" As suggested by (2009), a larger survey volume will show less effects of cosmic variance, and an earlier analysis of the 2MASS catalog by provides a hint that this may be the case."," As suggested by \citet{2009A&A...508...17S}, a larger survey volume will show less effects of cosmic variance, and an earlier analysis of the 2MASS catalog by \citet{2005ApJ...626..795S} provides a hint that this may be the case."
+" In the 2MASS analysis, found less variation between quadrants, on the order of 5% instead of the 20% we have found in the SDSS."," In the 2MASS analysis, \citet{2005ApJ...626..795S} found less variation between quadrants, on the order of $5\%$ instead of the $20\%$ we have found in the SDSS."
+" We note that after excluding regions close to the galactic plane using |||<20°, the 2MASS catalog covers thirds of the sky, a coverage that is more than three times that of the SDSS."," We note that after excluding regions close to the galactic plane using $|l| < 20^\circ$, the 2MASS catalog covers two-thirds of the sky, a coverage that is more than three times that of the SDSS."
+" Although the SDSS “great wall"" is contained within the low redshift range, it covers only a small fraction of the SDSS footprint in a stripe within 6° of the celestial equator."," Although the SDSS “great wall” is contained within the low redshift range, it covers only a small fraction of the SDSS footprint in a stripe within $6^\circ$ of the celestial equator."
+ For this reason we do not see much difference between the low redshift range and the high redshift range beyond differences in N because the overdensity in the “great wall” is not large enough to dominate over the rest of the survey., For this reason we do not see much difference between the low redshift range and the high redshift range beyond differences in $\overline{N}$ because the overdensity in the “great wall” is not large enough to dominate over the rest of the survey.
+" Indeed, because the “great wall” is most likely porous in three dimensions, it is not clear that “great wall” is an accurate description."," Indeed, because the “great wall” is most likely porous in three dimensions, it is not clear that “great wall” is an accurate description."
+ Our comparison of y for 2D projected cells and 3D spherical cells in redshift space shows that there is a difference in the exponent y of the two-point correlation function between the 2D and 3D sample., Our comparison of $\gamma$ for 2D projected cells and 3D spherical cells in redshift space shows that there is a difference in the exponent $\gamma$ of the two-point correlation function between the 2D and 3D sample.
+ The 3D samples have a value of y that is lower than that for the 2D samples because the 3D samples in redshift space are affected by peculiar velocity distortions which change the apparent clustering in redshift space., The 3D samples have a value of $\gamma$ that is lower than that for the 2D samples because the 3D samples in redshift space are affected by peculiar velocity distortions which change the apparent clustering in redshift space.
+ Using the well-known relation between the projected correlation function and the real space correlation function we find that the difference in y between the 2D and 3D samples is a measure of the redshift space distortions., Using the well-known relation between the projected correlation function and the real space correlation function we find that the difference in $\gamma$ between the 2D and 3D samples is a measure of the redshift space distortions.
+ Our findings for the value of y suggest that the difference between the projected sample and redshift space sample is about 0.2~0.3 in agreement with work by and using earlier catalogs., Our findings for the value of $\gamma$ suggest that the difference between the projected sample and redshift space sample is about $0.2 \sim 0.3$ in agreement with work by \citet{1994ApJ...425....1F} and \citet{2003MNRAS.346...78H} using earlier catalogs.
+" Comparing the low redshift and high redshift range, we find that there is a large difference in N between the low and high redshift range for samples with the same magnitude cutoff which may be caused by galaxy"," Comparing the low redshift and high redshift range, we find that there is a large difference in $\overline{N}$ between the low and high redshift range for samples with the same magnitude cutoff which may be caused by galaxy"
+concentrated narrowly towards K=0. Le. those with only a few significant modes. the low-&. large-scale modes will dominate for any finite field size.,"concentrated narrowly towards $\vec{k}=\vec{0}$, i.e. those with only a few significant modes, the $k$, large-scale modes will dominate for any finite field size."
+ But for a wide power spectrum. where many small-scale modes contribute. a sufficiently large field size should lead to approximate tsotropy.," But for a wide power spectrum, where many small-scale modes contribute, a sufficiently large field size should lead to approximate isotropy."
+ This is demonstrated for à two-dimensional field with a power spectrum Pik)=K in Fig. 5.., This is demonstrated for a two-dimensional field with a power spectrum $P(k) = k^{-2}$ in Fig. \ref{fig:pxi_2dim_isotropy}.
+ Both panels show the probability distributions for two separation vectors ¥ and 1” of the same length. but rotated by 45°. (," Both panels show the probability distributions for two separation vectors $\vec{x}$ and $\vec{x}^\prime$ of the same length, but rotated by $45^{\circ}$. ("
+"Or. more precisely. with X=(Qv.0) and V.=νιcos(45°+ey.x,—ey. with €=107"". to avoid double poles.)","Or, more precisely, with $\vec{x}=(x_1,0)$ and $\vec{x}^\prime=\left(x_1\cos(45^{\circ}+\epsilon), x_1\cos(45^{\circ}-\epsilon)\right)$, with $\epsilon=10^{-10}$, to avoid double poles.)"
+" For a large separation to field size ratio. 1/1,=0.3. seen in the left panel. the distributions for different separation vectors are quite different. while for 1/1.= 0.03. in the right panel. they have almost converged."," For a large separation to field size ratio, $|\vec{x}|/L=0.3$, seen in the left panel, the distributions for different separation vectors are quite different, while for $|\vec{x}|/L=0.03$ , in the right panel, they have almost converged."
+ Note that. for Pik)=&7 in 2 dimensions. the field size L cancels out of the ο. and therefore we could keep the power spectrum normalisation constant while changing |3]/L. without changing the scale of p(é).," Note that, for $P(k) = k^{-2}$ in 2 dimensions, the field size $L$ cancels out of the $\sigma_n$, and therefore we could keep the power spectrum normalisation constant while changing $|\vec{x}|/L$, without changing the scale of $p(\xi)$."
+ Isotropy is even easier to obtain for power spectra that are small both for very small and very large & and large only at intermediate wavelengths. since then the low-&. large scale modes do not harm isotropy.," Isotropy is even easier to obtain for power spectra that are small both for very small and very large $k$ and large only at intermediate wavelengths, since then the $k$, large scale modes do not harm isotropy."
+ The ACDM power spectrum in cosmology fulfils this condition. allowing tsotropie N-body simulations with reasonable field sizes.," The $\Lambda$ CDM power spectrum in cosmology fulfils this condition, allowing isotropic N-body simulations with reasonable field sizes."
+ However. these are no fundamental changes. and the resulting distribution function has all the same properties às the one-dimensional version.," However, these are no fundamental changes, and the resulting distribution function has all the same properties as the one-dimensional version."
+ So all our results can be readily applied to higher dimensions. as long às the computational difficulties can be handled.," So all our results can be readily applied to higher dimensions, as long as the computational difficulties can be handled."
+ We have considered the problem of accurate likelihood functions for Bayesian analyses of data from Gaussian random fields., We have considered the problem of accurate likelihood functions for Bayesian analyses of data from Gaussian random fields.
+ Making use of Fourier mode expansions and characteristic. funetions. we have derived analytically the probability distribution function of the correlation function for a one-dimensional finite Gaussian random field.," Making use of Fourier mode expansions and characteristic functions, we have derived analytically the probability distribution function of the correlation function for a one-dimensional finite Gaussian random field."
+ For general power spectra. we can only give a sum formula for the distribution function.," For general power spectra, we can only give a sum formula for the distribution function."
+ However. for the special power spectrum PUOxK7. we have found an explicit expression in terms of elliptic theta functions.," However, for the special power spectrum $P(k)\propto k^{-2}$, we have found an explicit expression in terms of elliptic theta functions."
+ We can also. for general power spectra. calculate arbitrary moments and cumulants of the distribution by much simpler sum expressions.," We can also, for general power spectra, calculate arbitrary moments and cumulants of the distribution by much simpler sum expressions."
+" Then. we continued the analytical approach for bivariate ""Sistributions as well. and found a similar. but even. more complicated sum formula às a result."," Then, we continued the analytical approach for bivariate distributions as well, and found a similar, but even more complicated sum formula as a result."
+" We have also outlined a general procedure for calculating arbitrary high multivariate ""Sistributions. though this ts not feasible in practice."," We have also outlined a general procedure for calculating arbitrary high multivariate distributions, though this is not feasible in practice."
+ Furthermore. we have considered the analytical properties of the new probability distribution function. which are well understood and consistent with numerical results.," Furthermore, we have considered the analytical properties of the new probability distribution function, which are well understood and consistent with numerical results."
+ We used the moments of the univariate distribution to construct an Edgeworth expansion. which ts able to closely approximate the true distribution function m the region of highest likelihood. and therefore could be a useful replacement of simple Gaussian approximations in Bayesian analyses.," We used the moments of the univariate distribution to construct an Edgeworth expansion, which is able to closely approximate the true distribution function in the region of highest likelihood, and therefore could be a useful replacement of simple Gaussian approximations in Bayesian analyses."
+ Finally. we found that all our results easily generalise to multi-dimensional fields.," Finally, we found that all our results easily generalise to multi-dimensional fields."
+" Consideri5 possible future applications of these results. we have to point out the importance of improving the correlation fuction likelihood. as well as the further steps that are necessary for a practical implementation,"," Considering possible future applications of these results, we have to point out the importance of improving the correlation function likelihood, as well as the further steps that are necessary for a practical implementation."
+ As already mentioned in the introduction. the Gaussian approximatio for the likelihood can lead to considerable deviations in. parameter estimation.," As already mentioned in the introduction, the Gaussian approximation for the likelihood can lead to considerable deviations in parameter estimation."
+ For example. in cosmic shear studies. this leads to significantly. reduced. accuracy (?)..," For example, in cosmic shear studies, this leads to significantly reduced accuracy \citep{Hartlap2009}."
+ Similar effects are to be expected in other fields where correlation functions are used., Similar effects are to be expected in other fields where correlation functions are used.
+ However. the work presented in this article has so far been purely mathematical. and the results are not readily applicable to real data.," However, the work presented in this article has so far been purely mathematical, and the results are not readily applicable to real data."
+ The main obstacle lies in the infeasibility of analytical caleulations. for higher multivariate distributions., The main obstacle lies in the infeasibility of analytical calculations for higher multivariate distributions.
+ If data of the correlation function. over N bins needs to be analysed. we would need the full N-variate distribution function.," If data of the correlation function over $N$ bins needs to be analysed, we would need the full $N$ -variate distribution function."
+ Therefore. we expect that a numerical approach. as by ?.. is best suited for practical computations.," Therefore, we expect that a numerical approach, as by \cite{Wilking2011}, is best suited for practical computations."
+" Still. their. ""quasi-Gaussian approach makes direct use of the analytical univariate distribution function. presented in. this article."," Still, their 'quasi-Gaussian' approach makes direct use of the analytical univariate distribution function presented in this article."
+ We also expect that the analytical results will yield important guidance and cross-checks for future numerical implementations., We also expect that the analytical results will yield important guidance and cross-checks for future numerical implementations.
+ We also note that our analytical results depend on the assumption of a Gaussian random field. whereas a lot of cosmological data probes the evolved density field on small scales. which ts far from Gaussian.," We also note that our analytical results depend on the assumption of a Gaussian random field, whereas a lot of cosmological data probes the evolved density field on small scales, which is far from Gaussian."
+ Nevertheless. applications for the analytical likelihood function could be found on very large scales. or in cosmic microwave background analysis. since the density fluctuations at that epoch were still either Gaussian or close to Gaussian.," Nevertheless, applications for the analytical likelihood function could be found on very large scales, or in cosmic microwave background analysis, since the density fluctuations at that epoch were still either Gaussian or close to Gaussian."
+ Furthermore. this work could also be relevant to fields outside of cosmology. for example in the common problem of time series analysis of Gaussiar random processes.," Furthermore, this work could also be relevant to fields outside of cosmology, for example in the common problem of time series analysis of Gaussian random processes."
+ So far. to keep calculations simple.we have considered simple poles only.," So far, to keep calculations simple,we have considered simple poles only."
+" However. it is entirely possible to have multiple poles in the characteristicfunction. re. to have some mode numbers 7z with C,;2 C,."," However, it is entirely possible to have multiple poles in the characteristicfunction, i.e. to have some mode numbers $\vec{m} \neq \vec{n}$ with $C_{\vec{m}} = C_{\vec{n}}$ ."
+ This could happen if, This could happen if
+"each molecular line was estimated by considering the range of projected radii where [,,(p)pdp, with I,, the intensity at the line center and p the impact parameter, exceeds half of its maximum value (see Fig. 3)).","each molecular line was estimated by considering the range of projected radii where $I_{\nu_0}(p)\,p\,dp$, with $I_{\nu_0}$ the intensity at the line center and $p$ the impact parameter, exceeds half of its maximum value (see Fig. \ref{fig:velocity}) )."
+ We note that the radial extent of the line formation region is almost insensitive to the exact value of the velocity structure., We note that the radial extent of the line formation region is almost insensitive to the exact value of the velocity structure.
+" Optical depths effects can strongly affect the observed line widths, and detailed radiative transfer modelling (as presented in this Letter) is required to determine the underlying velocity structure."," Optical depths effects can strongly affect the observed line widths, and detailed radiative transfer modelling (as presented in this Letter) is required to determine the underlying velocity structure."
+" Most of the observed lines have line widths in excess of kkm/s. However, a few lines are considerably narrower, and their line formation regions are located in that part of the envelope where the wind has not yet reached its terminal velocity kkm/s)."," Most of the observed lines have line widths in excess of km/s. However, a few lines are considerably narrower, and their line formation regions are located in that part of the envelope where the wind has not yet reached its terminal velocity km/s)."
+" Although the line formation regions are quite broad, we find the first observational evidence that the wind acceleration is slower than implied by the momentum equation, corroborating the results of ?.."," Although the line formation regions are quite broad, we find the first observational evidence that the wind acceleration is slower than implied by the momentum equation, corroborating the results of \citet{Decin2010}. ."
+" Using the classical 6-parametrization (e.g.,?) to simulate the velocity structure, where vp is the velocity at the dust condensation radius and Uo. is the terminal velocity, we find that 1€8<2."," Using the classical $\beta$ -parametrization \citep[e.g.,][]{Lamers1999isw..book.....L} to simulate the velocity structure, where $\varv_0$ is the velocity at the dust condensation radius and $\varv_\infty$ is the terminal velocity, we find that $1\le\beta\le2$."
+ The theoretical line predictions shown in this Letter are computed for B=1., The theoretical line predictions shown in this Letter are computed for $\beta=1$.
+" The momentum equation, in contrast, matches a much steeper velocity profile with 6=0.6."," The momentum equation, in contrast, matches a much steeper velocity profile with $\beta$ =0.6."
+ We note that the velocity structure derived by ? for y Cyg is compliant with a B-value of 0.9., We note that the velocity structure derived by \citet{Justtanont2010} for $\chi$ Cyg is compliant with a $\beta$ -value of 0.9.
+" As discussed by ?,, there are several possible causes for a less steep velocity structure."," As discussed by \citet{Decin2010}, there are several possible causes for a less steep velocity structure."
+" We summarize that a slower wind velocity may be caused by incomplete momentum coupling, dust emission being slightly optically thick to the stellar radiation, the fact that not all dust species are formed at the same time and at the same radial distance, and/or that some dust species are inefficient as wind drivers."," We summarize that a slower wind velocity may be caused by incomplete momentum coupling, dust emission being slightly optically thick to the stellar radiation, the fact that not all dust species are formed at the same time and at the same radial distance, and/or that some dust species are inefficient as wind drivers."
+" Using the HIFI spectrometer onboard Herschel, we have observed the evolved oxygen-rich Mira star, IK Tau, in 31 molecular emission lines."," Using the HIFI spectrometer onboard Herschel, we have observed the evolved oxygen-rich Mira star, IK Tau, in 31 molecular emission lines."
+" For the first time, several lines of ortho and para H,°0, and the rarer isotoplogs H,’O and H;80 have been detected."," For the first time, several lines of ortho and para $_2^{16}$ O, and the rarer isotoplogs $_2^{17}$ O and $_2^{18}$ O have been detected."
+ We have deduced a totalwater content, We have deduced a totalwater content
+ ‘rotational aud vibration-rotation trausitions of the 'eactive lous includiug the first quautitative 'edietion Of an astronomical spectrum of H5.,of rotational and vibration-rotation transitions of the reactive ions including the first quantitative prediction of an astronomical spectrum of $_2^+$.
+" It appea5S that superthermally excited eiissiou lines ""Πο aud H4 may be detectable at wavelengths between 3 axd 6 402. Strong thermal emission . H4..may. be measurable in the [ar infrared. at 20 to 150 jin wavelength.", It appears that superthermally excited emission lines of $_2^+$ and $_3^+$ may be detectable at wavelengths between 3 and 6 $\mu$ m. Strong thermal emission of $_3^+$ may be measurable in the far infrared at 50 to 150 $\mu$ m wavelength.
+ W[9B seems to be 1 outstanding candidate for the detection. with precicted ioiizalion rate more than [| orders of uagnuitude above Gaactic average.," W49B seems to be an outstanding candidate for the detection, with predicted ionization rate more than 4 orders of magnitude above Galactic average."
+ Futufe neasurernmerts by HHerschel and ALLA provide the opportunity to make observations of iorlzatloi signatures in 5JU/MC systems ad analyze their 'orrelatiou to the gaο emission., Future measurements by Herschel and ALMA provide the opportunity to make observations of ionization signatures in SNR/MC systems and analyze their correlation to the gamma-ray emission.
+ A yeceit observation stMODeeests an enhanced excitation and ibundance of ammolia toward W28 (Nicholasetal.2011)., A recent observation suggests an enhanced excitation and abundance of ammonia toward W28 \citep{nicholas2011}.
+. Using Radex. it is possible to explain lie observed spectruu as CR induced with an jonization rate of L.5-10I? 1 so a factor of >10! 1ove Calactic average.," Using , it is possible to explain the observed spectrum as CR induced with an ionization rate of $4.5\cdot 10^{-12}$ $^{-1}$ , so a factor of $>10^{4}$ above Galactic average."
+ Further. H4 etUssion observed from the direction of IC[13 iudicates an eibanced ionization evel of 2-1015 &1 SO a [actor of 10 above Calactic average 2010b).," Further, $_{3}^{+}$ emission observed from the direction of IC443 indicates an enhanced ionization level of $2\cdot 10^{-15}$ $^{-1}$, so a factor of 10 above Galactic average \citep{indriolo_ic443_h3p_2010}."
+". In tle future. more observatious need to be performed iu order to explicitly search for the expected sigtattwes of H, aud H4."," In the future, more observations need to be performed in order to explicitly search for the expected signatures of $_2^+$ and $_3^+$."
+ Treating the propagation of the CRs in detail will allow us to give even mo'eprecise predictions of the expected emission regions., Treating the propagation of the CRs in detail will allow us to give even moreprecise predictions of the expected emission regions.
+distinguished from a young stellar component with dust extinction.,distinguished from a young stellar component with dust extinction.
+ Hot Dust also has a significant contribution to the continuum., Hot Dust also has a significant contribution to the continuum.
+" It has a peak in the nuclear region, which, as reported by Thatte et al. ("," It has a peak in the nuclear region, which, as reported by Thatte et al. ("
+"1997), is a strong evidence of a dense circumnuclear torus.","1997), is a strong evidence of a dense circumnuclear torus."
+ It also has peaks in the, It also has peaks in the
+Ikobavashi et al. (,Kobayashi et al. (
+2010) showed that planetesinials mainly accreting onto embrvos have i=100m.,2010) showed that planetesimals mainly accreting onto embryos have $m = 100 m_0$.
+" For myZ1075 & (ryZ3«107 kk). final enibryo mnasses exceed 1074, at AAU in à MMSN disk. but embrvos caunot reach it within a disk lifetime due to their slow growth."," For $m_0 \ga 10^{23}\,$ g $r_0 \ga 3 \times 10^{3}$ km), final embryo masses exceed $10\,M_\oplus$ at AU in a MMSN disk, but embryos cannot reach it within a disk lifetime due to their slow growth."
+ The final mass M4 isindependent of Moy. while hich “oy inereases Mo because eas drag highly damps 6.," The final mass $M_{\rm
+ca}$ isindependent of $\Sigma_{\rm s,0}$, while high $\Sigma_{\rm g,0}$ increases $M_{\rm ca}$ because gas drag highly damps $\tilde e$."
+" For X4— (4104 MNSN). initial planctesimals with ryZ bÜkkm can produce an embryo with 1037, at AAU."," For $\Sigma_1 = 71 \, {\rm g\,cm}^{-2}$ $10\times$ MMSN), initial planetesimals with $r_0 \ga 50$ km can produce an embryo with $10\,M_\oplus$ at AU."
+ For comparison. we also show the final mass Al. in the same situation but ucelecting the atinosphiere (I&obavashi e al.," For comparison, we also show the final mass $M_{\rm c}$ in the same situation but neglecting the atmosphere (Kobayashi et al."
+" 2010): where InayaX.)n js estimated. from Equatiou(22)) with Ce,=0.1 for M=LAL, in the MMSN inodel."," 2010): where $\ln(\Sigma_{\rm s,0}/\Sigma_{\rm s}) \simeq 4.5$ is estimated from \ref{eq:sigma_limit}) ) with $C_{\Sigma_{\rm s}}=0.1$ for $M=0.1 M_{\oplus}$ in the MMSN model."
+" The collisional euliauceimoeut due to the atmosphere is inefficient for m=Ls107 ee: AL,<AL.", The collisional enhancement due to the atmosphere is inefficient for $m = 4 \times 10^{20}$ g; $M_{\rm ca} < M_{\rm a}$.
+" IbiXLx 10ο, the atinosphere contributes to embryo growth."," If $m \ga 4 \times 10^{22}\,$ g, the atmosphere contributes to embryo growth."
+ As described above. planetesimals are ground down bv collision cascade and resulting siuall ragiuents spiral iuto the central star by eas drag.," As described above, planetesimals are ground down by collision cascade and resulting small fragments spiral into the central star by gas drag."
+ In the steady state of collision cascade. he surface density of planetesimals is iuuch arecr than that of fragments.," In the steady state of collision cascade, the surface density of planetesimals is much larger than that of fragments."
+ However. when he erinding of plauetesinals is much quicker han the removal of small fragiieuts by eas drag. rasnients accumulate at the low-nass end of collision cascade and determine the total mass of )odies.," However, when the grinding of planetesimals is much quicker than the removal of small fragments by gas drag, fragments accumulate at the low-mass end of collision cascade and determine the total mass of bodies."
+" where C,~Lis a constant.", where $C_{\rm L} \sim 1$ is a constant.
+ we apply CL=0.5¢7.," we apply $C_{\rm
+L}=0.5$."
+. between stinius by clubrvos aud damping by eas drag is obtained ax (Ixobavashietal.2010) Usine Eqs. (13)). (26)).," between stirring by embryos and damping by gas drag is obtained as \citep{kobayashi+10}
+ Using Eqs. \ref{eq:tau}) ), \ref{eq:condition_for_lowmass}) ),"
+ aud (27)) under the Stokes regime. we have the fragiment mass jg at the low-iass cud of collision cascade: where For ο«hy. Ida&Nakazawa(19890). found that the dimneusiouless collision rate for e«Ay is given by (Pai=Ἡδν AR. where the coefficient is determined by Inabaetal. (2001). ," and \ref{eq:e_f}) ) under the Stokes regime, we have the fragment mass $m_{\rm f}$ at the low-mass end of collision cascade: where For $e \ll h_{\rm M}$, \citet{ida89} found that the dimensionless collision rate for $e
+\ll h_{\rm M}$ is given by $\langle P_{\rm col,low} \rangle = 11.3 \sqrt{\tilde R}$ , where the coefficient is determined by \citet{inaba01}. ."
+Since the atimosphere oeffectivelv enhances au enibrvo radius for the aceretion of bodies. the collision rate is modified to be (Iuaba&Tsoma Weobtain frou: Eqs. (12))," Since the atmosphere effectively enhances an embryo radius for the accretion of bodies, the collision rate is modified to be \citep{inaba_ikoma03}
+ Weobtain from Eqs. \ref{eq:dM_ori}) )"
+ and (30)) as, and \ref{eq:pcol_low}) ) as
+The detection of primordial anisotropies in the microwave background. radiation by the CODI satellite (Smoot 1992) has had an enormous impact on cosmology (see White. Scott and Silk 1994 and. Boned 1996 for reviews).,"The detection of primordial anisotropies in the microwave background radiation by the COBE satellite (Smoot 1992) has had an enormous impact on cosmology (see White, Scott and Silk 1994 and Bond 1996 for reviews)."
+" However. the relatively poor angular. resolution of COBL/DAIR (Fret,7Y ) limits the amount of information that can be extracted [from the CAIB."," However, the relatively poor angular resolution of COBE/DMR $\theta_{fwhm} \approx 7^\circ$ ) limits the amount of information that can be extracted from the CMB."
+ From the 4. vear CODE maps (Bennett 1996a). the overall amplitude of the CAIB power spectrum for a given spectral shape has been determined to an accuracy of 14 and a power law index characterizing the shape to £0.24 (Bond 1996).," From the 4 year COBE maps (Bennett 1996a), the overall amplitude of the CMB power spectrum for a given spectral shape has been determined to an accuracy of $7 \%$ and a power law index characterizing the shape to $\pm 0.24$ (Bond 1996)."
+ Constraints on other paranicters such as the spatial curvature ancl the cosmological constant A are weak., Constraints on other parameters such as the spatial curvature and the cosmological constant $\Lambda$ are weak.
+" 1t has long been known (Sunvaev and Zeldovich. 1970) that at angular resolutions smaller than 17017 (the angle subtenced by the Hubble. radius at. the. time of recombination) the CMD. power spectrum. will depend. on the sound speed of the barvon-photon Εις, and hence on a number of fundamental cosmological parameters. such as the densities of barvons. cold. and. hot. dark. matter. and the spatial curvature of the Universe."," It has long been known (Sunyaev and Zeldovich 1970) that at angular resolutions smaller than $\sim 1^\circ \Omega_0^{1/2}$ (the angle subtended by the Hubble radius at the time of recombination) the CMB power spectrum will depend on the sound speed of the baryon-photon fluid, and hence on a number of fundamental cosmological parameters, such as the densities of baryons, cold and hot dark matter, and the spatial curvature of the Universe."
+ In. adiabatic models. the acoustic motions of the matter radiation Iluid load. to a characteristic series of “Doppler peaks’ in the CAB power spectrum which have been investigated. in considerable detail numerically ancl semi-analvticallv ond 1996. Llu 1997).," In adiabatic models, the acoustic motions of the matter radiation fluid lead to a characteristic series of `Doppler peaks' in the CMB power spectrum which have been investigated in considerable detail numerically and semi-analytically Bond 1996, Hu 1997)."
+ Similar behaviour is expected qualitatively in defect (isocurvature) theories. though the pattern. of Doppler peaks is expected to be less. distinct and has not vet been calculated: to high. precision Turok 1996).," Similar behaviour is expected qualitatively in defect (isocurvature) theories, though the pattern of Doppler peaks is expected to be less distinct and has not yet been calculated to high precision Turok 1996)."
+ We therefore restrict the discussion in this paper to purely adiabatic perturbations obeving Ciaussian statistics. as expected in most inllationary models of the carly universe Linde 1990).," We therefore restrict the discussion in this paper to purely adiabatic perturbations obeying Gaussian statistics, as expected in most inflationary models of the early universe Linde 1990)."
+ Phe anisotropies in such models can be computed to high accuracy which. aswe will show in Section 2.. is essential for estimating the precision with which cosmological parameters can be determined [rom he CMB.," The anisotropies in such models can be computed to high accuracy which, aswe will show in Section \ref{sec:param}, is essential for estimating the precision with which cosmological parameters can be determined from the CMB."
+ Intermediate angle experiments have detected emperature anisotropies which are consistent with a ximordial origin and in rough agreement with acdiabatic heoryv predictions., Intermediate angle experiments have detected temperature anisotropies which are consistent with a primordial origin and in rough agreement with adiabatic theory predictions.
+" Llowever. the accuracy and: robustness of the results does not vet allow strong conclusions to of εναν, even when experimental results are combined ogether Zoned 1996. Bond Jalle 1997. Lineweaver 1997. Rocha Lancock 1997)."," However, the accuracy and robustness of the results does not yet allow strong conclusions to be drawn, even when experimental results are combined together Bond 1996, Bond Jaffe 1997, Lineweaver 1997, Rocha Hancock 1997)."
+" Xn experiment with an angular resolution of 85,4,5 can vield. useful information about the CMD Cy up to multipoles bevone a Gaussian filtering scale (ντvN1n30,1,uin~ 2000."," An experiment with an angular resolution of $\theta_{fwhm}
+\sim 5^\prime$ can yield useful information about the CMB ${\rm C}_\ell$ up to multipoles beyond a Gaussian filtering scale $\ell_s =
+\sqrt{8\ln2}\theta_{fwhm}^{-1} \sim 2000$ ."
+synchrotron model provides the most natural explanation for the properties of the extended emission. supporting this preassumption. (,"synchrotron model provides the most natural explanation for the properties of the extended emission, supporting this preassumption. ("
+2) The synchrotron photons are assumed to be IC scattered at most once. Le.. neglecting the second-order IC scattering.,"2) The synchrotron photons are assumed to be IC scattered at most once, i.e., neglecting the second-order IC scattering."
+ This is generally available since the electrons in early afterglow are very energetic. their further interactions with up-seattered synchrotron photons take place in deep Klein-Nishina (KN) regime.," This is generally available since the electrons in early afterglow are very energetic, their further interactions with up-scattered synchrotron photons take place in deep Klein-Nishina (KN) regime."
+ The other minor preassumptions about the afterglow parameters are below. (, The other minor preassumptions about the afterglow parameters are below. (
+"3) The postshock (downstream) electrons! energy equipartition parameter is e,20.1.",3) The postshock (downstream) electrons' energy equipartition parameter is $\eps_e\ga0.1$.
+ This value is inferred from afterglow modelling. and is consistent with. the clustering of explosion energy (Frailetal.2001) and X-ray afterglow luminosity (Freedman&Waxman2001).. (," This value is inferred from afterglow modelling, and is consistent with the clustering of explosion energy \citep{frail01} and X-ray afterglow luminosity \citep{fw01}. ("
+4) The postshock magnetic field energy equipartition parameter is in the range of 107<eg107.,4) The postshock magnetic field energy equipartition parameter is in the range of $10^{-5}\la\eps_B\la10^{-2}$.
+ In the afterglow modelling. values of 107XegὉ are usually obtained (e.g..Panaitescu&Kumar2001:Harrisonetal.2001).," In the afterglow modelling, values of $10^{-3}\la\eps_B\la10^{-2}$ are usually obtained \citep[e.g.,][]{pk01,harrison01}."
+" Moreover. the X-ray afterglow observations indicate B,>0.2, “mG. with n the medium density (LW06)."," Moreover, the X-ray afterglow observations indicate $B_u>0.2n_0^{5/8}$ mG, with $n$ the medium density (LW06)."
+" An lower limit to the downstream magnetic field strength is the shock-compressed field. BY,comp=40B,,."," An lower limit to the downstream magnetic field strength is the shock-compressed field, $B'_{d,\rm
+comp}=4\Gamma B_u$."
+ Here [is the Lorentz factor of postshock plasma and the prime denotes the quantities in the frame of the postshock plasma., Here $\Gamma$ is the Lorentz factor of postshock plasma and the prime denotes the quantities in the frame of the postshock plasma.
+ Thus the lower limit to ερ AS €pcomp=Becomp32sTnmyc?-B;2znmyc.," Thus the lower limit to $\eps_B$ is $\eps_{B,\rm comp}=B_{d,\rm comp}'^2/32\pi
+\Gamma^2nm_pc^2=B_u^2/2\pi nm_pc^2$."
+" With the lower limit to B, from X-ray observations. one has (We have also discussed the case of much lower eg. which does not change much the final constraints: see refsec:Bulimit))."," With the lower limit to $B_u$ from X-ray observations, one has (We have also discussed the case of much lower $\eps_B$, which does not change much the final constraints; see \\ref{sec:Bulimit}) )."
+" With these assumed values of e, and ep. we actually also have e,>eg. which is consistent with the IC component observed in some afterglows (e.g..GRBHarrisonetal.2001 ).. ("," With these assumed values of $\eps_e$ and $\eps_B$, we actually also have $\eps_e>\eps_B$, which is consistent with the IC component observed in some afterglows \citep[e.g., GRB 000926;][]{harrison01}. ("
+5) The medium density is assumed to be 107 IOem.,"5) The medium density is assumed to be $10^{-2}\la n\la10^2\rm
+cm^{-3}$ ."
+" The 5 value is not sure but this adopted range is consistent with afterglow modelling (e.g..Panaitescu&Ku-mar2001:Harrisonetal.2001) and with our knowledge of the interstellar medium, ("," The $n$ value is not sure but this adopted range is consistent with afterglow modelling \citep[e.g.,][]{pk01,harrison01} and with our knowledge of the interstellar medium. ("
+"6) The postshock injected electrons follow a energy distribution of dn,/d,X5,"" with the power-law index assumed to be 2«p=2.5.",6) The postshock injected electrons follow a energy distribution of $dn_e/d\gamma_e\propto\gamma_e^{-p}$ with the power-law index assumed to be $2<p\la2.5$.
+ These values are usually implied by the afterglow observations (e.g..Waxman1997a:Freed-man&Waxman2001.andreferences therein).. and are consistent with the theoretical value pz2.2 derived for isotropic diffusion of accelerated particles (1n the test particle limit) in both numerical calculations (e.g..Achterberget2001) and analytical analysis (Keshet&Waxman2005).," These values are usually implied by the afterglow observations \citep[e.g.,][and references
+therein]{W97a,fw01}, and are consistent with the theoretical value $p\approx2.2$ derived for isotropic diffusion of accelerated particles (in the test particle limit) in both numerical calculations \citep[e.g.,][]{acht01} and analytical analysis \citep{kw05}."
+ In the diffusive (Fermi) shock acceleration mechanism. high-energy charged particles are scattered by upstream and downstream magnetic field back and forth. respectively. crossing the shock front multiples times. and gradually gain energy in each crossing.," In the diffusive (Fermi) shock acceleration mechanism, high-energy charged particles are scattered by upstream and downstream magnetic field back and forth, respectively, crossing the shock front multiples times, and gradually gain energy in each crossing."
+ The particle acceleration depends on the upstream and downstream magnetic fields., The particle acceleration depends on the upstream and downstream magnetic fields.
+ We refer readers to LW06 for detailed discussion and the formula used below., We refer readers to LW06 for detailed discussion and the formula used below.
+" For a relativistic shock expanding into the medium. with postshock plasma ofLorentz factor [ss| and preshock magnetic field B,. the acceleration time of electrons to energy ~tees ds derived to be where the correction factor g accounts for that the upstream electrons are not deflected by an angle of 1/T just with a fraction 1/T of the Larmor time. because the preshock field structure may not be uniform (e.g.Lemoine&Pel-letier2003:Lemome&Revenu2006) and that the electrons do not return the downstream region once they are deflected by ΤΕ. because the shock velocity 1s decreasing (e.g..Piran&Nakar2010)."," For a relativistic shock expanding into the medium, with postshock plasma ofLorentz factor $\Gamma\gg1$ and preshock magnetic field $B_u$, the acceleration time of electrons to energy $\gamma_em_ec^2$ is derived to be where the correction factor $g$ accounts for that the upstream electrons are not deflected by an angle of $1/\Gamma$ just with a fraction $1/\Gamma$ of the Larmor time, because the preshock field structure may not be uniform \citep[e.g.,][]{lp03,lr06} and that the electrons do not return the downstream region once they are deflected by $1/\Gamma$, because the shock velocity is decreasing \citep[e.g.,][]{pn10}."
+. A conservative value is adopted às e~10 (LW06)., A conservative value is adopted as $g\approx10$ (LW06).
+ Hereafter. primed variables denote parameter values measured in the downstream frame. while non-primed ones in the upstream frame.," Hereafter, primed variables denote parameter values measured in the downstream frame, while non-primed ones in the upstream frame."
+ When an electron resides in the upstream region. it may suffer from energy loss due to IC scattering the synchrotron photons from the downstream region.," When an electron resides in the upstream region, it may suffer from energy loss due to IC scattering the synchrotron photons from the downstream region."
+ LWOG derived the IC cooling time of an upstream-residing electron with ος. measured in the downstream region. is Here Tn is the synchrotron radiation. energy density measured in the downstream frame.," LW06 derived the IC cooling time of an upstream-residing electron with $\gamma_e$, measured in the downstream region, is Here $u_{\rm ph}'$ is the synchrotron radiation energy density measured in the downstream frame."
+ Because the electrons are confined in a region with thickness much smaller than the size of the shock Κ. the synchrotron radiation energy densitiesin the upstream and downstream regions that we concern are similar. thus Tn can be either of them.," Because the electrons are confined in a region with thickness much smaller than the size of the shock $R$, the synchrotron radiation energy densitiesin the upstream and downstream regions that we concern are similar, thus $u_{\rm ph}'$ can be either of them."
+" The IC scattering may occur in KN regime. then 6ph should be considered as the ""effective"" energy density of scattered photons."," The IC scattering may occur in KN regime, then $u_{\rm ph}'$ should be considered as the ""effective"" energy density of scattered photons."
+ Let us define the photon frequency so that the IC scattering between electrons with (downstream-frame) Lorentz factor +f and photons at v>VSN enters KN regime., Let us define the photon frequency so that the IC scattering between electrons with (downstream-frame) Lorentz factor $\gamma_x'$ and photons at $\nu>\nukn_x$ enters KN regime.
+" Except for the spectrum of the seed photons are very hard (which ts not the case here). an electron of >, mainly interacts with photons below ΕΝ,"," Except for the spectrum of the seed photons are very hard (which is not the case here), an electron of $\gamma_x$ mainly interacts with photons below $\nukn_x$."
+ So if KN effect is important. the photon energy density 1n eq.(3)) should be replaced by that below ολοι). Tn»TESMN).," So if KN effect is important, the photon energy density in \ref{eq:ictime}) ) should be replaced by that below $\nukn(\gamma_e)$ , $u_{\rm ph}'\rightarrow u_{\rm ph}'(<\nukn)$."
+" An electron that has been accelerated to ο must satisfy that its upstream IC cooling time is longer than the acceleration time. 7> £,/T. which gives an upper limit to 77. For the following derivation. it is convenient to define the ""effective"" downstreamCompton parameter Y, for relevant electrons of 5, as with a,=B7/85 being the energy density of postshock magnetic field. and further replace Te1 in eq.(5)) with Yu."," An electron that has been accelerated to $\gamma_em_ec^2$ must satisfy that its upstream IC cooling time is longer than the acceleration time, $t_c'>t_a/\Gamma$ , which gives an upper limit to $\gamma_e'$ , For the following derivation, it is convenient to define the ""effective"" downstreamCompton parameter $Y_x$ for relevant electrons of $\gamma_x$ as with $u_B'=B_d'^2/8\pi$ being the energy density of postshock magnetic field, and further replace $u_{\rm ph}'$ in \ref{eq:gammamaxic}) ) with $u_{\rm
+ph}'(<\nukn)=Yu'_B$ ."
+via bipolar jets.,via bipolar jets.
+ The syncehrotron spectrum is cut. olf at wavelengths shorter than£j., The synchrotron spectrum is cut off at wavelengths shorter than.
+ As noted in section. 1. many cdillerent. classes of X-ray binaries emit svnchrotron radiation (lender 2003).," As noted in section 1, many different classes of X-ray binaries emit synchrotron radiation (Fender 2003)."
+ Rupen et al. (, Rupen et al. (
+2002) identified a weak (0.35£0.07 mJv) 4.86 ο radio source at the SAN JO929-314 position on 2002 May 3 and 7.,2002) identified a weak $\sim 0.35 \pm 0.07$ mJy) 4.86 GHz radio source at the SAX J0929-314 position on 2002 May 3 and 7.
+ Unfortunately there do not appear to have been any Ht observations during this outburst., Unfortunately there do not appear to have been any IR observations during this outburst.
+ We hypothesise that a rather similar phenomenon may have occurred during the 1998 outburst of the acereting millisecond pulsar SAN JISOS.43658., We hypothesise that a rather similar phenomenon may have occurred during the 1998 outburst of the accreting millisecond pulsar SAX J1808.4–3658.
+ Wang et al (2001) noted that the band Iluxes measured by the JI. 12m telescope on 1998. April 18.2 were about 0.2 magnitudes brighter than they were  0.5 days later when measured by the Mt Canopus 12m telescope., Wang et al (2001) noted that the band fluxes measured by the JKT 1-m telescope on 1998 April 18.2 were about 0.2 magnitudes brighter than they were $\sim$ 0.5 days later when measured by the Mt Canopus 1-m telescope.
+ μον assumed that the discrepancy was due to calibration uncertainties in the JP cata., They assumed that the discrepancy was due to calibration uncertainties in the JKT data.
+ Phere was a clear IR. (AHN) excess measured by the UINIIE telescope at 1998 April 18.6 just 0.4 clavs after the JINT measurements., There was a clear IR ) excess measured by the UKIRT telescope at 1998 April 18.6 just 0.4 days after the JKT measurements.
+ Wang et al. (, Wang et al. (
+2001) proposed a synchrotron origin for this Lt excess.,2001) proposed a synchrotron origin for this IR excess.
+ We sugeest that the svnchrotron excess was also present at the time of the JIT measurements and that it extended into the optical bands., We suggest that the synchrotron excess was also present at the time of the JKT measurements and that it extended into the optical bands.
+ It had disappeared at the time of the Mt Canopus measurements a few hours later., It had disappeared at the time of the Mt Canopus measurements a few hours later.
+ The optical counterpart of NTE JO929314 was variable on all timescales down to a few hours during the 2002 May observations., The optical counterpart of XTE J0929–314 was variable on all timescales down to a few hours during the 2002 May observations.
+ On one occasion lasting 4 hours the band Hux was modulated at the orbital period with amplitude 0.09+0.01 magnitudes., On one occasion lasting $\sim 4$ hours the band flux was modulated at the orbital period with amplitude $0.09 \pm 0.01$ magnitudes.
+ No variability was apparent in the X-rav measurements (Calloway ct al. 2002b).," No variability was apparent in the X-ray measurements (Galloway et al, 2002b)."
+ Phe peak of the orbital modulation occurs at a phase of 0.1940.05 relative to the X-ray ephemeris ancl appears to rule out N-rav. heating of the companion as the source of the modulation unless it is combined with emission from a hot spot on the disc., The peak of the orbital modulation occurs at a phase of $0.19 \pm 0.05$ relative to the X-ray ephemeris and appears to rule out X-ray heating of the companion as the source of the modulation unless it is combined with emission from a hot spot on the disc.
+ Broad band spectra taken on S. nights have an approximately power law distribution as expected for an optically thick accretion cise but with variable excesses in., Broad band spectra taken on 8 nights have an approximately power law distribution as expected for an optically thick accretion disc but with variable excesses in.
+1.2 Overall these excesses declined: ancl the spectra steepenec (became blucr) during the period of the observations., Overall these excesses declined and the spectra steepened (became bluer) during the period of the observations.
+" While variable 4, emission may be responsible for some of the excess inZ?. another explanantion is required for the band enhancements."," While variable $H_\alpha$ emission may be responsible for some of the excess in, another explanantion is required for the band enhancements."
+ We suggest. they may be due to emission from cool matter in the outer part of the dise following a transient. episode of mass transfer from the companion., We suggest they may be due to emission from cool matter in the outer part of the disc following a transient episode of mass transfer from the companion.
+ Alternatively. variable svnchrotron emission. cut olf at band wavelengths. contributes to the emission spectrum.," Alternatively, variable synchrotron emission, cut off at band wavelengths, contributes to the emission spectrum."
+ There is a clear need for [ast follow-up optical. LR ancl racio observations of millisecond. X-ray. pulsars.," There is a clear need for fast follow-up optical, IR and radio observations of millisecond X-ray pulsars."
+ These should include polarimetry and high time resolution optical and near LR: photometry to test the svnchrotron emission hypothesis., These should include polarimetry and high time resolution optical and near IR photometry to test the synchrotron emission hypothesis.
+ The feasibility of high5 speed photometry at the pulsar spin frequency. might. also be investigated., The feasibility of high speed photometry at the pulsar spin frequency might also be investigated.
+ Facilities for immediate data reduction and 5generation of light5 curves are essential in. order to optimise observing strategies [or these highly: variable objects., Facilities for immediate data reduction and generation of light curves are essential in order to optimise observing strategies for these highly variable objects.
+ We thank Ron Remillard for timely information on the occurrence of this transient and Duncan Calloway for providing PCA data., We thank Ron Remillard for timely information on the occurrence of this transient and Duncan Galloway for providing PCA data.
+ This research. has macle use of data obtained through the High Enerey Astrophysics Science Archive Research Center Online Service. provided by the NASA Goddard: Space Flight Center.," This research has made use of data obtained through the High Energy Astrophysics Science Archive Research Center Online Service, provided by the NASA / Goddard Space Flight Center."
+ We thank Don Alelrose. Mark Walker and the referee for helpful comments and gratefully acknowledge financial support for the Mt Canopus Observatory by Mr. David. Warren.," We thank Don Melrose, Mark Walker and the referee for helpful comments and gratefully acknowledge financial support for the Mt Canopus Observatory by Mr David Warren."
+ ABC thanks the University of Tasmania Antarctic CRC for the use of computer facilities., ABG thanks the University of Tasmania Antarctic CRC for the use of computer facilities.
+where m is the mass of the protostar (here assumed equal to the sink particle mass).,where $m$ is the mass of the protostar (here assumed equal to the sink particle mass).
+ The eas temperature around the voung stars is set to be the maximum of the temperature from either the barotropic equation of state or the radiative feedback., The gas temperature around the young stars is set to be the maximum of the temperature from either the barotropic equation of state or the radiative feedback.
+ his ensures a maximal elfect from the raciation., This ensures a maximal effect from the radiation.
+ To study the local environment of the proto-stars. we make maps of the column density of the core of gas surrounding the sink.," To study the local environment of the proto-stars, we make maps of the column density of the core of gas surrounding the sink."
+ These are mace from the perspective of the central proto-star looking outward., These are made from the perspective of the central proto-star looking outward.
+ We define the core to be all material within 7.=0.01 pe of the central sink., We define the core to be all material within $R_c = 0.01$ pc of the central sink.
+ In order to ensure that our core sample has a constant volume we do not use a clumplinding approach., In order to ensure that our core sample has a constant volume we do not use a clumpfinding approach.
+ This ensures that cilferences in the column densities are due to internal density enhancements. and not due to cdillering radial extents.," This ensures that differences in the column densities are due to internal density enhancements, and not due to differing radial extents."
+ In ος1200 it was found that the typical radius within which the bound material of a pre-stellar core was contained. was r=116.10? pe., In SCB09 it was found that the typical radius within which the bound material of a pre-stellar core was contained was $r=1.16\E^{-2}$ pc.
+ Therefore our selected size of 0.01. pe represents the volume which will first collapse., Therefore our selected size of $0.01$ pc represents the volume which will first collapse.
+ Although the Class 0 classification scheme is based on low mass YSOs. we do form high mass stars in this simulation.," Although the Class 0 classification scheme is based on low mass YSOs, we do form high mass stars in this simulation."
+ However. as shown in SLD09 there are no high mass pre-stellar cores in this simulation.," However, as shown in SLB09 there are no high mass pre-stellar cores in this simulation."
+" Fherefore we use the term ""Class 0 to refer to all of our collapsing cores.", Therefore we use the term `Class 0' to refer to all of our collapsing cores.
+ We use Hammer projections to show the column density surface., We use Hammer projections to show the column density surface.
+ In this projection. a point at (A.0) on a spherical surface will correspond to position Gr. on the projected Hat surface. according to the transformation:y) where This projection is not conformal. but area is conserved. which is essential for the following calculation of column density.," In this projection, a point at $(\lambda,\phi)$ on a spherical surface will correspond to position $(x,y)$ on the projected flat surface, according to the transformation; where This projection is not conformal, but area is conserved, which is essential for the following calculation of column density."
+ The first stage to make the Llamumner projections is to find the particles within the core volume for every sink just after it has formed., The first stage to make the Hammer projections is to find the particles within the core volume for every sink just after it has formed.
+" Cores that are within 2/2, of an existing sink are rejected.", Cores that are within $2 R_c$ of an existing sink are rejected.
+ In addition. to ensure that we have sullicient resolution. we only use cores with over 500 SPILL particles.," In addition, to ensure that we have sufficient resolution, we only use cores with over $500$ SPH particles."
+ This is 10 times higher than our minimum resolution of 50 particles., This is $10$ times higher than our minimum resolution of $50$ particles.
+ In. practice. the cores. generally have ~1500 particles We select. lines of sight [rom a uniform. (50 100) erid on the Hammer surface. in order to ensure equal area sampling.," In practice, the cores generally have $\sim1500$ particles We select lines of sight from a uniform $50\times100$ ) grid on the Hammer surface, in order to ensure equal area sampling."
+ This equates to 3895 lines of sieht within the projected elliptical surface., This equates to $3895$ lines of sight within the projected elliptical surface.
+ We calculate the density at cach point along the line of sight using the standard SPLL density equation. where m; is the mass of the jth particle. Wrrj.5) is the smoothing kernel and I is the number of particles within a smoothing length. h.," We calculate the density at each point along the line of sight using the standard SPH density equation, where $m_{j}$ is the mass of the $j$ 'th particle, $W(\mathbf{r}-\mathbf{r_{j}},h)$ is the smoothing kernel and $N$ is the number of particles within a smoothing length, $h$."
+ Phe densities are then integrated along the line of sight to obtain the column density., The densities are then integrated along the line of sight to obtain the column density.
+ A ram pressure map of the core can be made in a similar manner. by integrating the product of density and. velocity along the line of sight.," A ram pressure map of the core can be made in a similar manner, by integrating the product of density and velocity along the line of sight."
+ ‘To categorise the column density. surfaces. the angular covariance of the high density gas is calculated for cach core projection.," To categorise the column density surfaces, the angular covariance of the high density gas is calculated for each core projection."
+ “Phe covariance is where AI is the mean column density and where q;;=NiNj is the product of the cobumn density. AN. ab points / and j which have an angular separation in the range dé. and n is the number of pairs 7j separated. by d.," The covariance is where $\bar{N}$ is the mean column density and where $q_{ij}=N_i*N_j$ is the product of the column density, $N$ , at points $i$ and $j$ which have an angular separation in the range $d\theta$, and $n$ is the number of pairs $ij$ separated by $d\theta$."
+ Vhe covariance is calculated at [ive degree intervals. as this allows each bin to be well filled while still sampling the angular range finely enough to see small-scale structure.," The covariance is calculated at five degree intervals, as this allows each bin to be well filled while still sampling the angular range finely enough to see small-scale structure."
+ If the column clensity clistribution is uniform. then we expect cor(d@) to be zero everywhere as the column densities will be equal to the average.," If the column density distribution is uniform, then we expect $cov(d\theta)$ to be zero everywhere as the column densities will be equal to the average."
+ However. if there is structure in the column density map with a particular angular scale. [or example in the form of ten degree wide blobs. then the value of coe(d?) will be higher at this angular separation as the structure is highly correlated on this scale.," However, if there is structure in the column density map with a particular angular scale, for example in the form of ten degree wide blobs, then the value of $cov(d\theta)$ will be higher at this angular separation as the structure is highly correlated on this scale."
+ The column density projections of four representative cores are shown in LL., The column density projections of four representative cores are shown in \ref{CDprojections}.
+ Core (a) has a fairly uniform column density., Core (a) has a fairly uniform column density.
+ However. the other three cores have at loast one xeferred direction where the column density is higher.," However, the other three cores have at least one preferred direction where the column density is higher."
+ This distribution is consistent with the high density material being distributed. along a Lilament., This distribution is consistent with the high density material being distributed along a filament.
+ LW the sink xwticle is situated at the end. of the filament. then the column density is higher in the direction of the filament. as in core (b).," If the sink particle is situated at the end of the filament, then the column density is higher in the direction of the filament, as in core (b)."
+ If the sink particle is located within the filament. hen the column density peaks where the filament intersects he core volume. and falls olf perpendicular to the filament.," If the sink particle is located within the filament, then the column density peaks where the filament intersects the core volume, and falls off perpendicular to the filament."
+ Cores (c) and. (d) represent this situation., Cores (c) and (d) represent this situation.
+ Core (c) has a ilament which is denser on one side than the other. and core (d) has a filament which is equally dense on both sides. showing that the sink is located at the centre of the filament.," Core (c) has a filament which is denser on one side than the other, and core (d) has a filament which is equally dense on both sides, showing that the sink is located at the centre of the filament."
+ Phe column density seen by the central sink is almost entirely due to the core itself., The column density seen by the central sink is almost entirely due to the core itself.
+ 1. is caleulated. using only the material within 0.01 pe ofthe central sink. but when," \ref{CDprojections} is calculated using only the material within $0.01$ pc ofthe central sink, but when"
+"The IRAS Bright Galaxy Survey (hereafter BIRG survey) by Soifer.etal.(1989). and by Sandersetal.(1995) (southern extension) includes all the galaxies brighter than 5.4 Jv αἱ 60jan. The IRAS Bright galaxies are therefore. by definition. the brightest extragalactic objects in the sky al 6050. From this survey we learned a wealth of astrophysical information: (1) the far IR. (FIR) emission dominates the total luminosity in a significant fraction of ealaxies: (2) at luminosity loe(L;,./£.)Z11 (the so-called Luminous IInrared Galaxies. LIRGs). IR. selected. galaxies become more numerous (han optically selected Starburst ancl Sevlert galaxies of comparable bolometric luminosity.","The IRAS Bright Galaxy Survey (hereafter BIRG survey) by \citet{s89} and by \citet{s95} (southern extension) includes all the galaxies brighter than $5.4$ Jy at $60 \mu$ m. The IRAS Bright galaxies are therefore, by definition, the brightest extragalactic objects in the sky at $60 \mu$ m. From this survey we learned a wealth of astrophysical information: (1) the far IR (FIR) emission dominates the total luminosity in a significant fraction of galaxies; (2) at luminosity $\log(L_{ir}/ L_\odot)
+\simgt 11$ (the so-called Luminous Infrared Galaxies, LIRGs), IR selected galaxies become more numerous than optically selected Starburst and Seyfert galaxies of comparable bolometric luminosity."
+" At luminosity log(L;,/L..)Z12. the so-called Ultra-Luminous Infrared Galaxies (ULIRGs) exceed the space densities of QSOs by a [actor of 1.5-2 (Sanders&Mirabel(1996):: Sanders.Stace.&Ishida (1999)))."," At luminosity $\log(L_{ir}/ L_\odot) \simgt 12$, the so-called Ultra-Luminous Infrared Galaxies (ULIRGs) exceed the space densities of QSOs by a factor of 1.5-2 \citet{sm96}; \citet{ssi99}) )."
+ A considerable number of studies suggest a strong relation between galaxy interactions and the hiehest IR. luminosity., A considerable number of studies suggest a strong relation between galaxy interactions and the highest IR luminosity.
+ ULIBGs are often found to be interacting/mergine svstenis (Sanders.Surace.&Ishicla (1999)))., ULIRGs are often found to be interacting/merging systems \citet{ssi99}) ).
+" However. the environment of moderately luminous infrared galaxies (LOML.<Lyi,x10HL.. hereafter MIRGs). and luminous infrared galaxies is nol well known vet."," However, the environment of moderately luminous infrared galaxies $10^{10}L_{\odot} \leq L_{fir} \leq
+10^{11}L_{\odot}$, hereafter MIRGs), and luminous infrared galaxies is not well known yet."
+" In (his paper. we study the circumgalactie environment of 87 galaxies from the DIRG Survey. with luminosity range 10!'""L. Lig 107L..."," In this paper, we study the circumgalactic environment of 87 galaxies from the BIRG Survey, with luminosity range $10^{10}$ $\leq $ $\leq 10^{12}$."
+ The sample is composed by MIRGs and a few LIRGs., The sample is composed by MIRGs and a few LIRGs.
+ We also consider whether a correlation may be present between FIR. properties and the projected separation of the DIRG and its nearest companion (84.6))., We also consider whether a correlation may be present between FIR properties and the projected separation of the BIRG and its nearest companion \ref{corr}) ).
+ We then compare the BIRG environment {ο the one of Sevlert 1 and 2 ealaxies (85.2))., We then compare the BIRG environment to the one of Seyfert 1 and 2 galaxies \ref{seyf}) ).
+ Finally. we discuss the implications of interaction-induced eenhancement for secular evolution of galaxies and for the relationship between starbursting and active galaxies (85)).," Finally, we discuss the implications of interaction-induced enhancement for secular evolution of galaxies and for the relationship between starbursting and active galaxies \ref{disc}) )."
+ In the following discussion. we adopt a IIubble constant of |.," In the following discussion, we adopt a Hubble constant of $H_o =
+75~Km~s^{-1}~Mpc^{-1}$ ."
+ The Bright IRAS sample consists of 87 objects. and was compiled from the IRLAS Bright Galaxy Survey bv Soiler.etal.(1989). for the Northern hemisphere. aud by (1995) for the Southern one.," The Bright IRAS sample consists of 87 objects, and was compiled from the IRAS Bright Galaxy Survey by \citet{s89} for the Northern hemisphere, and by \citet{s95} for the Southern one."
+ All objects with galactic latitude [δη>30° wwere included in (he sample., All objects with galactic latitude $|b_{II}| \geq 30^{\circ}$ were included in the sample.
+ In (his way. we avoid sampling the galactic plane. where a bias in the detection of companions is expected. because of both absorption and crowding.," In this way we avoid sampling the galactic plane, where a bias in the detection of companions is expected because of both absorption and crowding."
+ We further restrict our selection to a volune-limited sample (redshift range 0.008<2 0.013).," We further restrict our selection to a volume-limited sample (redshift range $0.008 \leq z \leq
+0.018$ )."
+" A V/V, {test (SchmidtLOGS) gives a value of 0.47 4 0.05 (rms).", A $V/V_{max}$ test \citep{schmidt68} gives a value of 0.47 $\pm$ 0.05 (rms).
+ Sincethe DIRG Survey is highly, Sincethe BIRG Survey is highly
+(Siwak et al.,(Siwak et al.
+ 2010)., 2010).
+ Furthermore. a third light source (£4 ) was conside‘edd throughout the analyses.," Furthermore, a third light source $\ell_3$ ) was considered throughout the analyses."
+ This procedure showed acceptable photometric soltions only [or conact mode 3 auc indicated a somewlat broad range of 0.56 €cUq 0.6L where the fil-out [actor «ecreases toward larger mass ratios f‘oll  at q—0.56 to 5 at qc—0.61.," This procedure showed acceptable photometric solutions only for contact mode 3 and indicated a somewhat broad range of 0.56 $\le q \le$ 0.64, where the fill-out factor decreases toward larger mass ratios from 8 at $q$ =0.56 to 5 at $q$ =0.64."
+ But. the trials for a yossible {η failed to achieve convergeice.," But, the trials for a possible $\ell_3$ failed to achieve convergence."
+ lu he secoud stage. the previously determine parameters aud tje Inass ratio from Siwak οἱ al. (," In the second stage, the previously determined parameters and the mass ratio from Siwak et al. ("
+2010 were secl as the initial values.,2010) were used as the initial values.
+ Cousicerit& its effective temperature with au accuracy no better ian 1007200 Ix. the envelopes of the primary star should lie close to the lower unit between the raclaive ald convective atmospheres. so Ay atd gy were included as additional [ree variables.," Considering its effective temperature with an accuracy no better than $\sim$ 200 K, the envelopes of the primary star should lie close to the lower limit between the radiative and convective atmospheres, so $A_1$ and $g_1$ were included as additional free variables."
+ Final resits are given in Table 2 together with those of Siwak et al. (, Final results are given in Table 2 together with those of Siwak et al. (
+2010) for coiparison aud plottec i Figu'e 2. where. for clarity. individual observatiois have beeu compiled ito LOO inean poluts Ising bii widths of 0.01 in phase for each filtered light. οἱrve.,"2010) for comparison and plotted in Figure 2, where, for clarity, individual observations have been compiled into 100 mean points using bin widths of 0.01 in phase for each filtered light curve."
+ As seen in te ligure. the COLUDIUec| light curves describe the SOAO iultiband data ¢ute well.," As seen in the figure, the computed light curves describe the SOAO multiband data quite well."
+ Qi-— 'esult represeus the system as a contac binary with a fill-ot factor of LS aud with a COLSI‘ably large temperature cilference of LLI7 Ex. This itpies hat both component. stars are ia“aej1ially over-coLtact with respect to the iujer Roche obe axc te thermal contact between them is poor., Our result represents the system as a contact binary with a fill-out factor of 4.8 and with a considerably large temperature difference of 1447 K. This implies that both component stars are marginally over-contact with respect to the inner Roche lobe and the thermal contact between them is poor.
+ But contrarily. Siwak et al. (," But contrarily, Siwak et al. ("
+2010 reported WZ Cye ο be a NCB in which t1e two Conponenis are a or near their lobes.,2010) reported WZ Cyg to be a NCB in which the two components are at or near their lobes.
+ From our binary pa‘aleers aud the spectroscopic orbit o .Siwak et al. (, From our binary parameters and the spectroscopic orbit of Siwak et al. (
+2010). we obtained the absoute climensious Ort e system listed in Table 3. assuni18oO that the tem;»erature of each compouen has an error of 20 Ix and that the boometric magnitude of the Suu is Apr. =41-73.,"2010), we obtained the absolute dimensions for the system listed in Table 3, assuming that the temperature of each component has an error of 200 K and that the bolometric magnitude of the Sun is $M_{\rm bol}$$_\odot$ =+4.73."
+ For the absolute visual nagitucles (Αν). We Ised te bolometric corrections (BCs) [roi the scaling between logT aud BC recalculated by Torres (2010) [rom Flowers (1996) table.," For the absolute visual magnitudes $M_{\rm V}$ ), we used the bolometric corrections (BCs) from the scaling between $\log T$ and BC recalculated by Torres (2010) from Flower's (1996) table."
+ We deteined 15 times of ΙΙ light with the weighted means for the timines in each filter wy usiug the method of Iwee van Woerden (1956)., We determined 15 times of minimum light with the weighted means for the timings in each filter by using the method of Kwee van Woerden (1956).
+ Twenty-five additional timines were derived ον Us usiug {ie data [rom the WASP (Wide Anele Search for Planets) publie archive (Butters et al., Twenty-five additional timings were derived by us using the data from the WASP (Wide Angle Search for Planets) public archive (Butters et al.
+ 2010)., 2010).
+ From he database of Ixreiner et al. (, From the database of Kreiner et al. (
+2001) aud from more recent literature. 331 timings (33) shotographic plate. 213 visual. 35 photographie. LL photoelectric. aid 39 CCD) have been added O OUP east'eiuents.,"2001) and from more recent literature, 331 timings (33 photographic plate, 213 visual, 35 photographic, 11 photoelectric, and 39 CCD) have been added to our measurements."
+ Our period study is based ou a total o “S71 times of minimuai light spauniug nore than 112 years., Our period study is based on a total of 371 times of minimum light spanning more than 112 years.
+ All photoelectric and CCD timines are listed i1 Table Ll. wherein the second and third coluuus give tle HJED (Heliocentric Julia1 Ephemeris Dae) timiugs trausformed to the errestrial tit1ο scale (Bastiau 2000) and their uncelaiutles. respectively.," All photoelectric and CCD timings are listed in Table 4, wherein the second and third columns give the HJED (Heliocentric Julian Ephemeris Date) timings transformed to the terrestrial time scale (Bastian 2000) and their uncertainties, respectively."
+ Because alimost all but he CCD timines were ptblished without error infornation. he following staucard deviatious were assigned to the timing residuals based on observational metlod for tie period analysis of WZ Cyve: +£0.0165 d for photograpuc plate. £O.0067 d [or visual. £0.0018 d for photographic. aud 20.0029 for photoelectric minitia.," Because almost all but the CCD timings were published without error information, the following standard deviations were assigned to the timing residuals based on observational method for the period analysis of WZ Cyg: $\pm$ 0.0165 d for photographic plate, $\pm$ 0.0067 d for visual, $\pm$ 0.0048 d for photographic, and $\pm$ 0.0029 d for photoelectric minima."
+ Relative weights were then calculated as the inverse squares of these, Relative weights were then calculated as the inverse squares of these
+"The second order correction to the Jarskog CP invariant is given by The second order contributions to the 79.74 unitarity (triangle are given by The second order contributions to the 1.7, (langle are given by",The second order correction to the Jarskog CP invariant is given by The second order contributions to the $\nu_2.\nu_3$ unitarity triangle are given by The second order contributions to the $\nu_e .\nu_{\mu}$ triangle are given by
+ R. V., R. V.
+for0.2M.. donors will apply to larecr donors as well because the accretion rates deteriined by gravitational wave radiatiou for low-1iass main-sequence douors differ by a factor of a few at most.,for$0.2 \msol$ donors will apply to larger donors as well because the accretion rates determined by gravitational wave radiation for low-mass main-sequence donors differ by a factor of a few at most.
+" For example. the accretion rate for the most massive stably iiass-trausferiug svsteni (Mj=0.33M M=0.5 M.) is only 3 times larger at a given orbital period than ourfiducial example with AI,=0.2M. aud M=0.1AL..."," For example, the accretion rate for the most massive stably mass-transferring system $M_d = 0.33 \msol$, $M = 0.5 \msol$ ) is only 3 times larger at a given orbital period than ourfiducial example with $M_d = 0.2 \msol$ and $M=0.4 \msol$."
+ We neglect any potential variations of AL on the CN. recurrence time due to prolonged periods of “hibernation(Sharaetal.1986)., We neglect any potential variations of $\dot M$ on the CN recurrence time due to prolonged periods of “hibernation”\citep{shara86}.
+. We further uote that the RCM) relation we use assumes a douor with solar metallicity., We further note that the $R_d(M_d)$ relation we use assumes a donor with solar metallicity.
+ As Stehleetal.(1997) aud I&olb&Baratte(1999) note. the secondary’s composition docs have au effect ou its structure and thus ou the secular evolution of the CV.," As \cite{skr97} and \cite{kolbbar99} note, the secondary's composition does have an effect on its structure and thus on the secular evolution of the CV."
+ We have also neelected the zinall correction needed to re-adjust the πι1) relation for our specific AT history. which is lower by a factor of z1.5 than that iu Kolb&Daraffe(1999) due to our 0.1AL. WD (they modeled a 0.6AL. WD).," We have also neglected the small correction needed to re-adjust the $R_d(M_d)$ relation for our specific $\dot M$ history, which is lower by a factor of $\approx 1.5$ than that in \cite{kolbbar99} due to our $0.4 \msol$ WD (they modeled a $0.6 \msol$ WD)."
+ This would teud to slightly shorten the minima orbital period. as the mass-transfer timescale at a given donor mass is longer in our case.," This would tend to slightly shorten the minimum orbital period, as the mass-transfer timescale at a given donor mass is longer in our case."
+ However. given that the current calculations do not match the now observed mium orbital period in either case (Causickeetal.2009).. we will neelect the corrections due to the secondarys composition and AL history.," However, given that the current calculations do not match the now observed minimum orbital period in either case \citep{gaen09}, we will neglect the corrections due to the secondary's composition and $\dot{M}$ history."
+ The previous section showed that the relevant accretion rates are much lower than that considered bv Sharaetal.(1993)., The previous section showed that the relevant accretion rates are much lower than that considered by \cite{shara93}.
+.. Iu addition. bv the time the known pre-CVs come iutocontact and initiate stable mass fransfer. the Πο WD core will have cooled to T.=2?6«10°IK CAlthaus&Beuvennto 1997)..," In addition, by the time the known pre-CVs come intocontact and initiate stable mass transfer, the He WD core will have cooled to $T_c\approx 2-6\times 10^6\ {\rm K}$ \citep{altben97}. ."
+" Extensions of the work by Townsley&Bildsten on CN jenition masses to lower AM. AL aud T; vields Aha,m107AL..."," Extensions of the work by \cite{tb04} on CN ignition masses to lower $M$, $\dot M$ and $T_c$ yields $M_{\rm ign}\approx 10^{-3} \msol$."
+ This implies CN recurrence times of 10°105 yr during the many eigavears of evolution uuder —accretion., This implies CN recurrence times of $10^7-10^8$ yr during the many gigayears of evolution under accretion.
+ The long evolution time allows for sieuificaut thermal coupling between the freshly accreted II euvelope aud the He WD core during the acciunulatiou of fresh matter., The long evolution time allows for significant thermal coupling between the freshly accreted H envelope and the He WD core during the accumulation of fresh matter.
+ For this reason. we must track T; aud its evolution under the action of accretion for a prolonged period of hundreds of CN events.," For this reason, we must track $T_c$ and its evolution under the action of accretion for a prolonged period of hundreds of CN events."
+ Due to the long accumulation aud evolution time. much of the He WD core participates in the thermal balance of the CN evele.," Due to the long accumulation and evolution time, much of the He WD core participates in the thermal balance of the CN cycle."
+" We beein by estimating the timescale for heat transport between two mass shells at radial coordinates ry aud à (Houyey&L'Ecuyer 1969).. where ep is the specific heat. agp, is the Stefau- radiationconstant. and & is the opacity."," We begin by estimating the timescale for heat transport between two mass shells at radial coordinates $r_0$ and $r$ \citep{hl69}, , _d ^r dr ]^2 , where $c_P$ is the specific heat, $a_{\rm SB}$ is the Stefan-Boltzmann radiationconstant, and $\kappa$ is the opacity."
+ Electron coucuction determines leat transfer in the ≼↧↸∖∶↴∙⊾↸∖∐↸∖↥⋅⋜↧↑↸∖∐↸∖↸⊳∪↥⋅↸∖⋅↴∖↴∪∕⋅⊳∶∕⋅⊳⊺−∣≻−⋅↖↖⇁↕∐∖↥⋅↸∖∕⋅," Electron conduction determines heat transfer in the degenerate He core, so $\kappa = \kappa' T^2/\rho^2$ , where"
+ Electron coucuction determines leat transfer in the ≼↧↸∖∶↴∙⊾↸∖∐↸∖↥⋅⋜↧↑↸∖∐↸∖↸⊳∪↥⋅↸∖⋅↴∖↴∪∕⋅⊳∶∕⋅⊳⊺−∣≻−⋅↖↖⇁↕∐∖↥⋅↸∖∕⋅⊳," Electron conduction determines heat transfer in the degenerate He core, so $\kappa = \kappa' T^2/\rho^2$ , where"
+ Electron coucuction determines leat transfer in the ≼↧↸∖∶↴∙⊾↸∖∐↸∖↥⋅⋜↧↑↸∖∐↸∖↸⊳∪↥⋅↸∖⋅↴∖↴∪∕⋅⊳∶∕⋅⊳⊺−∣≻−⋅↖↖⇁↕∐∖↥⋅↸∖∕⋅⊳≈," Electron conduction determines heat transfer in the degenerate He core, so $\kappa = \kappa' T^2/\rho^2$ , where"
+ Electron coucuction determines leat transfer in the ≼↧↸∖∶↴∙⊾↸∖∐↸∖↥⋅⋜↧↑↸∖∐↸∖↸⊳∪↥⋅↸∖⋅↴∖↴∪∕⋅⊳∶∕⋅⊳⊺−∣≻−⋅↖↖⇁↕∐∖↥⋅↸∖∕⋅⊳≈∩," Electron conduction determines heat transfer in the degenerate He core, so $\kappa = \kappa' T^2/\rho^2$ , where"
+ Electron coucuction determines leat transfer in the ≼↧↸∖∶↴∙⊾↸∖∐↸∖↥⋅⋜↧↑↸∖∐↸∖↸⊳∪↥⋅↸∖⋅↴∖↴∪∕⋅⊳∶∕⋅⊳⊺−∣≻−⋅↖↖⇁↕∐∖↥⋅↸∖∕⋅⊳≈∩∖," Electron conduction determines heat transfer in the degenerate He core, so $\kappa = \kappa' T^2/\rho^2$ , where"
+ Electron coucuction determines leat transfer in the ≼↧↸∖∶↴∙⊾↸∖∐↸∖↥⋅⋜↧↑↸∖∐↸∖↸⊳∪↥⋅↸∖⋅↴∖↴∪∕⋅⊳∶∕⋅⊳⊺−∣≻−⋅↖↖⇁↕∐∖↥⋅↸∖∕⋅⊳≈∩∖E," Electron conduction determines heat transfer in the degenerate He core, so $\kappa = \kappa' T^2/\rho^2$ , where"
+ Electron coucuction determines leat transfer in the ≼↧↸∖∶↴∙⊾↸∖∐↸∖↥⋅⋜↧↑↸∖∐↸∖↸⊳∪↥⋅↸∖⋅↴∖↴∪∕⋅⊳∶∕⋅⊳⊺−∣≻−⋅↖↖⇁↕∐∖↥⋅↸∖∕⋅⊳≈∩∖EE," Electron conduction determines heat transfer in the degenerate He core, so $\kappa = \kappa' T^2/\rho^2$ , where"
+Siuce there is no allbsky map at a sufficiently. high. resolution in the NIB. it is cliffieult to calculate the cumulative intensity of stars at varving ealactic locations.,"Since there is no all-sky map at a sufficiently high resolution in the NIR, it is difficult to calculate the cumulative intensity of stars at varying galactic locations."
+ This calculation is vital in eliminating one of the main foregrounds to the CIRB between 1 aaud 5un., This calculation is vital in eliminating one of the main foregrounds to the CIRB between 1 and 5.
+. To get an estimate of the inteusity. (rom stars. theoretical inodels of the Galaxy have been coustructec from which number counts in any direction can be derived.," To get an estimate of the intensity from stars, theoretical models of the Galaxy have been constructed from which number counts in any direction can be derived."
+ The most recent aud successful 11O¢el of the Galaxy was an euhlancerent of the Waluscoatefal.(1992) model macle yy Cohen(1991) wlo divided the Galaxy inο 2 maln components: the disk. the spiral aruis. the molecual jug. the central bulge. aud the exeuded halo. as well as several minor Components: Could’s bel i.|ocal inolecular clouds. rellectiou ebulae. etc.," The most recent and successful model of the Galaxy was an enhancement of the \citet{WCVWS92} model made by \citet{Co94} who divided the Galaxy into 5 main components: the disk, the spiral arms, the molecular ring, the central bulge, and the extended halo, as well as several minor components: Gould's belt, local molecular clouds, reflection nebulae, etc."
+ Iu his model. each major component has 82 diffe“ell stellar vpes distriuted among tlle components with weighting techniques designed to reprocluc οςobservatlous of eacl1 COnponent in the NIR.," In his model, each major component has 82 different stellar types distributed among the components with weighting techniques designed to reproduce observations of each component in the NIR."
+ Aclitionally the spectral classes are characterized wy absolie nagnitucles at IR wavelengths. a magnitude dispersion. a scale height above |ie galactic plare. ancl a soar neighorlioocd cenusity.," Additionally the spectral classes are characterized by absolute magnitudes at IR wavelengths, a magnitude dispersion, a scale height above the galactic plane, and a solar neighborhood density."
+ Finally. dust in the Calaxy is modeled usine the Rieke&Lebolsky(1985) 1uillornm ext]uctiou law.," Finally, dust in the Galaxy is modeled using the \citet{RL85} uniform extinction law."
+ For the purpose of «etectiig the CIRB. the 11odel has to be accurate at high ealactic latitudes where tle πο count5 are lowest axl present the least amourt of contamination.," For the purpose of detecting the CIRB, the model has to be accurate at high galactic latitudes where the number counts are lowest and present the least amount of contamination."
+ This is where the cliffieulies with the 11oclel )ecome evident., This is where the difficulties with the model become evident.
+ 51ice the sun is siuated about 15 pe above the clisk aud we are looking throwel the disk:and tje halo when we are ooking at high οalactic latittcles. aby error 1 the halo-to«lisk weightine wil cenefate solue error in the star coLuts.," Since the sun is situated about 15 pc above the disk and we are looking through the disk the halo when we are looking at high galactic latitudes, any error in the halo-to-disk weighting will generate some error in the star counts."
+ An additional source of e‘ror will be the acttal distance «cof the suu a)»ve the cisk., An additional source of error will be the actual distance of the sun above the disk.
+ The greater the height aOVE the clisk. 1je lower the contribution of disς stars whe looking North vs. South.," The greater the height above the disk, the lower the contribution of disk stars when looking North vs. South."
+ These probelus alfect the οςαἱ» aπι he surface briglitness caletlations αἱ alJOU the leve (Cohen1996)., These problems affect the counts and the surface brightness calculations at about the level \citet{Co96}.
+. A lilleent so'e of error arises whet the inheretly statisical precictious of the mode are subtracted [rom oved data., A different source of error arises when the inherently statistical predictions of the model are subtracted from observed data.
+ Any model will not exactly recreate the nuijiber counts iu the ealactic region of ist., Any model will not exactly recreate the number counts in the galactic region of interest.
+ There will be a variatce [ro1 tlie actial number of stals., There will be a variance from the actual number of stars.
+ This variance is of little COLlseqence for din stars since eacli star contributes relatively little to the overall luteusitv. buU the val]ance becomes very imporant for the uxighter stars wrere the flux of each judividual slar COLributes sieuificautly to the overall inteusitN.," This variance is of little consequence for dim stars since each star contributes relatively little to the overall intensity, but the variance becomes very important for the brighter stars where the flux of each individual star contributes significantly to the overall intensity."
+" While he zocliacal light moclel is intriusicaly on the saije [lux scale as he DIRBE data. the 'elative callbration ol the exterual stellar data for wieght stars aud models for faint. sta""s the DIRBE cata neecs to be known with good accu""ACY."," While the zodiacal light model is intrinsically on the same flux scale as the DIRBE data, the relative calibration of the external stellar data for bright stars and models for faint stars the DIRBE data needs to be known with good accuracy."
+ Thefaitt star moclel of Aretdiefad.(1998) sec IlLNON €i zero uaguitude of Ρον)=612.3J vald FL)=285Jy (COBESupplement LOS)., Thefaint star model of \citet{AOWSH98} used fluxes at zero magnitude of $F_\circ(K) = 612.3\;\mbox{Jy}$ and $F_\circ(L) = 285\;\mbox{Jy}$ \citep{DIRBE-ExpSup}.
+ We have checked these vales by extracting the total flux in 16 pixel areasALOULL the bright stars 3 Aud. a Tau. a Aur. a )rl. a. Boo. a Her 3 Peg.," We have checked these values by extracting the total flux in 16 pixel areasaround the bright stars $\beta$ And, $\alpha$ Tau, $\alpha$ Aur, $\alpha$ Ori, $\alpha$ Boo, $\alpha$ Her $\beta$ Peg."
+ Tle median fF:(A) clerivec| from these stars is fly)611Jy with a mean and standard cleviation of he mean 616.6+12.5 Jv., The median $F_\circ(K)$ derived from these stars is $F_\circ(K) = 614\;\mbox{Jy}$ with a mean and standard deviation of the mean $616.6 \pm 12.5\;\mbox{Jy}$ .
+ This calibration agrees with Areuclefal.(1998) to better thar Ut.., This calibration agrees with \citet{AOWSH98} to better than .
+ But at 3.5, But at 3.5
+so strongly over predicted. that it appears unlikely that any secular processes in the minority barvonic component could redistribute sullicientIy large amounts of the clark halo profile to match the observation.,so strongly over predicted that it appears unlikely that any secular processes in the minority baryonic component could redistribute sufficiently large amounts of the dark halo profile to match the observation.
+" We fix Vay,=47.5hm/s from the observed. rotation curve. A2, from eq.(11). and Aly—3.22510M. from the observed light and gas profiles. and use £4,=4 and fF=0.06. as calibrated from the LSB cases. to fully constrain this model."," We fix $V_{M}=47.5 km/s$ from the observed rotation curve, $R_d$ from eq.(11) and $M_{d}=3.2 \times 10^{8} M_{\odot}$ from the observed light and gas profiles, and use $P_0=4$ and $F=0.06$, as calibrated from the LSB cases, to fully constrain this model."
+ Phe barvonic component in this galaxy is dominated by an extended gas cisk which is not exponential. but given this svstem is totally dark matter dominated. this does not influence the mocelec rotation curve significanth. although it does make estimating the initial value of A for this svstem rather mocdel-dependent.," The baryonic component in this galaxy is dominated by an extended gas disk which is not exponential, but given this system is totally dark matter dominated, this does not influence the modeled rotation curve significantly, although it does make estimating the initial value of $\lambda$ for this system rather model-dependent."
+ The resulting final rotation curve is shown in Figure (9). ancl is seen to reproduce the data reasonably well at all racii although the outer decreasing region appears further out in the model. (," The resulting final rotation curve is shown in Figure (9), and is seen to reproduce the data reasonably well at all radii, although the outer decreasing region appears further out in the model. ("
+Given that the observational relations used. to obtain f° and £% contain some spread. we could. fine-tune the numbers to improve the agreement.,"Given that the observational relations used to obtain $F$ and $P_0$ contain some spread, we could fine-tune the numbers to improve the agreement."
+ Further. there are observational uncertainties in the total gas and star content and on the distance to this galaxy. all of which could. be adjusted to present an optimized fit.," Further, there are observational uncertainties in the total gas and star content and on the distance to this galaxy, all of which could be adjusted to present an optimized fit."
+ We do not do this as 10 spirit of our approach is its universality.), We do not do this as the spirit of our approach is its universality.)
+ The mocoeled curve is thus not a fit. it is merely our standard LSB galaxy with the same Vy; as DDO 154.," The modeled curve is thus not a fit, it is merely our standard LSB galaxy with the same $V_M$ as DDO 154."
+ We note that although only representative and central observational constraints have been used in this analysis. the agreement with the whole rotation curve is quite good.," We note that although only representative and central observational constraints have been used in this analysis, the agreement with the whole rotation curve is quite good."
+ This. together with the morphological similarity of LSB and dwarf galaxies (de Blok et al 1996) suggests that a halo profile of the form described here is universally. applicable.," This, together with the morphological similarity of LSB and dwarf galaxies (de Blok et al 1996) suggests that a halo profile of the form described here is universally applicable."
+For the svstems where the Na Ll doublet has been used in the reconstructions. we have emploved this technique of adding a positive constant to cach data value.,"For the systems where the Na I doublet has been used in the reconstructions, we have employed this technique of adding a positive constant to each data value."
+ We find that this significantly improves the reliability of the technique to determine the component masses when noisy data are being used., We find that this significantly improves the reliability of the technique to determine the component masses when noisy data are being used.
+ As the cata for It er are from Gaussian fits to the He LE A4686AÀX emission line. there are no errors in the continuum and hence— we do not apply this technique.," As the data for HU Aqr are from Gaussian fits to the He II $\lambda$ emission line, there are no errors in the continuum and hence we do not apply this technique."
+" In ασΠλοία, we have accounted for the elfects of velocity smearing due to the motion of the secondary star over the duration of an exposure for all four svstems."," In addition, we have accounted for the effects of velocity smearing due to the motion of the secondary star over the duration of an exposure for all four systems."
+ This is done by calculating the profile at evenly. spaced. intervals over one exposure and then averaging them., This is done by calculating the profile at evenly spaced intervals over one exposure and then averaging them.
+ Although computationally more time-consuming. this reduces the ellects of smearing. and can also take into account. features that may cross the limb of the secondary. star during an exposure.," Although computationally more time-consuming, this reduces the effects of smearing, and can also take into account features that may cross the limb of the secondary star during an exposure."
+ Roche tomograms are greatly influencecl by parameters such as the component masses. inclination. ancl svstemic velocity. of the binary.," Roche tomograms are greatly influenced by parameters such as the component masses, inclination and systemic velocity of the binary."
+ Errors in the assumed: values of these parameters degraces the quality of the surface map and. in general. results in additional structure in. the Roche tomograms.," Errors in the assumed values of these parameters degrades the quality of the surface map and, in general, results in additional structure in the Roche tomograms."
+ The correct. parameters are. therefore. those that produce the map containing the least artefacts. corresponding to the map of highest entropy (Paper D).," The correct parameters are, therefore, those that produce the map containing the least artefacts, corresponding to the map of highest entropy (Paper I)."
+ By carrving out reconstructions for many pairs of component masses (iterating to the sanie XE on each occasion) we can construct an ‘entropy landscape’ (e.g... Eο. 3)).," By carrying out reconstructions for many pairs of component masses (iterating to the same $\chi^2$ on each occasion) we can construct an `entropy landscape' (e.g., Fig. \ref{fig:landscapes}) ),"
+ where each point corresponds to the entropy value obtained in a reconstruction for a particular pair of component masses., where each point corresponds to the entropy value obtained in a reconstruction for a particular pair of component masses.
+ In order to determine the binary parameters of the CVs studied. in. this. paper. we constructed à. series. of entropy landscapes assuming a wide range of dilferent values for the orbital inclination and systemic velocity.," In order to determine the binary parameters of the CVs studied in this paper, we constructed a series of entropy landscapes assuming a wide range of different values for the orbital inclination and systemic velocity."
+ For each combination of ὁ and 5. we picked the maximum entropy value in the corresponding entropy landscape. and. plotted this value in the //—5 plane. which is shown for each object on the lelt-hanel side of Fig. 1.," For each combination of $i$ and $\gamma$, we picked the maximum entropy value in the corresponding entropy landscape, and plotted this value in the $i-\gamma$ plane, which is shown for each object on the left-hand side of Fig. \ref{fig:incl}."
+ In all cases we found that a unique or optimal svstemic velocity. could. be selected that consistently eave the map of maximum. entropy., In all cases we found that a unique or optimal systemic velocity could be selected that consistently gave the map of maximum entropy.
+ 1n addition. the value of the optimal svstemic velocity was found to be largely independent. of the orbital inclination assumed curing the reconstructions. as demonstrated in the left-hand. plots of Fig. 1..," In addition, the value of the optimal systemic velocity was found to be largely independent of the orbital inclination assumed during the reconstructions, as demonstrated in the left-hand plots of Fig. \ref{fig:incl}."
+ For example. in reconstructions carried: out over a range of inclinations spanning 307. the svstemic velocity we obtain for both IP Peg and AM Ler only varies bv 1 km +.," For example, in reconstructions carried out over a range of inclinations spanning $^{\circ}$, the systemic velocity we obtain for both IP Peg and AM Her only varies by 1 km $^{-1}$."
+ The inclination. however. is not as well constrained [or any of the CVs.," The inclination, however, is not as well constrained for any of the CVs."
+ The right-hand plots on Fig., The right-hand plots on Fig.
+ 1. show cuts through the plots in the ¢5. plane at the optimal svstemic velocity. and the scatter clearly demonstrates the difficulty) in selecting an optimal inclination. although the deterioration in the quality of the reconstruction for some inclinations allow limits to be set.," \ref{fig:incl} show cuts through the plots in the $i-\gamma$ plane at the optimal systemic velocity, and the scatter clearly demonstrates the difficulty in selecting an optimal inclination, although the deterioration in the quality of the reconstruction for some inclinations allow limits to be set."
+ The dillieultv in constraining the inclination can be explained by considering that there are two ways of etermining the inclination of non-eclipsing CVs from the μαvectra of the secondary star alone., The difficulty in constraining the inclination can be explained by considering that there are two ways of determining the inclination of non-eclipsing CVs from the spectra of the secondary star alone.
+ First. from the variation in the projected radius of the Itoche-Iobe shaped secondary star às it is viewed at clillerent aspects. corresponding to a variation in the shape (or measured. esin;) of the line profiles (see. 2).," First, from the variation in the projected radius of the Roche-lobe shaped secondary star as it is viewed at different aspects, corresponding to a variation in the shape (or measured $v\sin i$ ) of the line profiles (see \citealt{shahbaz98}) )."
+ The second. is from the variation in the μαrength of the profiles. since a high inclination svstemi will exhibit a much greater variation in the line profile strength 1an a low inclination svstem with the same surface intensity istribution.," The second is from the variation in the strength of the profiles, since a high inclination system will exhibit a much greater variation in the line profile strength than a low inclination system with the same surface intensity distribution."
+ Unfortunately. the variation in esin Is expected to be. ab most. around 2025 kms + for high inclination/ svstenis.," Unfortunately, the variation in $v\sin i$ is expected to be, at most, around 20–25 km $^{-1}$ for high inclination systems."
+ Although this is comparable to the velocity. resolution of our data. it is buried in noise and hence the mapping," Although this is comparable to the velocity resolution of our data, it is buried in noise and hence the mapping"
+As will become clearer below. it is important to know how lar to trust the derived peculiar velocities.,"As will become clearer below, it is important to know how far to trust the derived peculiar velocities."
+ They are uncertain because of the uncertainties in (hie model parameters. and because of uneertainties in distance (uncertainties in angle having negligible effect).," They are uncertain because of the uncertainties in the model parameters, and because of uncertainties in distance (uncertainties in angle having negligible effect)."
+ For situations in which a derived quantity Q is a function of various independent. parameters p;. each with an uncertainty op;. I have used the formula which is sufliciently accurate as long as the errors are not large fractions of the values.," For situations in which a derived quantity $Q$ is a function of various independent parameters $p_i$, each with an uncertainty $\delta p_i$, I have used the formula which is sufficiently accurate as long as the errors are not large fractions of the values."
+ For the uncertainties in peculiar velocities due to model parameter uncertainties and clistance errors (0.0) this technique gives And for the uncertainties in absolute magnitude M due to errors in apparent magnitude m and distance D. All B photometry is assigned a om of 0.4 magnitude. while A' uncertainties are taken from the source.," For the uncertainties in peculiar velocities due to model parameter uncertainties and distance errors $\delta D$ ) this technique gives And for the uncertainties in absolute magnitude $M$ due to errors in apparent magnitude $m$ and distance $D$, All $B$ photometry is assigned a $\delta m$ of 0.4 magnitude, while $K$ uncertainties are taken from the source."
+ The results are listed in Table (4))., The results are listed in Table \ref{data3}) ).
+ The uncertainties in absolute magnitude for those galaxies without photometry reflects errors in distances alone. and are left in the table to give an idea of (he relative importance of distance and photometry uncertainties.," The uncertainties in absolute magnitude for those galaxies without photometry reflects errors in distances alone, and are left in the table to give an idea of the relative importance of distance and photometry uncertainties."
+ We now turn to smaller-scale motions. those which can be connected with individual ealaxies or galaxy groups.," We now turn to smaller-scale motions, those which can be connected with individual galaxies or galaxy groups."
+ For (his the dvnamical vouth of the Local Volume is invoked: given the peculiar velocity dispertion of 70-80 kins 1. we expect that satellite galaxies around a nass concentration to be mostly still infalling.," For this the dynamical youth of the Local Volume is invoked: given the peculiar velocity dispertion of 70-80 km $^{-1}$, we expect that satellite galaxies around a mass concentration to be mostly still infalling."
+ To look for this. we pick out the brightest galaxies. (hose wilh Mj of -20 or brighter. (," To look for this, we pick out the brightest galaxies, those with $M_B$ of -20 or brighter. ("
+While one can argue about (he role of starbursts. for example. in skewing (he mass-to-light ratio. ib is certain that Mg of -20does select the very biggest concentrations of luminous natter in the Local Volume.,"While one can argue about the role of starbursts, for example, in skewing the mass-to-light ratio, it is certain that $M_B$ of -20 select the very biggest concentrations of luminous matter in the Local Volume."
+ They are also the brightest in A.), They are also the brightest in $K$ .)
+ We turn our attention to (hose galaxies between 0.5 and 1.5 Mpc away trom the bright ones. to isolate a sample which should feel (heir eravity but which probably havent crossecl from one side to the other (near," We turn our attention to those galaxies between 0.5 and 1.5 Mpc away from the bright ones, to isolate a sample which should feel their gravity but which probably haven't crossed from one side to the other (near"
+ ~0.022... Z.. ~LOOAL. Zeit~10! :215 Z4; ," $\sim 0.02
+Z_\odot$ $Z_\odot$ $\sim 10^3 \msun$ $Z_{\rm crit} \sim 10^{-4}$ $z \gsim 15$ $Z_{\rm crit}$ "
+at the telescope site.,at the telescope site.
+ In addition. Sucll’s law of refraction states that product sin: between the heielt-dependent iudex of refraction and the sine of the anele of refraction at that heigl: remains constant alone the path(12)).," In addition, Snell's law of refraction states that product $n\sin z$ between the height-dependent index of refraction $n$ and the sine of the angle of refraction at that height remains constant along the path."
+. Therefore the optical path delay is D=μα:bugsin cy: it changes as the astronomical objec changes position iu :. or. supposed Pb is fixed. according to atmospheric parameters accessible at the eround level.," Therefore the optical path delay is $D=b\sin z=bn_0\sin z_0$ ; it changes as the astronomical object changes position in $z$, or, supposed $b$ is fixed, according to atmospheric parameters accessible at the ground level."
+ The benefit of this analysis is that both. the pointing correction R in ((1)) and the OPD mieasured on the erouud. are functions of the index of refraction at the clescope site. not functionals of the entire layered atiuospheroe.," The benefit of this analysis is that both, the pointing correction $R$ in \ref{eq.Rofnflat}) ) and the OPD measured on the ground, are functions of the index of refraction at the telescope site, not functionals of the entire layered atmosphere."
+ The theme of this paper are corrections to hese statements considering a non-turbuleut atinosphiere covering an earth surface of constant. ut non-neelieible curvature.," The theme of this paper are corrections to these statements considering a non-turbulent atmosphere covering an earth surface of constant, but non-negligible curvature."
+ The rest of refSec.lateral— shortlv describes the staucard heorv of refraction and defines two ciffereut vascline Ieneths., The rest of \\ref{Sec.lateral} shortly describes the standard theory of refraction and defines two different baseline lengths.
+ rofsec.zeni concentrates on the integral formulation of the OPD calculation throueh the atmosphere: econietries with coustrained azimuths suffice to introduce all relevant concepts: general star positions are then reduced to the constrained case., \\ref{sec.zeni} concentrates on the integral formulation of the OPD calculation through the atmosphere: geometries with constrained azimuths suffice to introduce all relevant concepts; general star positions are then reduced to the constrained case.
+ Things are more complicated if we start to look at the more realistic model of a spherical earth., Things are more complicated if we start to look at the more realistic model of a spherical earth.
+ A telescope distance of b=100 im on an earth of radius p—6365 kin leads to a pointing mismatch of purely geometric origin of about b/pz3.2 arcsec rofPoiutRho.ps))., A telescope distance of $b=100$ m on an earth of radius $\rho=6368$ km leads to a pointing mismatch of purely geometric origin of about $b/\rho\approx 3.2$ arcsec \\ref{PointRho.ps}) ).
+ The baseline b is the distance between the telescope locations ou the earth. the leneth of the straight secant drawn in and 3.. The inversion of this series reads," The baseline $b$ is the distance between the telescope locations on the earth, the length of the straight secant drawn in \\ref{PointRho.ps} and \ref{Base.ps}, , The inversion of this series reads"
+"survey limit (however. the precise number of bins below the limit is difficult to determine because the redshift bin width is not presented),","survey limit (however, the precise number of bins below the limit is difficult to determine because the redshift bin width is not presented)."
+ At very low luminosities. particularly at luminosities where the survey limit starts cutting into the number of objects in a bin. the number densities will start decreasing because the incomplete bins. due to the shape of the flux cutoff. will preferentially sample a different part of the LF than would be the case for a complete bin.," At very low luminosities, particularly at luminosities where the survey limit starts cutting into the number of objects in a bin, the number densities will start decreasing because the incomplete bins, due to the shape of the flux cutoff, will preferentially sample a different part of the LF than would be the case for a complete bin."
+ We will show in Section ??. that the position of the dominant break will depend on the minimum redshift of the bin in the case of a non-evolving luminosity function., We will show in Section \ref{sec:Method} that the position of the dominant break will depend on the minimum redshift of the bin in the case of a non-evolving luminosity function.
+ A final example ts that of the binned luminosity function of the quarter-Jansky sample from Walletal.(2005)., A final example is that of the binned luminosity function of the quarter-Jansky sample from \citet{Wall05}.
+. The flattening of the LF of this sample at low luminosities produces the illusion of the luminosity evolution. which was recognized by the authors as such.," The flattening of the LF of this sample at low luminosities produces the illusion of the luminosity evolution, which was recognized by the authors as such."
+ The authors attributed the flattening of the LF at lower luminosities to an intrinsic spread in radio spectral indices., The authors attributed the flattening of the LF at lower luminosities to an intrinsic spread in radio spectral indices.
+ Since this flattening (or steepening in the case of negative evolution) is caused by undersampling of the bins cut by the survey limit. one solution to eliminate these spurious breaks is to stop binning as soon as the lowest luminosity bin reaches the survey limit.," Since this flattening (or steepening in the case of negative evolution) is caused by undersampling of the bins cut by the survey limit, one solution to eliminate these spurious breaks is to stop binning as soon as the lowest luminosity bin reaches the survey limit."
+ This. however. can lead to the exclusion of data points lying between this last complete bin and the survey limit.," This, however, can lead to the exclusion of data points lying between this last complete bin and the survey limit."
+ LaFranca&Cristiani(1997);Miyajietal.(2001) proposed to correct for biases in the undersampled bins by prorating them for the missing parts using a model LF determined by independent means.," \citet{LaFranca97,Miyaji01} proposed to correct for biases in the undersampled bins by prorating them for the missing parts using a model LF determined by independent means."
+ This approach. however. has the disadvantage of depending on an assumed model.," This approach, however, has the disadvantage of depending on an assumed model."
+" Efstathiouetal.(1988) have developed a non-parametric “step-wise maximum likelihood"" method that allows a binned LF to be fit directly to the unbinned data.", \citet{Efstathiou88} have developed a non-parametric “step-wise maximum likelihood” method that allows a binned LF to be fit directly to the unbinned data.
+ While these methods their own advantages and disadvantages. we have chosen to focus in this paper on the method of Page Carrera.," While these methods their own advantages and disadvantages, we have chosen to focus in this paper on the method of Page Carrera."
+ We note that any method that computes the values of the LF over incomplete bins must. in one way or another. use some sort of extrapolation to the parts of the bins lying below the flux limit.," We note that any method that computes the values of the LF over incomplete bins must, in one way or another, use some sort of extrapolation to the parts of the bins lying below the flux limit."
+ In this paper we show that the biases associated with undersampled bins can produce artificial breaks in the derived lummosity functions. particularly in the case of strong evolution.," In this paper we show that the biases associated with undersampled bins can produce artificial breaks in the derived luminosity functions, particularly in the case of strong evolution."
+ As evolution of the LF is typically judged by the shifts of features in the LF. these artificial breaks. when interpreted as real. can produce the illusion. of luminosity evolution.," As evolution of the LF is typically judged by the shifts of features in the LF, these artificial breaks, when interpreted as real, can produce the illusion of luminosity evolution."
+ We show that the flattening (or breaking) of the binned LF for bins close to the flux cutoff of the survey can be minimized if the binning i5 performed in the flux-redshift plane instead of the luminosity-redshift plane as it 1s usually done., We show that the flattening (or breaking) of the binned LF for bins close to the flux cutoff of the survey can be minimized if the binning is performed in the flux–redshift plane instead of the luminosity–redshift plane as it is usually done.
+ The improvement is most apparent at lower redshifts where. due to the steepness of the luminosity cutoff in the logL—z space. many bins are undersampled. and therefore give strongly biased estimates of the luminosity function.," The improvement is most apparent at lower redshifts where, due to the steepness of the luminosity cutoff in the $\log L-z$ space, many bins are undersampled, and therefore give strongly biased estimates of the luminosity function."
+" While our analysis of the distortion of the shape of the binned LFs is performed using the method of Page&Car-rera(2000). for constructing binned luminosity functions. this analysis equally applies to the 1/V, method as well."," While our analysis of the distortion of the shape of the binned LFs is performed using the method of \citet{PC00} for constructing binned luminosity functions, this analysis equally applies to the $1/V_{\mathrm{a}}$ method as well."
+" The binned LFs constructed using the 1/V, method will exhibit larger deviations from the true LF than the method of Page&Carrera(2000) at extremely low luminosities GL<μα. see equation 9))."," The binned LFs constructed using the $1/V_{\mathrm{a}}$ method will exhibit larger deviations from the true LF than the method of \citet{PC00} at extremely low luminosities $L<L_{\mathrm{break,3}}$, see equation \ref{eqn:PosBreak3}) )."
+" However. in the case of non-evolving LFs. the 1/V, method. unlike the method of Page Carrera. will not distort the shape of the LF for Lois<b<Loa because it assigns a variable weight to each object in the bin and can correctly handle the flux limit down to lower lummosities ρα)."," However, in the case of non-evolving LFs, the $1/V_{\mathrm{a}}$ method, unlike the method of Page Carrera, will not distort the shape of the LF for $L_{\mathrm{break,3}}<L<L_{\mathrm{break,1}}$ because it assigns a variable weight to each object in the bin and can correctly handle the flux limit down to lower luminosities $L_{\mathrm{break,3}}$ )."
+ Regardless of the method used. any determination of the LF below Ly3 is pure extrapolation and should be avoided.," Regardless of the method used, any determination of the LF below $L_{\mathrm{break,3}}$ is pure extrapolation and should be avoided."
+ The outline of the paper is as follows: In ?? we describe both binning methods and discuss the main advantages of the binning in the flux-redshift plane over the traditional luminosity-redshift plane binning method., The outline of the paper is as follows: In \ref{sec:Method} we describe both binning methods and discuss the main advantages of the binning in the flux–redshift plane over the traditional luminosity–redshift plane binning method.
+ In ?? we apply these two binning methods to several model luminosity functions and show that the LF binned in the flux-redshift plane is a better approximation to the model LF., In \ref{sec:MCsim} we apply these two binning methods to several model luminosity functions and show that the LF binned in the flux–redshift plane is a better approximation to the model LF.
+ We construct the binned LF of the MOJAVE sample using both methods in 22.., We construct the binned LF of the MOJAVE sample using both methods in \ref{sec:MOJAVELF}. .
+ We summarize our findings and present our conclusions in 22.., We summarize our findings and present our conclusions in \ref{sec:conclusions}.
+" Throughout this paper we assume a cosmology with 0,20.3. 0420.7. ο20 and Hy=70kms7!Mpe7!."," Throughout this paper we assume a cosmology with $\Omega_m=0.3$, $\Omega_{\Lambda}=0.7$ , $\Omega_{r}=0$ and $H_0=70\,\mbox{km}\,\mbox{s}^{-1}\,\mbox{Mpc}^{-1}$ ."
+ All luminosities are quoted as monochromatic luminosities at specific frequency 7., All luminosities are quoted as monochromatic luminosities at specific frequency $\nu$.
+" We also adopt the following convention for the spectral index. oj: S,xir and assume that ©pag20 for all model LFs and plots in this paper."," We also adopt the following convention for the spectral index, $\alpharad$: $S_{\nu}\propto\nu^{\alpharad}$ and assume that $\alpharad=0$ for all model LFs and plots in this paper."
+ The differential luminosity function of a sample of astrophysical objects is defined as the number of objects per unit comoving volume per unit luminosity interval: but. because of the typically large span of the luminosities. it is often defined in terms of logL: and the relationship between the two forms is given by: One of the advantages of working with the LF defined in terms of logL is that for the most common forms of the LF used in model fitting. 1.e.. the simple power law the plot of logo vs. logL will show a simple linear dependence: with a slope (ΕΤ).," The differential luminosity function of a sample of astrophysical objects is defined as the number of objects per unit comoving volume per unit luminosity interval: but, because of the typically large span of the luminosities, it is often defined in terms of $\log L$: and the relationship between the two forms is given by: One of the advantages of working with the LF defined in terms of $\log L$ is that for the most common forms of the LF used in model fitting, i.e., the simple power law the plot of $\log \phi$ vs. $\log L$ will show a simple linear dependence: with a slope $\left( \alpha + 1 \right)$."
+ In equations (4)) and (5)) L. is an arbitrary positive constant., In equations \ref{eqn:singlepow}) ) and \ref{eqn:singlepowloglog}) ) $L_*$ is an arbitrary positive constant.
+" The usual way to construct the binned LF of a sample of objects is to divide the logL—z plane into luminosity-redshift bins and to estimate the space density in a bin using. for example. the 1/V,4, method or other methods."," The usual way to construct the binned LF of a sample of objects is to divide the $\log L-z$ plane into luminosity–redshift bins and to estimate the space density in a bin using, for example, the $1/V_{\max}$ method or other methods."
+" In this paper we use the improved method of Page&Carrera(2000).. that is. we assume that the LF at the center of a bin with a luminosity interval logLj;, and logZ,4, and a redshift interval zii; and μαςcan be approximatedas: whereN is the number of sources detected within the bin limits and Lijs(z) ts the minimum luminosity within the bin at which we can still detect an object."," In this paper we use the improved method of \citet{PC00}, that is, we assume that the LF at the center of a bin with a luminosity interval $\log L_{\min}$ and $\log L_{\max}$ and a redshift interval $z_{\min}$ and $z_{\max}$can be approximatedas: where$N$ is the number of sources detected within the bin limits and $L_{\min}(z)$ is the minimum luminosity within the bin at which we can still detect an object."
+ In Figure 1. we, In Figure \ref{fig:LzLsBinning} we
+"These simulations cannot resolve star-forming events on parsec scales, so our canonical runs employed an artificial minimum pressure to prevent Jeans fragmentation on smaller scales than we could resolve.","These simulations cannot resolve star-forming events on parsec scales, so our canonical runs employed an artificial minimum pressure to prevent Jeans fragmentation on smaller scales than we could resolve."
+" By including this minimum pressure described in ?,, we assured that the Jeans length is always refined by at least 8 cells, and therefore the Truelove criteria was met (seealso???).."," By including this minimum pressure described in \citet{Machacek:2001p1047}, we assured that the Jeans length is always refined by at least 8 cells, and therefore the Truelove criteria was met \citep[see also][]{Truelove:1997p1046, Robertson:2008p1017, Ceverino:2010p1012}."
+" It has been suggested that artificial disk fragmentation leads to large gas clumps that lose angular momentum via dynamical friction, and hence result in angular momentum loss."," It has been suggested that artificial disk fragmentation leads to large gas clumps that lose angular momentum via dynamical friction, and hence result in angular momentum loss."
+" For modified run H26S, we turned off this minimum pressure support."," For modified run H26S, we turned off this minimum pressure support."
+ The results of that run are presented in Figure 10 and 11.., The results of that run are presented in Figure \ref{fig:modified_runs1_SFR} and \ref{fig:modified_runs1_rotation_curves}.
+" The results are nearly identical for both measures, indicating that the minimum pressure support did not have a significant effect on the star formation rate or the distribution of gas and stars in our simulated system."," The results are nearly identical for both measures, indicating that the minimum pressure support did not have a significant effect on the star formation rate or the distribution of gas and stars in our simulated system."
+" The star formation model we adopt in this work is based on the Kennicutt-Schmidt relation, but has an efficiency parameter that is not well-constrained."," The star formation model we adopt in this work is based on the Kennicutt-Schmidt relation, but has an efficiency parameter that is not well-constrained."
+" It has recently been suggested by ? that star formation efficiency is a key parameter controlling the distribution of stars in the disk, and hence the rotation curve."," It has recently been suggested by \citet{Agertz:2010p461} that star formation efficiency is a key parameter controlling the distribution of stars in the disk, and hence the rotation curve."
+" Our efficiency is already fairly low; however in order to examine this suggestion in more detail, and also to probe the impact of decreasing the efficiency in general, we adopt ege=107? in run H26SPML."," Our efficiency is already fairly low; however in order to examine this suggestion in more detail, and also to probe the impact of decreasing the efficiency in general, we adopt $\epsilon_{\rm SF} = 10^{-3}$ in run H26SPML."
+" The star formation history, in Figure 10,, shows a significant delay to later time, as would be expected, although the net amount of stars produced is nearly identical."," The star formation history, in Figure \ref{fig:modified_runs1_SFR}, shows a significant delay to later time, as would be expected, although the net amount of stars produced is nearly identical."
+" However, this does not translate into a flatter rotation curve, as can be seen in Figure 11.."," However, this does not translate into a flatter rotation curve, as can be seen in Figure \ref{fig:modified_runs1_rotation_curves}."
+" 'To understand this result a little more, we look at the rotation curve for this run, at early times."," To understand this result a little more, we look at the rotation curve for this run, at early times."
+ Figure 9 shows the rotation curves at z—4.5 of this low star formation efficiency run H26SPML compared against our canonical run H26SPM., Figure \ref{fig:lowSFE_rotation_curves} shows the rotation curves at $z=4.5$ of this low star formation efficiency run H26SPML compared against our canonical run H26SPM.
+" During this epoch, much of the halo has formed, but star formation has not yet converted most of the gas into stars."," During this epoch, much of the halo has formed, but star formation has not yet converted most of the gas into stars."
+" This plot demonstrates that gas dominates the rotation curve at early times, and yet despite this, the gas clumps have already lost their angular momentum and formed a central cusp of compressed gas."," This plot demonstrates that gas dominates the rotation curve at early times, and yet despite this, the gas clumps have already lost their angular momentum and formed a central cusp of compressed gas."
+" In runs with higher star formation rates the canonical run the gas clumps are partially(e.g. converted to stars before H26SPM),accreting so the cusp is primarily composed of stars, but the net result is the same — the clumps lose angular momentum and form a centrally peaked rotation curve."," In runs with higher star formation rates (e.g. the canonical run H26SPM), the gas clumps are partially converted to stars before accreting so the cusp is primarily composed of stars, but the net result is the same – the clumps lose angular momentum and form a centrally peaked rotation curve."
+" Note that it is still possible that star formation efficiency coupled with feedback may be important — in particular these results are consistent with the idea that low star formation rates combined with sufficient feedback to puff up gas-dominated clumps would suppress angular momentum loss (indeed, the runs in ? generally include feedback and always include a stiff equation of state, P~ "," Note that it is still possible that star formation efficiency coupled with feedback may be important -- in particular these results are consistent with the idea that low star formation rates combined with sufficient feedback to puff up gas-dominated clumps would suppress angular momentum loss (indeed, the runs in \citet{Agertz:2010p461} generally include feedback and always include a stiff equation of state, $P \sim \rho^2$ )."
+"However, we see that by itself, a low star formation efficiencyp?). does not change the distribution ofmatter in the simulated galaxies."," However, we see that by itself, a low star formation efficiency does not change the distribution ofmatter in the simulated galaxies."
+" Also shown in Figure 9 are the rotation curves at z=4.5 for the fixed physical resolution run (H26SPMR), as well as the simulation including feedback, cooling suppression, and fixed resolution (H26SPMFCR)."," Also shown in Figure \ref{fig:lowSFE_rotation_curves} are the rotation curves at $z=4.5$ for the fixed physical resolution run (H26SPMR), as well as the simulation including feedback, cooling suppression, and fixed resolution (H26SPMFCR)."
+" As can be seen, the cusp already appears in the fixed resolution run (although not quite as high as in the canonical run), while gas dominates in the H26SPMFCR simulation."," As can be seen, the cusp already appears in the fixed resolution run (although not quite as high as in the canonical run), while gas dominates in the H26SPMFCR simulation."
+" For that run, the additional pressure from feedback/cooling suppression has allowed the infalling clumps to be incorporated in the disk before losing a significant amount of angular momentum."," For that run, the additional pressure from feedback/cooling suppression has allowed the infalling clumps to be incorporated in the disk before losing a significant amount of angular momentum."
+" For the rest of the modified physics runs (not only the cooling suppression simulations differshown), significantly from the canonical run, and they are similar to the H26SPMFCR simulation, but with a somewhat more pronounced cusp."," For the rest of the modified physics runs (not shown), only the cooling suppression simulations differ significantly from the canonical run, and they are similar to the H26SPMFCR simulation, but with a somewhat more pronounced cusp."
+" Our canonical set of simulations used a fixed maximum refinement level, which translates to a best resolution achievable in comoving spatial units."," Our canonical set of simulations used a fixed maximum refinement level, which translates to a best resolution achievable in comoving spatial units."
+ Another common prescription is to fix the highest resolution as a fixed physical length scale (e.g.?).., Another common prescription is to fix the highest resolution as a fixed physical length scale \citep[e.g.][]{Agertz:2010p461}.
+ We carry out a simulation of Halo 26 in which we vary the maximum allowed refinement level so as to keep as close as possible to a physical resolution of 425 pc., We carry out a simulation of Halo 26 in which we vary the maximum allowed refinement level so as to keep as close as possible to a physical resolution of 425 pc.
+" Note that because of the factor-of-two refinement in AMR, this implies discrete resolution changes at various times during the simulation."," Note that because of the factor-of-two refinement in AMR, this implies discrete resolution changes at various times during the simulation."
+ The star formation history for this simulation is shown in Figure 10.., The star formation history for this simulation is shown in Figure \ref{fig:modified_runs1_SFR}.
+" The lower resolution at early times (for example, at z—5, the spatial resolution is four times worse than in the canonical run) results in lower densities and hence a shift of the bulk of the star formation to later times in the lowered star formation efficiency run H26SPML)."," The lower resolution at early times (for example, at $z=5$, the spatial resolution is four times worse than in the canonical run) results in lower densities and hence a shift of the bulk of the star formation to later times (as in the lowered star formation efficiency run H26SPML)."
+" (asThe rotation curve is shown in Figure 11,, and again, there is not a significant change in the distribution of mass in the core."," The rotation curve is shown in Figure \ref{fig:modified_runs1_rotation_curves}, , and again, there is not a significant change in the distribution of mass in the core."
+this first investigation was extended in Soleretal.(2011a.hereafterSOB) to the case of propagating waves.,this first investigation was extended in \citet[hereafter SOB]{solerspatial} to the case of propagating waves.
+ The use of the single-fluid approximation. as in Soleretal.(2009b) and SOB. seems reasonable as the expected values of the collision frequencies in the solar atmosphere are much larger than the observed wave frequencies.," The use of the single-fluid approximation, as in \citet{solerpartialres} and SOB, seems reasonable as the expected values of the collision frequencies in the solar atmosphere are much larger than the observed wave frequencies."
+ However. the single-fluid approximation misses important effects when small length scales and/or high frequencies are involved.," However, the single-fluid approximation misses important effects when small length scales and/or high frequencies are involved."
+ In such cases. the multifluid description is a more suitable approach (see.e.g..Zagarashvilietal.201 la.b).," In such cases, the multifluid description is a more suitable approach \citep[see, e.g.,][]{temury,temuryhelium}."
+. For resonant waves. perturbations develop very small length scales in the vicinity of the resonance position (see.e.g..Tirry&Goossens1996;RudermanWright1999;Vasquez2005:Terradasetal. 2006)... where multifluid effects may play a relevant role.," For resonant waves, perturbations develop very small length scales in the vicinity of the resonance position \citep[see, e.g.,][]{tirry,rudermanwright,vasquez,terradas2006}, where multifluid effects may play a relevant role."
+ Here we perform a general description of propagating resonant Alfvénnie waves in partially ionized plasmas using the multifluid treatment., Here we perform a general description of propagating resonant Alfvénnic waves in partially ionized plasmas using the multifluid treatment.
+ Our work extends the investigation of SOB by considering arbitrary values of the collision frequencies and by performing a more in-depth analysis of the resonant process., Our work extends the investigation of SOB by considering arbitrary values of the collision frequencies and by performing a more in-depth analysis of the resonant process.
+ Section 2. contains a description of the equilibrium configuration and the basic equations., Section \ref{sec:basic} contains a description of the equilibrium configuration and the basic equations.
+ In Section 3 the behavior of wave perturbations around the Alfvénn resonance position are investigated and jump relations for the perturbations are derived., In Section \ref{sec:resonant} the behavior of wave perturbations around the Alfvénn resonance position are investigated and jump relations for the perturbations are derived.
+ Later. in Section 4. we perform an application to resonant kink waves in straight tubes and obtain expressions for the damping lengths due to resonant absorption and ion-neutral collisions.," Later, in Section \ref{sec:tube} we perform an application to resonant kink waves in straight tubes and obtain expressions for the damping lengths due to resonant absorption and ion-neutral collisions."
+ The approximate analytical theory 1s complemented with fully numerical eigenvalue computations., The approximate analytical theory is complemented with fully numerical eigenvalue computations.
+ In Section 5. the implications of our results for the particular case of chromospheric waves ts discussed., In Section \ref{sec:chromos} the implications of our results for the particular case of chromospheric waves is discussed.
+ Finally. we give our main conclusions in Section 6..," Finally, we give our main conclusions in Section \ref{sec:con}."
+ We consider a partially 10nized. plasma composed of ions. electrons. and neutrals.," We consider a partially ionized plasma composed of ions, electrons, and neutrals."
+ We use cylindrical coordinates. namely F. yg. and z for the radial. azimuthal. and longitudinal coordinates.," We use cylindrical coordinates, namely $r$, $\varphi$, and $z$ for the radial, azimuthal, and longitudinal coordinates."
+ The medium is permeated by an equilibrium magnetic field. B. of the form In general. both azimuthal. B-. and longitudinal. B-. components are functions r.," The medium is permeated by an equilibrium magnetic field, $\bf B$, of the form In general, both azimuthal, $B_\varphi$, and longitudinal, $B_z$, components are functions $r$."
+ The equations governing the dynamics of a magnetized multifluid plasma are discussed in classical works as. e.g.. Cowling(1956) and Braginskii (1965).. or more recently in Birketal.(1996) and Pintoetal. (2008).," The equations governing the dynamics of a magnetized multifluid plasma are discussed in classical works as, e.g., \citet{cowling} and \citet{brag}, , or more recently in \citet{birk} and \citet{pinto}."
+. Recent investigations which make extensive use of the multifluid description in the context of MHD waves are Zagarashvilietal.(2011a.b).," Recent investigations which make extensive use of the multifluid description in the context of MHD waves are \citet{temury,temuryhelium}."
+. We refer the reader to these two papers for details about the derivation of the equations., We refer the reader to these two papers for details about the derivation of the equations.
+ In brief. à partially ionized multifluid plasma is governed by the equations of the different species. which contain terms that couple the various fluids by means of collisions.," In brief, a partially ionized multifluid plasma is governed by the equations of the different species, which contain terms that couple the various fluids by means of collisions."
+ Here we study linear perturbations superimposed on the equilibrium state., Here we study linear perturbations superimposed on the equilibrium state.
+ Due to the very small momentum of electrons. we neglect the collisions of electrons with ions and neutrals.," Due to the very small momentum of electrons, we neglect the collisions of electrons with ions and neutrals."
+ In addition. we consider a magnetically dominated plasma and neglect the gas pressure of tons compared to the Lorentz force. and the gas pressure of neutrals compared to the collisional friction with tons.," In addition, we consider a magnetically dominated plasma and neglect the gas pressure of ions compared to the Lorentz force, and the gas pressure of neutrals compared to the collisional friction with ions."
+" This simplification neglects the plasma displacement along the magnetic field direction and removes the slow or cusp continuum from the equilibrium,", This simplification neglects the plasma displacement along the magnetic field direction and removes the slow or cusp continuum from the equilibrium.
+ Thus. in the present work we focus on resonance absorption in the Alfvénn continuum only.," Thus, in the present work we focus on resonance absorption in the Alfvénn continuum only."
+ The inclusion of gas pressure is a subject for future works., The inclusion of gas pressure is a subject for future works.
+" Under these conditions. the basic equations of ourinvestigation are where v; and v, are the velocities of ions and neutrals. respectively. b is the magnetic field perturbation. o; and p, are the densities of the ton and neutral fluids. respectively. µ 1s the magnetic permittivity. and γι is the 10n-neutral collision rate coefficient per unit mass."," Under these conditions, the basic equations of ourinvestigation are where ${\bf v}_{\rm i}$ and ${\bf v}_{\rm n}$ are the velocities of ions and neutrals, respectively, ${\bf b}$ is the magnetic field perturbation, $\rhoi$ and $\rhon$ are the densities of the ion and neutral fluids, respectively, $\mu$ is the magnetic permittivity, and $\gamma_{\rm in}$ is the ion-neutral collision rate coefficient per unit mass."
+ Instead of using yj. in the remaining of this paper we use the ton-neutral collision frequency. vi. which has a more obvious physical meaning.," Instead of using $\gamma_{\rm in}$, in the remaining of this paper we use the ion-neutral collision frequency, $\nu_{\rm in}$, which has a more obvious physical meaning."
+ The ton-neutral collision frequency is defined as Equations (2)) and (3)) are the linearized momentum equations of ions and neutrals. respectively.," The ion-neutral collision frequency is defined as Equations \ref{eq:momion}) ) and \ref{eq:momneu}) ) are the linearized momentum equations of ions and neutrals, respectively."
+ Equation (4)) is the linearized 1duction equation. in which we have omitted all nonideal. diffusion terms as. e.g.. magnetic resistivity.," Equation \ref{eq:induct}) ) is the linearized induction equation, in which we have omitted all nonideal, diffusion terms as, e.g., magnetic resistivity."
+ We do so because the goal of the present paper is to assess the particular role of ton-neutral collisions on the resonance absorption process., We do so because the goal of the present paper is to assess the particular role of ion-neutral collisions on the resonance absorption process.
+ The effects of other dissipative mechanisms as. e.g. resistivity or viscosity. have beenextensively studied in the existing literature (seetherecentreviewbyGoossensetal.2011) In this work the equilibrium quantities are functions of r alone. so that the equilibrium is uniform in both azimuthal and longitudinal directions.," The effects of other dissipative mechanisms as, e.g, resistivity or viscosity, have beenextensively studied in the existing literature \citep[see the recent review by][]{goossens11}
+ In this work the equilibrium quantities are functions of $r$ alone, so that the equilibrium is uniform in both azimuthal and longitudinal directions."
+ Hence we can write the perturbed quantities proportional to exp(img+Ik-z). where a and κ: are the azimuthal and longitudinal wavenumbers. respectively.," Hence we can write the perturbed quantities proportional to $\exp \left( i m \varphi + i k_z z \right)$, where $m$ and $k_z$ are the azimuthal and longitudinal wavenumbers, respectively."
+ We express the temporal dependence of the perturbations as exp(-fwt). with w the frequency.," We express the temporal dependence of the perturbations as $\exp \left( - i \omega t \right)$, with $\omega$ the frequency."
+ Since there are no equilibrium flows. the Lagrangian displacement of ions. £. and neutrals. £j. are From Equation (3)) it is straightforward to derive the relation between & and &. namely Equation (7)) informs us that & and & are related to each other.," Since there are no equilibrium flows, the Lagrangian displacement of ions, ${\bf \xi}_{\rm i}$, and neutrals, ${\bf \xi}_{\rm n}$, are From Equation \ref{eq:momneu}) ) it is straightforward to derive the relation between ${\bf \xi}_{\rm n}$ and ${\bf \xi}_{\rm i}$, namely Equation \ref{eq:xiixin}) ) informs us that ${\bf \xi}_{\rm n}$ and ${\bf \xi}_{\rm i}$ are related to each other."
+ The factor of proportionally is a complex function of the wave frequency. ω. and the ion-neutral collision frequency. Vine ," The factor of proportionally is a complex function of the wave frequency, $\omega$ , and the ion-neutral collision frequency, $\nu_{\rm in}$ ."
+Hence. there is a phase difference between the motions of tons and neutrals.," Hence, there is a phase difference between the motions of ions and neutrals."
+ Physically. this phase difference can be interpreted as the “timedelay” between the motions of tons and the corresponding reaction of neutrals.," Physically, this phase difference can be interpreted as the “timedelay” between the motions of ions and the corresponding reaction of neutrals."
+ Thus. for vi; Sw.," Thus, for $\nu_{\rm in} \gg \omega$ ,"
+ where ry is the Schwarzschild radius of the hole. ep is in units of 10kms| and we have assumed specific angular momentum conservation inskle rage.," where $r_S$ is the Schwarzschild radius of the hole, $v_D$ is in units of $10\,\kms$ and we have assumed specific angular momentum conservation inside $r_{\rm acc}$."
+ The accretion disc has therefore a size of order of the trapping radius:unma. ie. (he radius at which radiation is trapped as the infall speed of the gas is larger than the diffusion speed of the radiation.," The accretion disc has therefore a size of order of the trapping radius:, i.e. the radius at which radiation is trapped as the infall speed of the gas is larger than the diffusion speed of the radiation."
+ The trapping radius is equivalent to the radius of spherization defined by Shakura Suuvaev (1973). where the thickness of the disc becomes of (he same order as rj: at smaller radii the (he inflow is quasi-spherical.," The trapping radius is equivalent to the radius of spherization defined by Shakura Sunyaev (1973), where the thickness of the disc becomes of the same order as $r_{\rm tr}$; at smaller radii the the inflow is quasi-spherical."
+ Beeelman (1979) and Beeelnan Meier (1982) studied super-criGical accretion onto a BIL in spherical geometry ancl quiescent thick discs respectively., Begelman (1979) and Begelman Meier (1982) studied super-critical accretion onto a BH in spherical geometry and quiescent thick discs respectively.
+ In. (he spherical case. radiation pressure cannot prevent the accretion rate from being super-crilical. while the enmergent Iuminosity is limited to ZiZia.," In the spherical case, radiation pressure cannot prevent the accretion rate from being super-critical, while the emergent luminosity is limited to $\lta L_{\rm Edd}$."
+ The case with angular momentum is nore complicated (especially with regard to the role of winds)., The case with angular momentum is more complicated (especially with regard to the role of winds).
+ Though this issue remains unclear. it still seems possible that when the inflow rate is super-critical. (he radiative efficiency. drops so that the hole can accept the material without greatly exceeding the Eddington luminosity.," Though this issue remains unclear, it still seems possible that when the inflow rate is super-critical, the radiative efficiency drops so that the hole can accept the material without greatly exceeding the Eddington luminosity."
+ The efficiency could be low either because most radiation is trapped and acdvectecl inward. or because the flow adjusts so that the material can plunge in [rom an orbit with small binding enerev CAbramowiez Lasota 1980).," The efficiency could be low either because most radiation is trapped and advected inward, or because the flow adjusts so that the material can plunge in from an orbit with small binding energy (Abramowicz Lasota 1980)."
+ Despite (he uncertainties it seems worthwhile to explore the consequences Lor black hole erowth of an early phase when the MDII can accept most the mass infalling at a rate that can be estimated using the Dondi-IHovle Formula. and can therefore grow much more rapidly than the Eddington rate would allow.," Despite the uncertainties it seems worthwhile to explore the consequences for black hole growth of an early phase when the MBH can accept most the mass infalling at a rate that can be estimated using the Bondi-Hoyle formula, and can therefore grow much more rapidly than the Eddington rate would allow."
+2001b).,.
+. Both cover a relatively large sky area. reducimg he problem of cosmic variance.," Both cover a relatively large sky area, reducing the problem of cosmic variance."
+ Both are deep enough o allow the detection of several galaxies. thus uiininiziugs the effect of shot noise.," Both are deep enough to allow the detection of several galaxies, thus minimizing the effect of shot noise."
+ Accurate photoiuetric redshifts (At/(tspee11)x0.03 or the FDF and <0.056 for the GOODS-South field) with onlv zz1% catastrophic fülures allow us to nicasure he UV Iuuiunositv fiction down to faiuter linüts than spectroscopic samples.," Accurate photometric redshifts $\Delta z / (z_{spec}+1) \le 0.03$ for the FDF and $\le 0.056$ for the GOODS-South field) with only $\approx
+1$ catastrophic failures allow us to measure the UV luminosity function down to fainter limits than spectroscopic samples."
+ A detailed comparison of the UV mnumositv functions of the FDF with the LF derived in huge surveys was presented iu Cabascletal.(2001) and shows eood agreement in the overlapping magnitude range at all redshifts., A detailed comparison of the UV luminosity functions of the FDF with the LF derived in large surveys was presented in \citet{gabasch:1} and shows good agreement in the overlapping magnitude range at all redshifts.
+ Finally aud iost importantly. the two fields provide us with D-baud. baud and EK-baud selected: catalogues. making it possible to assess the dependence of the SFR on the detection baud aud ealaxy colors aud the associated selection biases.," Finally and most importantly, the two fields provide us with B-band, I-band and K-band selected catalogues, making it possible to assess the dependence of the SFR on the detection band and galaxy colors and the associated selection biases."
+ The Letter is organized as follows., The Letter is organized as follows.
+ In 82. we discuss the photometry aud the photometric redshifts of the two fields. in 83) we present our results on the SER lhüstorv. and in δι woe draw our conclusions.," In \ref{sec_photometry} we discuss the photometry and the photometric redshifts of the two fields, in \ref{sec_sfr} we present our results on the SFR history, and in \ref{sec_conclusions} we draw our conclusions."
+ Throughout the paper we use AB magnitudes aud adopt acoucordance cosmology with0.3...0.7.. and .," Throughout the paper we use AB magnitudes and adopt a concordance cosmology with, and ."
+ The present results are based on photometric catalogues derived for the FDF (Ππααetal.2003:Gabaschetal.2001.UDgRLES21 um..zJIxs aud the GOODS-South fields (Salvatoctal.2001.UBVRI-JIIs bands)...," The present results are based on photometric catalogues derived for the FDF \citep[][ UBgRI,\mbox{834~nm}, ,zJKs and the GOODS-South fields \citep[][ UBVRIJHKs bands]{salvato:1}."
+ The two fields cover approximately the same area (39.81 arcnuün? for ΕΡΕ aud 50 arcmin? for GOODS): the FDF reaches effective absolute magnitude limits z1 inag deeper than COODS-South (see below)., The two fields cover approximately the same area (39.81 $^2$ for FDF and 50 $^2$ for GOODS); the FDF reaches effective absolute magnitude limits $\approx 1$ mag deeper than GOODS-South (see below).
+ We use the Ebaud aud D-baud selected PDF catalogues as derived in Ueidtetal.(2003). aud Cabaschetal. (2001)., We use the I-band and B-band selected FDF catalogues as derived in \citet{fdf_data} and \citet{gabasch:1}.
+".. The B and I selectedcatalogues list 5188 aud 5557 bona-fide galaxies (having excluded the oue known bright quasar in the field) down to δν=27.6 aud lj,=26.8. respectively."," The B and I selectedcatalogues list 5488 and 5557 bona-fide galaxies (having excluded the one known bright quasar in the field) down to $B_{lim}=27.6$ and $I_{lim}=26.8$, respectively."
+ The I)D catalogue obtained by increie these two contains 6756 entries;, The I+B catalogue obtained by merging these two contains 6756 entries.
+ Photometric redshifts for the ΕΤΕ galaxies have au accuracy of with ouly ~L4 catastrophic outlicrs (taper(Cabaschetal.2001)., Photometric redshifts for the FDF galaxies have an accuracy of with only $\sim 1$ catastrophic outliers \citep{gabasch:1}.
+". Qur IKk-haned selected catalogue for the GOODS-South Seld is based on the 8 2.5« arcmin? J. Π. Ks VLT-ISAAC images publically available. taken with seeing in he range OL"" 0.5%."," Our K-band selected catalogue for the GOODS-South field is based on the 8 $2.5\times2.5$ $^2$ J, H, Ks VLT-ISAAC images publically available, taken with seeing in the range $0.4''-0.5''$ ."
+ The U and I images are from GOODS/EIS public survey. while D. V. Ro are. taken your the Carching-Boun Deep Survey.," The U and I images are from GOODS/EIS public survey, while B V R are taken from the Garching-Bonn Deep Survey."
+ Data reduction is described in Arnoutsetal.(2001) aud Sehirineretal (2003).. respectivelv.," Data reduction is described in \citet{arnouts:1} and \citet{schirmer:1}, respectively."
+ The data for the COODS field were analyzed in a very sunilay wav to the data of the FDF., The data for the GOODS field were analyzed in a very similar way to the data of the FDF.
+ The objects were detected in the EK-baud inages closely following the procedure used for the ΕΤΕ I aud D band detection (IIeidtetal.2003).. using both SExtractor (Bertin&Arnouts1996) and the YODA package (Ότοιν2003).," The objects were detected in the K-band images closely following the procedure used for the FDF I and B band detection \citep{fdf_data}, using both SExtractor \citep{bertin} and the YODA package \citep{drory:3}."
+. A detailed description of the procedure can be found in Salvatoetal.(2001)., A detailed description of the procedure can be found in \citet[][]{salvato:1}.
+.. We detected 3367 objects in EK-baud for which we derived magnitudes (fixi aperture and total) im all baud., We detected 3367 objects in K-band for which we derived magnitudes (fixed aperture and total) in all bands.
+ Number couuts match the literature values down to Azx25. [. which is the completeness limit of the catalogue. iu agreement with the nuuber obtained following(2002).," Number counts match the literature values down to $K\approx 25.4$ , which is the completeness limit of the catalogue, in agreement with the number obtained following."
+ Note that much deeper ACS based catalogues are available (Cdavaliscoetal.200[b).. but as we are focusing on the I&-selection they are not rolevaut in this context.," Note that much deeper ACS based catalogues are available \citep{giavalisco:1}, but as we are focusing on the K-selection they are not relevant in this context."
+ We computed photometric redshifts following Benderetal.(2001.2001.inpreparation) aud using the same SED template spectra as for the FDF.," We computed photometric redshifts following \citet[ 2004, in
+ preparation]{bender:1} and using the same SED template spectra as for the FDF."
+ The comparison with the spectroscopic redshifts of the VIMOS team (LeFevreetal.2001). and the FORS2 spectra released at littp://www.cso.org/scicuce(goods/. shows that the photometric redshifts have au accuracy Ati|1)<0.056.," The comparison with the spectroscopic redshifts of the VIMOS team \citep{lefevre:1} and the FORS2 spectra released at http://www.eso.org/science/goods/, shows that the photometric redshifts have an accuracy $\Delta z /(z_{spec}+1) \le 0.056$."
+ Similar results are obtained when comparing(cspee to the COMDO-17 (Wolfetal.2001). data., Similar results are obtained when comparing to the COMBO-17 \citep{wolf:1} data.
+ Tlus is nearly a factor of 2 better than Mobasheretal.(2001) obtained using eround-based plus HST/ACS data., This is nearly a factor of 2 better than \citet{mobasher:1} obtained using ground-based plus HST/ACS data.
+ Stars are ideutified and excluded as in etal.(2001).. as well as known ACN (Szokolvctal.2003).. leaving 3297 bona-fide galaxies used in the further analysis.," Stars are identified and excluded as in \citet{gabasch:1}, as well as known AGN \citep{szokoly:1}, leaving 3297 bona-fide galaxies used in the further analysis."
+ Fig., Fig.
+ l shows the distribution of galaxies (slightly smoothed with a Catussian kernel) in the rest frame 1500A. aabsolute magnitude AM399 vs. redshift plane. computed bv iutegratiug the best fitting SED over the baud definition 100A)).," \ref{fig:sfr:cont_absmag_fdf_goods} shows the distribution of galaxies (slightly smoothed with a Gaussian kernel) in the rest frame $1500 $ absolute magnitude $M_{1500}$ vs. redshift plane, computed by integrating the best fitting SED over the band definition $1500\pm 100 $ )."
+ The contours agree remarkably well at the(15002E bright-end showing that the nuuber density of bright galaxies does not significantly depeud on the waveleugth at which they were selected (for the D-baud. this is of course onlv true up to τz3).," The contours agree remarkably well at the bright-end showing that the number density of bright galaxies does not significantly depend on the wavelength at which they were selected (for the B-band, this is of course only true up to $z\approx 3$ )."
+ For better colmparability of the FDF and COODS-South samples at the faiut-cud. we chose a consistent magnitude cut-off for all samples in Fie. l..," For better comparability of the FDF and GOODS-South samples at the faint-end, we chose a consistent magnitude cut-off for all samples in Fig. \ref{fig:sfr:cont_absmag_fdf_goods}."
+ This magnitude cut-off corresponds to the completeness limit of the GOODS-South sample aud is about oue magnitude brighter than the completeness lanits of theFDF D aud I sauiples., This magnitude cut-off corresponds to the completeness limit of the GOODS-South sample and is about one magnitude brighter than the completeness limits of theFDF B and I samples.
+ For the redshift bius defined bv the limits 0.15. 0.51. 1.21. 1.61. 2.13. 3.01. LOI. 5.01. the cut-offs in που ave at -15. -16. -17. -18. -19. -20. -20.," For the redshift bins defined by the limits 0.45, 0.81, 1.21, 1.61, 2.43, 3.01, 4.01, 5.01, the cut-offs in $M_{1500}$ are at -15, -16, -17, -18, -19, -20, -20."
+ The Misoy LEs aud the related parameters AM.. © and o of the D. E. and (|B) selected FDF samples are almost identical.," The $M_{1500}$ LFs and the related parameters $M_\ast$, $\Phi_\ast$ and $\alpha$ of the B, I, and (I+B) selected FDF samples are almost identical."
+ We derived consistent values for a 0.01) for all three samples considered here. similar to that deserihed in Gabaschetal.(2001).," We derived consistent values for $\alpha$ ) for all three samples considered here, similar to that described in \citet{gabasch:1}."
+. Consistent fünt-eud slopes (a=LOL+ 0.08) were obtained using a brighter subset of the data set. (.6. 1 nae brighter than the completeness limit)., Consistent faint-end slopes $\alpha = -1.01 \pm 0.08$ ) were obtained using a brighter subset of the data set (i.e. 1 mag brighter than the completeness limit).
+ Objects that were detected in only one baud and not im both are all faint aud do not contribute significautly to SER determined from the iuteeral over the LE., Objects that were detected in only one band and not in both are all faint and do not contribute significantly to SFR determined from the integral over the LF.
+ Cabaschetal.(2001). show that the steeper slope of other surveys is largely. due to shallower linitiug maguitudes., \citet{gabasch:1} show that the steeper slope of other surveys is largely due to shallower limiting magnitudes.
+ This is supported by an analysis of dropouts from COODS and the Hubble Ultra Deep Parallel Ficlds (Bowesetal.2001) where ana of -1.15 was derived., This is supported by an analysis of dropouts from GOODS and the Hubble Ultra Deep Parallel Fields \citep{bouwens:2004} where an $\alpha$ of -1.15 was derived.
+ Compared o LF parameters of the optically selected. samples. the Misogy LF of the I-selected suple has slishtlv brighter values of M... siguificautly lower values iu $. aud. within he large errors. a similar f£ünt-eud slope o.," Compared to LF parameters of the optically selected samples, the $M_{1500}$ LF of the K-selected sample has slightly brighter values of $M_\ast$ , significantly lower values in $\Phi_\ast$ and, within the large errors, a similar faint-end slope $\alpha$ ."
+ Since the slightly shallower Ix-selected sample does not allow us ο constrain the f£ünut-eud slope to the same level as our FDF sample (but is consistent with the faint cud slope a=1.07 determined for that field). we adopt this value or ourIy-selected. sample.," Since the slightly shallower K-selected sample does not allow us to constrain the faint-end slope to the same level as our FDF sample (but is consistent with the faint end slope $\alpha=-1.07$ determined for that field), we adopt this value for ourK-selected sample."
+ We examine the consequences of these findings for the, We examine the consequences of these findings for the
+The presence of bars in spiral galaxies is common (?) and the Milky Way is no exception.,The presence of bars in spiral galaxies is common \citep[][]{eskridge00} and the Milky Way is no exception.
+ Different studies based on various techniques have suggested that there is a large-scale bar with ~2.5 kpe in radius oriented 15-30 degrees with respect to the Galactic center line of sight with the near end towards positive longitudes.," Different studies based on various techniques have suggested that there is a large-scale bar with $\sim$ 2.5 kpc in radius oriented 15-30 degrees with respect to the Galactic center line of sight \citep[e.g.][]{dwek95,stanek97,bissantz02,babusiaux05,cabrera07,rattenbury07} with the near end towards positive longitudes."
+ However. observations of external galaxies have led to the conclusion that the existence of double bars ts not unusual (e.g??) and thus have provided important constraints on their scales.," However, observations of external galaxies have led to the conclusion that the existence of double bars is not unusual \citep[e.g][]{laine02,erwin11} and thus have provided important constraints on their scales."
+ Secondary bars are found to be generally small. with a typical size of ~500 pe or a relative length of about of that of the main bar (?)..," Secondary bars are found to be generally small, with a typical size of $\sim500$ pc or a relative length of about of that of the main bar \citep[][]{erwin11}."
+ Some observations 1n the very inner regions of the Milkv Way have indeed suggested that there is another structure inside the Galactic bar. namely a nuclear bar (e.g.22)..," Some observations in the very inner regions of the Milky Way have indeed suggested that there is another structure inside the Galactic bar, namely a nuclear bar \citep[e.g.][]{alard01,nishiyama05}."
+ ? used 2MASS star counts to show evidence of an excess in the inner regions of the Milky Way. which is interpreted as the presence of a small. lopsided. nuclear bar.," \citet[][]{alard01} used 2MASS star counts to show evidence of an excess in the inner regions of the Milky Way, which is interpreted as the presence of a small, lopsided, nuclear bar."
+ In particular. ?.hereafterNOS used the magnitude of red clump (RC) stars at a latitude of b=+1° to measure the properties of the inner regions (7107.«/ 107) detecting a shallower inclination angle for the Galactic bar at fields with longitudes |/]<4°.," In particular, \citet[][hereafter N05]{nishiyama05} used the magnitude of red clump (RC) stars at a latitude of $b=+1^\circ$ to measure the properties of the inner regions $-10^\circ<l<10^\circ$ ) detecting a shallower inclination angle for the Galactic bar at fields with longitudes $|l|<4^\circ$."
+ This was interpreted by NOS as the effect of a distinct inner bar structure., This was interpreted by N05 as the effect of a distinct inner bar structure.
+ The existence of such a nuclear bar has strong effects on the gas distribution in the inner parts of the Galaxy., The existence of such a nuclear bar has strong effects on the gas distribution in the inner parts of the Galaxy.
+ ? concluded that indeed a nuclear bar ts likely to be responsible for the observed so-called central molecular zone., \citet[][]{rodriguez08} concluded that indeed a nuclear bar is likely to be responsible for the observed so-called central molecular zone.
+ In this Letter. we make use of the ESO Vista Variables in the Via Lactea survey (VVV) to trace the magnitudes of RC giants in the longitude range -10°</10° at b.=41° to constrain the," In this Letter, we make use of the ESO Vista Variables in the Via Lactea survey (VVV) to trace the magnitudes of RC giants in the longitude range $-10^\circ<l<10^\circ$ at $b=\pm1^\circ$ to constrain the"
+We study the origin of the observed variability and search for spectral variations related to the changes in. X-ray brightness by determining the respective hardness ratio HR=(H-S)/(H+S) for each time bin. with § 00.7 keV and H 22.0 keV being the respective photon energy bands (see lower panel. reflc)).,"We study the origin of the observed variability and search for spectral variations related to the changes in X-ray brightness by determining the respective hardness ratio $HR=(H-S)/(H+S)$ for each time bin, with $S$ 0.7 keV and $H$ 2.0 keV being the respective photon energy bands (see lower panel, \\ref{lc}) )."
+ Spectral lines are unresolved in the ACIS spectrum. however the emission below 0.7 keV is dominated by cooler plasma with line peak formation temperatures in the range of —44 MK. while above kkeV lines from hotter plasma (z5 MK) dominate.," Spectral lines are unresolved in the ACIS spectrum, however the emission below 0.7 keV is dominated by cooler plasma with line peak formation temperatures in the range of 4 MK, while above keV lines from hotter plasma $\gtrsim 5$ MK) dominate."
+ If enhanced magnetic activity is the origi for the X-ray brightening. one expects a correlation of the X-ray brightness with the hardness of the emission.," If enhanced magnetic activity is the origin for the X-ray brightening, one expects a correlation of the X-ray brightness with the hardness of the emission."
+ The error oi the individual hardness ratio bins 1s about £0.25. but indeec harder emission is observed during X-ray brighter phases.," The error on the individual hardness ratio bins is about $\pm 0.25$, but indeed harder emission is observed during X-ray brighter phases."
+ A Spearman's rank correlation gives a chance probability of 0.07 that the correlation is due to a statistical fluctuation for the usec ten time bins., A Spearman's rank correlation gives a chance probability of 0.07 that the correlation is due to a statistical fluctuation for the used ten time bins.
+ Thus variable magnetic activity is the likely cause for the observed moderate X-ray brightness variability. whereas more extreme flaring events as observed on active stars are not present on HR 8799 during our observation.," Thus variable magnetic activity is the likely cause for the observed moderate X-ray brightness variability, whereas more extreme flaring events as observed on active stars are not present on HR 8799 during our observation."
+ We determine the global spectral properties of HR 8799 by fitting the ACIS spectrum with multi-temperature spectral models., We determine the global spectral properties of HR 8799 by fitting the ACIS spectrum with multi-temperature spectral models.
+ The metallicity of the coronal plasma can only be poorly constrained with the existing data and was fixed., The metallicity of the coronal plasma can only be poorly constrained with the existing data and was fixed.
+ Using the photospheric value of HR 8799. determined to 0.33 solar abundances (?).. already a one temperature model is sufficient to deseribe the data.," Using the photospheric value of HR 8799, determined to 0.33 solar abundances \citep{gray99}, already a one temperature model is sufficient to describe the data."
+ The corona therefore also seems to exhibit the photospheric sub-solar composition. but Adopting solar abundances results in correspondingly lower emission measure. albeit the quality of the fit is slightly poorer for a one temperature model.," The corona therefore also seems to exhibit the photospheric sub-solar composition, but Adopting solar abundances results in correspondingly lower emission measure, albeit the quality of the fit is slightly poorer for a one temperature model."
+ The corona therfore is better described by the photospheric sub-solar composition. but the indications are sparse; we note that a similarly good fit can be obtained with solar metallicity and two. however poorly constrained. temperature components.," The corona therfore is better described by the photospheric sub-solar composition, but the indications are sparse; we note that a similarly good fit can be obtained with solar metallicity and two, however poorly constrained, temperature components."
+ We show the X-ray spectrum and the best fit one temperature model in refspec.. the derived spectral properties are summarized in reffit..," We show the X-ray spectrum and the best fit one temperature model in \\ref{spec}, the derived spectral properties are summarized in \\ref{fit}."
+" From the best fit model we obtain an X-ray luminosity of Lx=1.3x107? erg/s in the 0.2—22.0 keV band. corresponding to a surface flux of logFx=5.1 cem"" ss7! and an activity level of log Lx/Ly4=—6.16. indicating that HR 8799 ts despite its youth a rather weakly active star."," From the best fit model we obtain an X-ray luminosity of $L_{\rm X}= 1.3 \times 10^{28}$ erg/s in the 2.0 keV band, corresponding to a surface flux of $\log F_{\rm X}= 5.1$ $^{-2}$ $^{-1}$ and an activity level of log $L_{\rm X}$ $L_{\rm bol} = -6.16$, indicating that HR 8799 is despite its youth a rather weakly active star."
+ Since HR 8799 is about an order of magnitude X-ray brighter than the average Sun. the X-ray flux at the position of the closest of the detected planets corresponcs to the solar X-ray flux that is received by Saturn.," Since HR 8799 is about an order of magnitude X-ray brighter than the average Sun, the X-ray flux at the position of the closest of the detected planets corresponds to the solar X-ray flux that is received by Saturn."
+ To cornpare our result with the RASS measurement. we also deternine the X-ray luminosity for the ROSAT kkeV bàid from the spectral model. leading to a value of Lx=1.6x107? erg/s. The ROSAT X-ray luminosity of Ly=1.5x107? erg/s. obtained from a hardness dependent energy conversion factor. is basically the same as the one obtained with more than 25 years later. pointing to the presence of long-term stable X-ray emission.," To compare our result with the RASS measurement, we also determine the X-ray luminosity for the ROSAT keV band from the spectral model, leading to a value of $L_{\rm X}= 1.6 \times 10^{28}$ erg/s. The ROSAT X-ray luminosity of $L_{\rm X} = 1.5 \times 10^{28}$ erg/s, obtained from a hardness dependent energy conversion factor, is basically the same as the one obtained with more than 25 years later, pointing to the presence of long-term stable X-ray emission."
+ The activity level as well as the X-ray luminosity of HR 8799 are about an order of magnittde larger than those of magnetically active mid to late A-type stars like Alderamin (vy Cep) or Altair (a AqD. both ¢ff spectral type A7.," The activity level as well as the X-ray luminosity of HR 8799 are about an order of magnitude larger than those of magnetically active mid to late A-type stars like Alderamin $\alpha$ Cep) or Altair $\alpha$ Aql), both of spectral type A7."
+ These stars have X-ray luminosities of o1e to à few times 107 ere/s and activity levels of log Lx/Lpo€-7., These stars have X-ray luminosities of one to a few times $10^{27}$ erg/s and activity levels of log $L_{\rm X}$ $L_{\rm bol}\lesssim -7$.
+ Compared to Altair. the hottest magnetically active star studied in detail at X-ray energies (?).. the X-ray surface flux of HR 8799 is by about a factor 20 higher: further the average coronal temperature of HR 8799 is with 3.0 MK slightly higher than those of Altair that is around 2.5 MK. but within the errors. this difference is not very significant.," Compared to Altair, the hottest magnetically active star studied in detail at X-ray energies \citep{rob09a}, the X-ray surface flux of HR 8799 is by about a factor 20 higher; further the average coronal temperature of HR 8799 is with 3.0 MK slightly higher than those of Altair that is around 2.5 MK, but within the errors this difference is not very significant."
+ This matches the finding that Altair's activity level is with log Lx/Loai—7.4 very low. even when compared to the weakly active stars HR 8799.," This matches the finding that Altair's activity level is with log $L_{\rm X}$ $L_{\rm bol} = -7.4$ very low, even when compared to the weakly active stars HR 8799."
+ However. the ultra-fast rotator Altair ts already around its maximum possible activity level and a significant spin-up would disrupt the star.," However, the ultra-fast rotator Altair is already around its maximum possible activity level and a significant spin-up would disrupt the star."
+ Consequently. the dynamo mechanism in HR 8799 needs to be more efficient than that in Altair.," Consequently, the dynamo mechanism in HR 8799 needs to be more efficient than that in Altair."
+" The efficiency of a solar-type dynamo is proportional to the inverse square of the Rossby-Number P/r,.). with P denoting the rotational period and τ, the convective turnover time."," The efficiency of a solar-type dynamo is proportional to the inverse square of the Rossby-Number $P/\tau_{c}$ ), with $P$ denoting the rotational period and $\tau_{c}$ the convective turnover time."
+ A significantly shorter rotation period for HR $799 is ruled out when considering stellar dimensions. Vsiii and geometry: thus a deeper convection zone would provide a natural explanation for the required efficiency.," A significantly shorter rotation period for HR 8799 is ruled out when considering stellar dimensions, $Vsini$ and geometry; thus a deeper convection zone would provide a natural explanation for the required efficiency."
+As the orbital phase advances. the spectra show rapid changes with the development of a pronounced hump bluesvard of 2.10 jam. aid an increasingly sleeper continuum bevond 2.2 yan. This blue hump is better defined in the NIRI spectra due to their extended blue coverage.,"As the orbital phase advances, the spectra show rapid changes with the development of a pronounced hump blueward of 2.10 $\mu$ m, and an increasingly steeper continuum beyond 2.2 $\mu$ m. This blue hump is better defined in the NIRI spectra due to their extended blue coverage."
+ Note that the Celluric water vapor absorption below 2.0 jan is verv strong. and (hus (the shape of the continuum here is highly sensitive to the telluric correction.," Note that the telluric water vapor absorption below 2.0 $\mu$ m is very strong, and thus the shape of the continuum here is highly sensitive to the telluric correction."
+ While the blueward rise begins well before this difficulty. the exact form of the slope of the continuum below 2.0 jm is not well-defined with the current data set.," While the blueward rise begins well before this difficulty, the exact form of the slope of the continuum below 2.0 $\mu$ m is not well-defined with the current data set."
+ As we have just noted (see Fig., As we have just noted (see Fig.
+ 5). combining the entire NIRSPEC data set has allowed us lo tentativelv detect the secondary star in EF Evi.," 5), combining the entire NIRSPEC data set has allowed us to tentatively detect the secondary star in EF Eri."
+ While the 5/N of these data are not hieh. if the secondary star were relatively normal. it should be more prominent (han seen here.," While the S/N of these data are not high, if the secondary star were relatively normal, it should be more prominent than seen here."
+ By modeling the optical spectrum of EF Eri. Deuermann et al. (," By modeling the optical spectrum of EF Eri, Beuermann et al. ("
+2000) have derived distances for EF Evi in the range 80 < d x 128 pc. while Thorstensen (2003) has measured a distance of 163 4 58 pe via parallax.,"2000) have derived distances for EF Eri in the range 80 $\leq$ d $\leq$ 128 pc, while Thorstensen (2003) has measured a distance of 163 $\pm$ 58 pc via parallax."
+ Thorstensen finds that an astrometric solution with a distance of 113 pe is possible by adjusting some of the input criteria., Thorstensen finds that an astrometric solution with a distance of 113 pc is possible by adjusting some of the input criteria.
+ As shown in llarrison οἱ al. (, As shown in Harrison et al. (
+2003). the optical photometry of EF Exi is well-modeled by a 9500 Ix white dwarl. and this white dwarf provides no more than of the flux in the A-band.,"2003), the optical photometry of EF Eri is well-modeled by a 9500 K white dwarf, and this white dwarf provides no more than of the flux in the -band."
+" If we assume (hal the secondary star is responsible for the remainder of the A-band flux al phase 0.0(A = 15.55). the absolute magnitude of EF Eri is in the range 10.0 < AI,— 11.0."," If we assume that the secondary star is responsible for the remainder of the -band flux at phase 0.0 = 15.55), the absolute magnitude of EF Eri is in the range 10.0 $\leq$ $_{\rm K} \leq$ 11.0."
+ This absolute magnitude corresponds to a spectral tvpe in the range M3 to L1 Genoring anv ellects due to the Roche lobe filling nature of this object)., This absolute magnitude corresponds to a spectral type in the range M8 to L1 (ignoring any effects due to the Roche lobe filling nature of this object).
+ A secondary star with a spectral tvpe near M9 should have strong CO features. as well as à detectable Na I doublet.," A secondary star with a spectral type near M9 should have strong CO features, as well as a detectable Na I doublet."
+ We have compared (he combined JF aud. A-band spectra of EF En with published spectra for stars and brown clwarls in the range AITV to Ts., We have compared the combined - and -band spectra of EF Eri with published spectra for stars and brown dwarfs in the range M7V to T8.
+ While, While
+with a distance modulus of10.03: we also correct for passive evolution(?).,with a distance modulus of$40.03$; we also correct for passive evolution.
+. The lage spectroscopic sample obtained for the CL 1358162 field was taken at the Multiple Mirror Telescope and the William Herschel Telescope., The large spectroscopic sample obtained for the CL 1358+62 field was taken at the Multiple Mirror Telescope and the William Herschel Telescope.
+ A detailed description of the reduction and analysis of the spectroscopic sanuple are described in? aud FOS., A detailed description of the reduction and analysis of the spectroscopic sample are described in and F98.
+ Twenty slit-inasks were designed to cover a 10«11' field: targets were chosen basedon their AR imagnuitude from imaging taken at the Whipple Observatory 1.214 telescope., Twenty slit-masks were designed to cover a $10'\times11'$ field; targets were chosen basedon their $R$ magnitude from imaging taken at the Whipple Observatory 1.2m telescope.
+ The spectral resolution was ~ΤΑat the WOT aud -~ at the MIAIT., The spectral resolution was $\sim13$at the WHT and $\sim20$ at the MMT.
+ Of the 387 redshifts ποατα QUAarecluster iienibers;," Of the 387 redshifts measured, 232 are cluster members."
+" Defining galaxies with 0.31161redshiftο-Le=0,31201 as cluster members. CL 1358|62s mica and velocity dispersion are 0.3283d0.0003 and 10273 respectively (E98)."," Defining galaxies with $0.31461<z<0.34201$ as cluster members, CL 1358+62's mean redshift and velocity dispersion are $0.3283\pm0.0003$ and $1027^{+51}_{-45}$ respectively (F98)."
+ From the redshift catalog. 192 cluster members fall ou thes«s TST ΑΕΡΟΣ mosaic.," From the redshift catalog, 192 cluster members fall on the $8'\times8'$ $HST$ WFPC2 mosaic."
+" The completeness witliu the LST iosaic is >90% to R—21 Gay,~ 20.2). decreasing to ~30% at Ro=22."," The completeness within the $HST$ mosaic is $>90\%$ to $R=21$ $m_{814}\sim20.2$ ), decreasing to $\sim30\%$ at $R=22$."
+ The incompleteness is nof due to an inability to measure redshifts at these mmaenitudes but by the limited umber of galaxies observed., The incompleteness is not due to an inability to measure redshifts at these magnitudes but by the limited number of galaxies observed.
+ The redshift success rate is not strongly color dependent for R<23.5 (F9s)., The redshift success rate is not strongly color dependent for $R<23.5$ (F98).
+ Of the 192 galaxies that are on the AST mosaic. we fit galaxy models only to 173 as 19 fall too close to a chip edge for proper analysis of the surface brightucss distribution.," Of the 192 galaxies that are on the $HST$ mosaic, we fit galaxy models only to 173 as 19 fall too close to a chip edge for proper analysis of the surface brightness distribution."
+ The 1723 galaxies range in absolute 0. magnitude from 20.8) images (Il375: BCC) to 15.9 mags (11829)., The 173 galaxies range in absolute $B_z$ magnitude from $-20.8$ mags (H375; BCG) to $-15.9$ mags (I1829).
+ We include the published [OIIA3727 line strength measmrements (F98) aud internal velocity dispersious for a subset in our analysis., We include the published $\lambda3727$ line strength measurements (F98) and internal velocity dispersions for a subset in our analysis.
+ We use the GIM2D package to fud the best-fit PSF-couvolved. 2D bulge|disk models to the surface brightuess profiles of the cluster members.," We use the GIM2D package to find the best-fit PSF-convolved, 2D bulge+disk models to the surface brightness profiles of the cluster members."
+ The program has aimaninuuu of 12 fitting parameters: the flux (Εν} iu the model integrated tor=x: the bulge/total luminosity οτη Protas tho somianajor axis effective radius of the bulge ο: the bulee elliptiitv e=1b/a where a and b are the bulee senü-nmiajor and semi-niünor axes respectively: the bulgeposition angle o5: the semianajor axis exponential disk scale leugth ry the iuclinatiou of the disk / (7=0 for face-on): the disk position angle ov: the subpixel die aud dy offsets of the ealaxys ceuter: the residual background level db: aud the Sévrsic iudex η of the bulge.," The program has a maximum of 12 fitting parameters: the flux $F_{total}$) in the model integrated to $r=\infty$; the bulge/total luminosity $B/T \equiv
+F_{bulge}/F_{total}$ ; the semi-major axis effective radius of the bulge $r_e$; the bulge ellipticity $e\equiv1-b/a$ where $a$ and $b$ are the bulge semi-major and semi-minor axes respectively; the bulge position angle $\phi_b$; the semi-major axis exponential disk scale length $r_d$; the inclination of the disk $i$ $i \equiv 0$ for face-on); the disk position angle $\phi_d$; the subpixel $dx$ and $dy$ offsets of the galaxy's center; the residual background level $db$; and the Sérrsic index $n$ of the bulge."
+" Both o5 and oa, are measured clockwise from the positive y-axis of the nage.", Both $\phi_b$ and $\phi_d$ are measured clockwise from the positive $y$-axis of the image.
+ The best-fit paramecters and them confidence intervals are deterumued using the Metropolis aleorithiu which uses the A test to determiue the region of maxima likelihood iu the parameter space., The best-fit parameters and their confidence intervals are determined using the Metropolis algorithm which uses the $\chi^2_{\nu}$ test to determine the region of maximum likelihood in the multi-parameter space.
+" The bulge profile is defined as where M(r) is the surface brightuess at ¢ along the seni-lnajor axis. aud X, is the effective surface brightucss,"," The bulge profile is defined as where $\Sigma(r)$ is the surface brightness at $r$ along the semi-major axis, and $\Sigma_e$ is the effective surface brightness."
+ This bulee profile is also kuowu as the Sévrsic profile(?)., This bulge profile is also known as the Sérrsic profile.
+. The pirezüueter & ds equal to (1.99920.3271). à value that defines +. to be the projected radius euclosiug half of the light in the bulee component(7).," The parameter $k$ is equal to $(1.9992n -
+0.3271)$, a value that defines $r_e$ to be the projected radius enclosing half of the light in the bulge component."
+. The classical de Vaucouleurs profile is a special case of Equation 3. with n=l., The classical de Vaucouleurs profile is a special case of Equation \ref{1358bulge-prof} with $n = 4$.
+ The disk profile is defined as where Ny is the (face-on) central surface briehtness., The disk profile is defined as where $\Sigma_0$ is the (face-on) central surface brightness.
+ We note (as do ??)) that the bulge/disk nomenclature adopted here to describe our surface brightuess models may not reflect the internal kinematics of ifs components., We note (as do ) that the bulge/disk nomenclature adopted here to describe our surface brightness models may not reflect the internal kinematics of its components.
+ A vbulec™ may not be a centralized. dynamically hot spheroid but could be a central starburst.," A “bulge” may not be a centralized, dynamically hot spheroid but could be a central starburst."
+ Similarly. a “disk” ay not necessarily be a cold. co-rotatiug population.," Similarly, a “disk” may not necessarily be a cold, co-rotating population."
+ For example. dvuamically “hot” systems such as faint dwiuf ellipticals are best fit bv exponential disks(?7).," For example, dynamically “hot” systems such as faint dwarf ellipticals are best fit by exponential disks."
+. Defore a 2D model to the surface brightuess of even galaxy can be fitted. the egalaxws isophotal area and an appropriate poiut spread function (PSF) ums be determined.," Before a 2D model to the surface brightness of a given galaxy can be fitted, the galaxy's isophotal area and an appropriate point spread function (PSF) must be determined."
+" To define the isophotal area. we use the ealaxv photometry package SExtractor V2.0 witha detection threshold of psy,=2L anaes/U"" (equivalen to leo of the skv noise) aud minimi debleudiug contrast parameter of 0.01."," To define the isophotal area, we use the galaxy photometry package SExtractor V2.0 with a detection threshold of $\mu_{814}=24.4$ $\Box''$ (equivalent to $\sim1\sigma$ of the sky noise) and a minimum deblending contrast parameter of 0.01."
+" As we fit galaxies in two filters (F606W and FSilw). we use the isophotal area. defiuec in the (redder) FaliW πμαρο,"," As we fit galaxies in two filters (F606W and F814W), we use the isophotal area defined in the (redder) F814W image."
+ To generate a PSF for cach ealaxy in each filter. we use Tiu TinVL.," To generate a PSF for each galaxy in each filter, we use TinyTim V4.4."
+As the PSF changes across cach WE chip. a PSF ος is generate every 50 pixels aud the nearest one to the galaxy is chosen for the GIM2D analysis.," As the PSF changes across each WF chip, a PSF model is generated every 50 pixels and the nearest one to the galaxy is chosen for the GIM2D analysis."
+ Both the PSF and galaxy models are subsampled by a factor of five as WE image are uudersampled aud the pixclization cau affect the shape of sinall galaxies such as those in our cluster suple., Both the PSF and galaxy models are subsampled by a factor of five as WF image are undersampled and the pixelization can affect the shape of small galaxies such as those in our cluster sample.
+ The centering of the galaxw is also improved by subsample the data: this point can be critical for determination of the galaxy residuals(?)., The centering of the galaxy is also improved by subsampling the data; this point can be critical for determination of the galaxy residuals.
+. By fitting modelsρα to the surface brightuess distribution of these galaxies. we ucasure the structural properties à. DB. reo ha. Op. Oy. and the half-lisht radiusryyo.," By fitting models to the surface brightness distribution of these galaxies, we measure the structural properties $n$, $B/T$, $r_e$, $r_d$, $\phi_b$ , $\phi_d$ , $i$ , and the half-light radius."
+. The semi-major axis halflieht radius is computed bv inteerating the sun of Equatious 3. aud | to r= x: uotethat as the models include bulge ellipticity aud disk inclination. circular svaiunietryv is nof assumed.," The semi-major axis half-light radius is computed by integrating the sum of Equations \ref{1358bulge-prof} and \ref{1358disk-prof} to $r=\infty$ ; notethat as the models include bulge ellipticity and disk inclination, circular symmetry is assumed."
+ The ealaxy’s Ge.gy) center is also determined frou the best fit mioclel.," The galaxy's $(x,y)$ center is also determined from the best fit model."
+ The asvunmetric duaee residual fiux is quantified by the asviuinietrv index Ba (2). defined as," The asymmetric image residual flux is quantified by the asymmetry index $R_{A}$ , defined as"
+Although these comprise an analysis of the scaling relations at low redshift. the method used to treat the Planck data is based on the assumption of a universal profile and scaling law with parameters derived from X-ray data alone.,"Although these comprise an analysis of the scaling relations at low redshift, the method used to treat the Planck data is based on the assumption of a universal profile and scaling law with parameters derived from X-ray data alone."
+ In these analyses. the measurements are restricted to the central part of clusters.," In these analyses, the measurements are restricted to the central part of clusters."
+ Our work shows that it should be possible to address the possible non-standard history of the thermal content of clusters. provided that the bias induced in Y by selection effects is properly taken into account.," Our work shows that it should be possible to address the possible non-standard history of the thermal content of clusters, provided that the bias induced in $Y$ by selection effects is properly taken into account."
+ In the wake of anew era for SZ cluster studies. the use of SZ flux measurements together with X-ray data can undoubtedly shed some light on the evolution of the properties of these massive objects. which can result in important consequences for cosmological applications.," In the wake of a new era for SZ cluster studies, the use of SZ flux measurements together with X-ray data can undoubtedly shed some light on the evolution of the properties of these massive objects, which can result in important consequences for cosmological applications."
+infall of fresh metal-poor material from. the LGAL,infall of fresh metal-poor material from the IGM.
+ Since he delay of SNla-cnrichnient is assumed to be independent of environment. the time at which trajectories first. move downward is independent: of radius.," Since the delay of SNIa-enrichment is assumed to be independent of environment, the time at which trajectories first move downward is independent of radius."
+ ὃν the fact that. the imescale of SNla-enrichment. does not. vary with racius either. the ISM trajectories tend to be nearly aligned.," By the fact that the timescale of SNIa-enrichment does not vary with radius either, the ISM trajectories tend to be nearly aligned."
+ Thus he point of turndown is at higher metallicities (further o the right) for populations closer to the Galactic centre. where the ISAT is enriched. faster by more intense star ormation relative to the present gas mass.," Thus the point of turndown is at higher metallicities (further to the right) for populations closer to the Galactic centre, where the ISM is enriched faster by more intense star formation relative to the present gas mass."
+ The colours and green. contours in show the density of stars within cach plane., The colours and green contours in show the density of stars within each plane.
+ In cach panel two ridges of high density are apparent one at high a/Ec]. which we call the metal-poor thick cise and one at low aLE]. which is associated with the thin disc.," In each panel two ridges of high density are apparent – one at high $\afe$, which we call the metal-poor thick disc and one at low $\afe$, which is associated with the thin disc."
+ Crucially. the disc ridge runs at a large angle to the black trajectories of the ISAL.," Crucially, the thin-disc ridge runs at a large angle to the black trajectories of the ISM."
+ Thus the thin-clise ridge in no sense traces the chemica history of the thin disc: insteac it reflects the spread in radii of birth of local stars. which gives rise to a spread in. Fe/1l by virtue of the metallicity eracicnt within the ISM (which is larger in the SBOO model than in traditional models of chemical evolution).," Thus the thin-disc ridge in no sense traces the chemical history of the thin disc; instead it reflects the spread in radii of birth of local stars, which gives rise to a spread in $\feh$ by virtue of the metallicity gradient within the ISM (which is larger in the SB09 model than in traditional models of chemical evolution)."
+ In a similar manner the thick-disk ridge follows the evolution of all rings at low metallicities. bu stretches significantly to higher. Fe/1l] than the point a which the solar annulus leaves it.," In a similar manner the thick-disk ridge follows the evolution of all rings at low metallicities, but stretches significantly to higher $\feh$ than the point at which the solar annulus leaves it."
+ The depression in the stellar density between the ridges in is à consequence of the rapid. downward motion of all trajectories after the onset of SNIa. and. of the ISM approaching a steady state as it enters the thin-disc ridge line: relatively few. stars are formed at intermediate values of a/Fe].," The depression in the stellar density between the ridges in is a consequence of the rapid downward motion of all trajectories after the onset of SNIa, and of the ISM approaching a steady state as it enters the thin-disc ridge line; relatively few stars are formed at intermediate values of $\afe$."
+ Variations in the timescales of enrichment. will change the depth of the depression for example. a shorter timescale for the ἄοσαν of SNla —progenitors will cause trajectories to move downwards Faster. leading to fewer stars in the intermediate region.," Variations in the timescales of enrichment will change the depth of the depression – for example, a shorter timescale for the decay of SNIa progenitors will cause trajectories to move downwards faster, leading to fewer stars in the intermediate region."
+ Our models use a prescription for SNla that is standard for models of chemical evolution. with no SNla until 0.15Civr after star formation. and then an exponential decay in the rate of SNla with time constant 1.5CGwvr.," Our models use a prescription for SNIa that is standard for models of chemical evolution, with no SNIa until $0.15 \Gyr$ after star formation, and then an exponential decay in the rate of SNIa with time constant $1.5 \Gyr$."
+" Mannuccictal.(2006). suggest that half. SNla explode promptly (within 0.1Civr of star formation) andthe rest explode at a rate that declines exponentially with time constant (ανν,", \cite{Mannucci06} suggest that $\sim$ half SNIa explode promptly (within $0.1\Gyr$ of star formation) andthe rest explode at a rate that declines exponentially with time constant $3\Gyr$.
+ TFhe existence of prompt SNla would not materially allect our work as long as a significant fraction of SNla are in the population with a long time constant. since he prompt SNla will lower the alpha-cnhancement level of the thick disc component. but not alfect the evolution )tween the two density ridges in the abundance. plane.," The existence of prompt SNIa would not materially affect our work as long as a significant fraction of SNIa are in the population with a long time constant, since the prompt SNIa will lower the alpha-enhancement level of the thick disc component, but not affect the evolution between the two density ridges in the abundance plane."
+ Forsteretal.(2006). showed that timescales are very weakly constrained by SNla counts due to uncertainties in the ormation histories., \cite{Foerster06} showed that timescales are very weakly constrained by SNIa counts due to uncertainties in the star-formation histories.
+ Since both Ca anc O are a-elements. the distributions in he upper and lower panels of are qualitatively similar. and O will be the only a-element explicitly discussed below.," Since both Ca and O are $\alpha$ -elements, the distributions in the upper and lower panels of are qualitatively similar, and O will be the only $\alpha$ -element explicitly discussed below."
+ As Haywood(2008) has pointed out. the principal racers of radial mixing are the [Large dispersion in the metallicities of stars in the solar neighbourhood and. the strong increase in this dispersion with age. which is caused »v immigration of hieh-metallicity stars from inwares ancl ow-metallicityv stars from outwards.," As \cite{Haywood08} has pointed out, the principal tracers of radial mixing are the large dispersion in the metallicities of stars in the solar neighbourhood and the strong increase in this dispersion with age, which is caused by immigration of high-metallicity stars from inwards and low-metallicity stars from outwards."
+ In. fact. as 91209 demonstrated. it is impossible to fit the shape of the local metallicity distribution under plausible assumptions without racial mixing.," In fact, as SB09 demonstrated, it is impossible to fit the shape of the local metallicity distribution under plausible assumptions without radial mixing."
+ As well as generating a large dispersion. in the metallicities of old stars. radial migration has a big impact on the interdependence of kinematies ancl chemistry.," As well as generating a large dispersion in the metallicities of old stars, radial migration has a big impact on the interdependence of kinematics and chemistry."
+ This impact is illustrated. by4.. which is another plot. of the (CFefH].O/Fe]) plane. but now with colour indicating the mean rotation velocity. of stars at. cach point blue indicates low rotation velocities and red large ones.," This impact is illustrated by, which is another plot of the $(\feh, \ofe)$ plane, but now with colour indicating the mean rotation velocity of stars at each point – blue indicates low rotation velocities and red large ones."
+ We see that at any given value of a/Ec]. there is a tight correlation between Fe/11] and rotation velocity in the sense that high. Fe/LH] implies low rotation velocity because stars with high Fe/L] are migrants from small radii and. tend to be deficient in angular momentum. while stars with low οΕΕ are migrants from large radii.," We see that at any given value of $\afe$, there is a tight correlation between $\feh$ and rotation velocity in the sense that high $\feh$ implies low rotation velocity because stars with high $\feh$ are migrants from small radii and tend to be deficient in angular momentum, while stars with low $\feh$ are migrants from large radii."
+ Phe black lines that show the trajectories of the ISM almost constitute contours of constant. mean rotation velocity. but there is a barely perceptible tendeney for the rotation velocity to decrease as one moves up along a black line.," The black lines that show the trajectories of the ISM almost constitute contours of constant mean rotation velocity, but there is a barely perceptible tendency for the rotation velocity to decrease as one moves up along a black line."
+ While the correlation between Ο/Η) and. rotation velocity seen in is qualitative consistent with stars xung scattered. to more eccentric orbits while retaining heir angular momenta (“blurring”). quantitatively changes in angular momentum (“churning”) play a big role in structuring4.," While the correlation between $\feh$ and rotation velocity seen in is qualitative consistent with stars being scattered to more eccentric orbits while retaining their angular momenta (“blurring”), quantitatively changes in angular momentum (“churning”) play a big role in structuring."
+". While churning does occasionally move he guicing-centre radius of a star from fe«Ry to RycLo and thus increase the mean rotation speed at large Fe/1I]. he dominant. οσοι of churning is to move guiding centres rom e,«Ry to Ry< such that avery metal-rich star is ound in the solar neighbourhood at a relatively low rotation velocity."," While churning does occasionally move the guiding-centre radius of a star from $R_{\rm g}<R_0$ to $R_{\rm g}>R_0$ and thus increase the mean rotation speed at large $\feh$, the dominant effect of churning is to move guiding centres from $R_{\rm g}\ll R_0$ to $R_{\rm
+g}<R_0$ such that a metal-rich star is found in the solar neighbourhood at a relatively low rotation velocity."
+ To illustrate the impact of churning quantitatively.," To illustrate the impact of churning quantitatively,"
+For the other lines listed in Table 2. the strongest is the O Ίνα absorption at 21.6 (Fig. 1)):,"For the other lines listed in Table 2, the strongest is the O $\alpha$ absorption at 21.6 (Fig. \ref{fig1}) );"
+ the corresponding O Ixy) is visible at 13.67 (Fig. 2))., the corresponding O $\beta$ is visible at 18.67 (Fig. \ref{fig2}) ).
+ The 21.6 feature has already been seen in MB. 421 itself (Nicastvoetal.2001) and in other 4 AGNs Canizares2003) and it is usually attributed to a WIIIM within our local group ol galaxies (e.g.Nicastroοἱal.2002). or to radiativelv cooling gas inside our Galaxy Heckmanetal. 2002).," The 21.6 feature has already been seen in MRK 421 itself \citep{nic01} and in other 4 AGNs \citep{nic02,fan02,cag03,mat02,fan03} and it is usually attributed to a WHIM within our local group of galaxies \citep[e.g.][]{nic02}
+ or to radiatively cooling gas inside our Galaxy \citep[e.g.][]{hec02}."
+. I also report the detection of zero redshift O Lvo at 18.97 (Fig., I also report the detection of zero redshift O $\alpha$ at 18.97 (Fig.
+ 2) and of Ίνα at 13.4 (Fig., 2) and of $\alpha$ at 13.4 (Fig.
+ 3)., 3).
+ Note however that these lines are affected. by small features in the effective The EW of the zero redshilt leatures (O νιο and Nerx) are compatible with those measured with lor PINS2155-304 CNicastroetal.2002:Fang2002) and [or Mix 421 CNicastroetal.2001).," Note however that these lines are affected by small features in the effective The EW of the zero redshift features (O ,O and ) are compatible with those measured with for PKS2155-304 \citep{nic02,fan02} and for MRK 421 \citep{nic01}."
+. Ia particular. the EW derived fixing the lines FWA to ihe value of the 21.6 line. (i.e. 0.33. 0.27 and 0.20 eV respectively) are consistent. with the lines being unsaturated and with column densities of ~107 atoms 7 per ion species (seeFig.4inNicastroetal. 2002).," In particular, the EW derived fixing the lines FWHM to the value of the 21.6 line, (i.e. 0.33, 0.27 and 0.20 eV respectively) are consistent with the lines being unsaturated and with column densities of $\sim 10^{16}$ atoms $^{-2}$ per ion species \citep[see Fig.~4 in][]{nic02}. ."
+. For a gas density of ~10° atoms *. and OQ ratios indicate gas temperatures of e(0.4—2)xLO? K. and ~(1—4)10° K. respectively (Fig.5ofNicastroοἱal.2002).," For a gas density of $\sim 10^{-6}$ atoms $^{-3}$ , and O ratios indicate gas temperatures of $\sim (0.4-2) \times 10^6$ K, and $\sim (1-4) 10^6$ K, respectively \citep[Fig. 5 of][]{nic02}."
+. These ranges move to e(1.5—2.5)x105 Ik and ~(3—4)x109 IX in case of Galactic gas densitv., These ranges move to $\sim (1.5-2.5) \times 10^6$ K and $\sim (3-4) \times 10^6 $ K in case of Galactic gas density.
+ Even if these estimates have to be regarded! with caution because of the possible modifications of the effective area features. they are all consistent wilh the range of temperatures predicted for the WILIIM.," Even if these estimates have to be regarded with caution because of the possible modifications of the effective area features, they are all consistent with the range of temperatures predicted for the WHIM."
+ ] report. strong evidence of WIIIM. outside our local group. of galaxies (hirough the detection of a 21.82 O Ίνα absorption line in the RGS-1 spectrum of MRI 421., I report strong evidence of WHIM outside our local group of galaxies through the detection of a 21.82 O $\alpha$ absorption line in the RGS-1 spectrum of MRK 421.
+ The line redshift is ~0.01. in agreement with that derived for the strongIL Lva absorption line detected by HIST. and related (o an absorber in a cosmic void.," The line redshift is $\sim 0.01$, in agreement with that derived for the strongH $\alpha$ absorption line detected by HST and related to an absorber in a cosmic void."
+ The three closest. galaxies are al 2.14. 3.98 and 4.12 fh.) Mpe (Penton.Stocke.&Shull2000).," The three closest galaxies are at 2.14, 3.98 and 4.12 $h_{70}^{-1}$ Mpc \citep{pen00}."
+.. This detection is a firma prool the existence of a WILIIM filament al 2e0.01 possibly connecting the nearby galaxies with other denser regions of the skv. as expected by the models of the cosnuc structures formation and evolution.," This detection is a firm proof the existence of a WHIM filament at $z\sim 0.01$ possibly connecting the nearby galaxies with other denser regions of the sky, as expected by the models of the cosmic structures formation and evolution."
+" Iderive anupper limit on the gas temperature of ~2.5—4x10° IX lor a gas density of10.""—1 atom .", Iderive anupper limit on the gas temperature of $\sim 2.5-4 \times 10^6$ K for a gas density of$10^{-6} - 1$ atom $^{-3}$ .
+ The only known similar, The only known similar
+lt is becoming increasingly clear that Population LLL star formation is not as cifferent [rom that in the present clay as was once thought.,It is becoming increasingly clear that Population III star formation is not as different from that in the present day as was once thought.
+ Several authors have now shown that insteacl of a single massive star being formed at the centre ofa primordial halo. a small eluster is formed Machidaüs:: 7. 2011a.b: 2??73).," Several authors have now shown that instead of a single massive star being formed at the centre of a primordial halo, a small cluster is formed \\citealt{Machida08}; ; \citeauthor{Clark11a}~ 2011a,b; \citealt{Turk09,Stacy10,Prieto11}) )."
+ There are frequent dvnamical interactions within these clusters that can stronely alter the aceretion rates of the protostars (?).., There are frequent dynamical interactions within these clusters that can strongly alter the accretion rates of the protostars \citep{Greif11}.
+ As such. Population UL star formation begins to resemble simulations of local clustered star formation (c.g???)..," As such, Population III star formation begins to resemble simulations of local clustered star formation \citep[e.g][]{Bonnell11,Offner10,Girichidis11}."
+ For a review of present day clustered star formation see ?.., For a review of present day clustered star formation see \citet{MacLow04}.
+ A parallel strand. in the study of Population ILL star formation has been the investigation of the stellar evolution of the protostar as ib grows in mass., A parallel strand in the study of Population III star formation has been the investigation of the stellar evolution of the protostar as it grows in mass.
+ As the protostar grows. the consequences of its protostellar evolution become increasingly significant.," As the protostar grows, the consequences of its protostellar evolution become increasingly significant."
+ For example. during the protostellar stage of the evolution. radiation from the protostar alters the thermal ancl chemical structure of the eircumstellar disk.," For example, during the protostellar stage of the evolution, radiation from the protostar alters the thermal and chemical structure of the circumstellar disk."
+ As a result. the nature of disk fragmentation via the gravitational instability is alfectecl by the protostellar radiation (7. 2011b: 7)).," As a result, the nature of disk fragmentation via the gravitational instability is affected by the protostellar radiation \citeauthor{Clark11b}~ 2011b; \citealt{Smith11b}) )."
+ As the star grows in mass ancl approaches the main sequence. ib also starts to emit UV radiation.," As the star grows in mass and approaches the main sequence, it also starts to emit UV radiation."
+ This creates an LLL region which can prevent further mass accretion onto the star and limit the final stellar mass (7: ? 2011b: ?))., This creates an HII region which can prevent further mass accretion onto the star and limit the final stellar mass \citealt{McKee08}; \citeauthor{Hosokawa11b}~ 2011b; \citealt{Stacy11c}) ).
+ The stellar evolution also determines the protostellar radius. which is an important quantity [or understanding the cdvnamic evolution of a voung cluster. as this determines the likelihood of protostellar collisions and mergers (Lor an analysis of present clay protostellar clusters see ?)).," The stellar evolution also determines the protostellar radius, which is an important quantity for understanding the dynamic evolution of a young cluster, as this determines the likelihood of protostellar collisions and mergers (for an analysis of present day protostellar clusters see \citealt{Baumgardt11}) )."
+ Aluch of the pioneering. work on the stellar evolution of Population LLL protostars was done by Ὁ and ?? who numerically solved. the interior structure equations assuming various constant accretion rates and showed that the evolution quantitatively changes with dillerent. rates.," Much of the pioneering work on the stellar evolution of Population III protostars was done by \citet{Stahler86a} and \citet{Omukai01,Omukai03} who numerically solved the interior structure equations assuming various constant accretion rates and showed that the evolution quantitatively changes with different rates."
+ llowever. as was pointed out in ?.hereafterS11.. the protostars in a Population HIE. minihalo do not accrete ab a constant rate. but at a highly variable one. as the environment from which they acerete is not uniform.," However, as was pointed out in \citet[][hereafter S11]{Smith11b}, the protostars in a Population III minihalo do not accrete at a constant rate, but at a highly variable one, as the environment from which they accrete is not uniform."
+ Time-variable mass accretion is also ubiquitous in present-dav star formation (e.g.2??7) and its effects on protostellar evolution have begun to be explored (222: ? 2011a: ?2)).," Time-variable mass accretion is also ubiquitous in present-day star formation \citep[e.g.][]{Schmeja04,Vorobyov06,Machida11} and its effects on protostellar evolution have begun to be explored \citealt{Wuchterl01,Baraffe09,Baraffe10}; \citeauthor{Hosokawa11}~ 2011a; \citealt{Hartmann11}) )."
+ In this paper. we seek to combine the strands of hydrodynamic simulations of fragmentation with stellar modelling and study how time-variable mass acceretion alfects the evolution of primordial protostars.," In this paper, we seek to combine the strands of hydrodynamic simulations of fragmentation with stellar modelling and study how time-variable mass accretion affects the evolution of primordial protostars."
+ We use. protostellar aceretion data taken from the simulations of SII., We use protostellar accretion data taken from the simulations of S11.
+" These are smoothed particle hivelrodyvnamies (SPILL)re-simulations of five. cosmological— minihalos presented in ? in which the heating of the eas bv the accretion luminosity produced. by acereting ""sink", These are smoothed particle hydrodynamics (SPH)re-simulations of five cosmological minihalos presented in \citet{Greif11} in which the heating of the gas by the accretion luminosity produced by accreting `sink
+roughly match the observations of Coles and Harmon (1989) in coronal holes.,roughly match the observations of Coles and Harmon (1989) in coronal holes.
+ Total density fluctuation spectra (solid) are shown for f=0.1.1.10. with the decomposition into passive-scalar component (dotted) and “active” KAW component (dashed) included for each case.," Total density fluctuation spectra (solid) are shown for $f=0.1,1,10$, with the decomposition into passive-scalar component (dotted) and “active” KAW component (dashed) included for each case."
+ The results in Figure 2. are similar to those in Figure | except that the active KAW component of the density fluctuations is more prominent at low β;.," The results in Figure \ref{fig:dn_cor} are similar to those in Figure \ref{fig:dn_sw}
+ except that the active KAW component of the density fluctuations is more prominent at low $\beta_i$."
+" The questionable assumption of balanced (zero cross-helicity) turbulence in coronal holes is discussed in detail at the end of this section,", The questionable assumption of balanced (zero cross-helicity) turbulence in coronal holes is discussed in detail at the end of this section.
+ Figure 3 focuses on the density fluctuations due to linear KAWs. showing results as a function kp; for several D;.," Figure \ref{fig:beta} focuses on the density fluctuations due to linear KAWs, showing results as a function $k_\perp \rho_i$ for several $\beta_i$."
+" For B<I. the relation between the density and velocity fluctuations of à KAW at wavevector k. denoted 6; and àv,;. can be derived analytically and is given by (Hollweg 1999, Eqn."," For $\beta \lesssim 1$, the relation between the density and velocity fluctuations of a KAW at wavevector ${\bf k}$, denoted $ \delta n_{k}$ and $\delta v_{ \perp k}$ , can be derived analytically and is given by (Hollweg 1999, Eqn."
+ 51) assuming that ky>>[Kj] and 6<Q;. where @ is the wave frequency and y; is the adiabatic index of the ions.," 51) assuming that $k_\perp \gg |k_\parallel|$ and $\omega \ll \Omega_{i}$, where $\omega$ is the wave frequency and $\gamma_i$ is the adiabatic index of the ions."
+ In addition to showing the (normalized) ratio of density fluctuations to velocity fluctuations from linear kinetic theory (solid). Figure 3 also shows the values predicted by equation (2)) (dotted) with y;=|.," In addition to showing the (normalized) ratio of density fluctuations to velocity fluctuations from linear kinetic theory (solid), Figure \ref{fig:beta} also shows the values predicted by equation \ref{eq:deltan0}) ) (dotted) with $\gamma_i=1$."
+ This figure demonstrates that the density fluctuations are suppressed by a factor of ~p when D21. as expected because the fluctuations become increasingly incompressible for high D.," This figure demonstrates that the density fluctuations are suppressed by a factor of $\sim \beta$ when $\beta \gtrsim 1$, as expected because the fluctuations become increasingly incompressible for high $\beta$."
+ An additional change in the predicted density fluctuations for D=| — as can occur in the solar wind at ~| AU —- is that the damping time of slow waves and entropy modes becomes shorter than their expected cascade times.," An additional change in the predicted density fluctuations for $\beta
+\gtrsim 1$ – as can occur in the solar wind at $\sim 1$ AU – is that the damping time of slow waves and entropy modes becomes shorter than their expected cascade times."
+ Because the solar wind is collisionless. this occurs even on large scales.," Because the solar wind is collisionless, this occurs even on large scales."
+ As a result. the passive contribution to the density. fluctuations may be significantly diminished when B21 (Lithwick Goldreich 2001).," As a result, the passive contribution to the density fluctuations may be significantly diminished when $\beta \gtrsim 1$ (Lithwick Goldreich 2001)."
+ Equation (2)) can be used to estimate the ratio between the passive-scalar density fluctuation and the density fluctuation arising from KAW compressions., Equation \ref{eq:deltan0}) ) can be used to estimate the ratio between the passive-scalar density fluctuation and the density fluctuation arising from KAW compressions.
+" We assume that /8v;, equals the ratio 8,/8v; for linear KAWs bngiven by equation (2))."," We assume that $\delta n_k^{\rm
+(comp)}/\delta v_{\perp k}$ equals the ratio $\delta n_{k}/\delta v_{\perp k}$ for linear KAWs given by equation \ref{eq:deltan0}) )."
+" When B=| and the cross helicity is zero. onOv, is independent of & within the inertial range."," When $\beta
+\lesssim 1$ and the cross helicity is zero, $\delta n_k^{ \rm
+ (ps)}/\delta v_{\perp k}$ is independent of $k$ within the inertial range."
+" Assuming that onÓv|, remains roughly constant from the outer scale all the way to ky—p;|. equation (2)) implies that for balanced turbulence with D<1 At | AU. the value of én.πρ atk=ky ranges from roughly 0.1 to 0.3. as the assumed outer-scale time scale (as measured in the spacecraft frame) of the turbulence is varied from |12 hrs (Tu Marsch 1995. Fig."," Assuming that $\delta n_k^{ \rm (ps)}/\delta v_{\perp k}$ remains roughly constant from the outer scale all the way to $k_\perp \sim
+\rho_i^{-1}$, equation \ref{eq:deltan0}) ) implies that for balanced turbulence with $\beta \lesssim 1$ At 1 AU, the value of $\delta n_k^{ \rm (ps)}/n_0$ at $k=k_0$ ranges from roughly 0.1 to 0.3, as the assumed outer-scale time scale (as measured in the spacecraft frame) of the turbulence is varied from $1-12$ hrs (Tu Marsch 1995, Fig."
+ 2-14)., 2-14).
+ If we set v4=77 km/s as in Celnikier's period II data and ὃν230 km/s at k=ky at | AU (see Cranmer van Ballegooijen 2005. Fig.9). then f~0.070.6.," If we set $v_{\rm A} =
+77$ km/s as in Celnikier's period II data and $\delta v_{\perp k} 
+\simeq 30 $ km/s at $k=k_0$ at 1 AU (see Cranmer van Ballegooijen 2005, Fig.9), then $f \simeq 0.07 - 0.6$."
+" Equation (3)) thus suggests that KAW compressions usually dominate the density fluctuations on small seales (&,p; I) at 1 AU.", Equation \ref{eq:rel}) ) thus suggests that KAW compressions usually dominate the density fluctuations on small scales $k_\perp \rho_i \sim 1$ ) at 1 AU.
+ It is important to note several theoretical uncertainties in our predictions for the density fluctuations that should accompany low frequeney Alfvénnic turbulence in the solar wind., It is important to note several theoretical uncertainties in our predictions for the density fluctuations that should accompany low frequency Alfvénnic turbulence in the solar wind.
+ First. the relative amplitudes of the slow waves. entropy modes. and Alfvénn waves on large scales are not precisely known.," First, the relative amplitudes of the slow waves, entropy modes, and Alfvénn waves on large scales are not precisely known."
+ Therefore. we present several values of f in Figures | and 2..," Therefore, we present several values of $f$ in Figures \ref{fig:dn_sw} and \ref{fig:dn_cor}. ."
+" A second. and more significant. limitation is that we have assumed that the turbulence is ""balanced"" (zero cross-helicity). ie.. that the energy in waves propagating towards the sun (in the solar-wind frame) equals the energy in waves propagating away from the sun."," A second, and more significant, limitation is that we have assumed that the turbulence is “balanced” (zero cross-helicity), i.e., that the energy in waves propagating towards the sun (in the solar-wind frame) equals the energy in waves propagating away from the sun."
+ This is generally not observed to be true at ~| AU (particularly in thefast wind: Grappin 1990; Tu Marsch 1990)., This is generally not observed to be true at $\sim 1$ AU (particularly in thefast wind; Grappin 1990; Tu Marsch 1990).
+ Closer to the Sun there is an even larger difference between the energies of Sunward and anti-Sunward waves (Roberts etal 1987: Bavassano et al 2000)., Closer to the Sun there is an even larger difference between the energies of Sunward and anti-Sunward waves (Roberts etal 1987; Bavassano et al 2000).
+" Non-zero cross helicity or ""imbalance"" affects strong Alfvénnie. turbulence in several", Non-zero cross helicity or “imbalance” affects strong Alfvénnic turbulence in several
+'ejonization occurred.,reionization occurred.
+ Less massive galaxies like spirals aud field galaxies collapsed later. so fewer Xue GCs were produced iu them before GC formation was halted.," Less massive galaxies like spirals and field galaxies collapsed later, so fewer blue GCs were produced in them before GC formation was halted."
+ The natural result is that more uassive (aud thus more luminous) galaxies will have more blue. metal-poor GCs per unit mass tha ess massive objects.," The natural result is that more massive (and thus more luminous) galaxies will have more blue, metal-poor GCs per unit mass than less massive objects."
+ Within the framework of this scenario. tle observed spatial distributions of the uetal-poor GCs cau help place coustraints ou the redshifts at which the protogalactic fragiuents collapsed. as well as on the masses of the fragments.," Within the framework of this scenario, the observed spatial distributions of the metal-poor GCs can help place constraints on the redshifts at which the protogalactic fragments collapsed, as well as on the masses of the fragments."
+ In. hierarchical models. the most overcleuse ragments that collapse at the highest redshifts (and can begin to form GCs) are more than those with lower deusities that collapse later.," In hierarchical models, the most overdense fragments that collapse at the highest redshifts (and can begin to form GCs) are more centrally-concentrated than those with lower densities that collapse later."
+ Likewise. fragments witli more mass will also collapse earlier and be more concentrated toward the galaxy ceuter.," Likewise, fragments with more mass will also collapse earlier and be more concentrated toward the galaxy center."
+ Santos Las simulatect this iu detail for the Milky Way GC system. but has not vet produced testable preclictious for the GC spatial distributious of other galaxies.," Santos has simulated this in detail for the Milky Way GC system, but has not yet produced testable predictions for the GC spatial distributions of other galaxies."
+ La future. however. our observed GC spatial proliles lor ealaxies with a variety of masses aud iu different envirorunents will provide valuable coustraints [οἱ this theoretical work.," In future, however, our observed GC spatial profiles for galaxies with a variety of masses and in different environments will provide valuable constraints for this theoretical work."
+ We have uudertakeu a wicle-lield CCD survey of the GC systems of a large sample of spiral aud elliptical galaxies with tLe primary goal of testing models for the origin of ellipticals. tuchidine the idea that they are formed via spiral-spiral mergers.," We have undertaken a wide-field CCD survey of the GC systems of a large sample of spiral and elliptical galaxies with the primary goal of testing models for the origin of ellipticals, including the idea that they are formed via spiral-spiral mergers."
+ We have completed our study of NCC 781I. au Sab spiral that was [oui in past work to have au extremely high. GC specific frequeucy.," We have completed our study of NGC 7814, an Sab spiral that was found in past work to have an extremely high GC specific frequency."
+ Our inain results are as follows: ]., Our main results are as follows: 1.
+ Our data provide racial coverage of NGC 7810s GC system to 7 (230 kpe) from the ealaxy center., Our data provide radial coverage of NGC 7814's GC system to $\sim$ $\arcm$ $\sim$ 30 kpc) from the galaxy center.
+ The surface «leusity in the racial distribution of GCs is cousistent with zero by ~3- 1 (11-12 kpc) from tje galaxy center. suggesting that we have οerved the full extent of NGC Talls GC system.," The surface density in the radial distribution of GCs is consistent with zero by $\sim$ $-$ $\arcm$ $\sim$ $-$ 15 kpc) from the galaxy center, suggesting that we have observed the full extent of NGC 7814's GC system."
+ 9., 2.
+ NGC 7811 hasatota population of ~200 CC's aud a IumOsitv-normalized specific [requeney Sv of 1.3 EO.L., NGC 7814 has a total population of $\sim$ 200 GCs and a luminosity-normalized specific frequency $S_N$ of 1.3 $\pm$ 0.4.
+ NGC το is uot a lieh-Sa spiral as previously clalmed. but iustead has a GC system that is comparable o those of the two most well-sttdied spirals. the Milky Way and M31.," NGC 7814 is not a $S_N$ spiral as previously claimed, but instead has a GC system that is comparable to those of the two most well-studied spirals, the Milky Way and M31."
+ 3., 3.
+ A natu alcOLsequewe of the fact tha NGC τοι has many [fewer GCs than oeviously hought is tha the 1umber of metal-poor QCs must also )e substantialy reduced., A natural consequence of the fact that NGC 7814 has many fewer GCs than previously thought is that the number of metal-poor GCs must also be substantially reduced.
+ We compare (GC 781Ls blue €C popuatiou to that of ellidticals [rom our studs iu order to test a oredictiou of the Ashimat Zeyf iue‘eer model., We compare NGC 7814's blue GC population to that of ellipticals from our study in order to test a prediction of the Ashman Zepf merger model.
+ We fi that mergiig the blue GC populatious oL spirals ike NGC 7811 Ine10wh to account for the e GC's ina uoderate-Iumiuosity field elliptical like (GC 3379. bu 110 ist La giant cluster elliptic: like NGC |172.," We find that merging the blue GC populations of spirals like NGC 7814 is enough to account for the blue GCs in a moderate-luminosity field elliptical like NGC 3379, but not in a giant cluster elliptical like NGC 4472."
+ This suggests that simpe Inergers uay be able ο expal ithe origius of tvpical-inosity ellipticals but tha some aclclitional process is ueeded to accouil or the GC systems of1‘e massive ellipticals., This suggests that simple mergers may be able to explain the origins of typical-luminosity ellipticals but that some additional process is needed to account for the GC systems of more massive ellipticals.
+ Our fiudiugs lor NGC 781E significantly aler the overall picture of extagalactic GC systems., Our findings for NGC 7814 significantly alter the overall picture of extragalactic GC systems.
+ Iu Paper HI of this se‘ies. we will present resulS or the remainiug early-tyye galaxies iu the sample," In Paper III of this series, we will present results for the remaining early-type galaxies in the sample"
+width measurement depending on the location of the pseudo-continuum level (Pavlenko 2002. priv.,"width measurement depending on the location of the pseudo-continuum level (Pavlenko 2002, priv."
+ comm.)., comm.).
+ When our Li curves of growth are compared with. Pavlenko's (2002. priv.," When our Li curves of growth are compared with Pavlenko's (2002, priv."
+ comm.), comm.)
+ model data which also takes the effect of TiO bands into account. our curves of growth differ by 15mmá at common Li abundance points (logN(L/)=1.0& 1.3).," model data which also takes the effect of TiO bands into account, our curves of growth differ by $\lesssim15$ at common Li abundance points $\log N(Li)=1.0\,\&\,1.3$ )."
+ Using the estimated effective temperature of the secondary from section 2.2.1 (3030 K) and the measured equivalent width of 290mÁ.. from Figure 2. we conservatively estimate the Li abundance of the secondary to lie in the range logNCLi)=0.85—1.10.," Using the estimated effective temperature of the secondary from section \ref{Teff}
+ (3030 K) and the measured equivalent width of 290, from Figure \ref{LiCurve} we conservatively estimate the Li abundance of the secondary to lie in the range $\log N(Li)=0.85-1.10$."
+ If we adopt (Li)=3.31 as the interstellar Li abundance. then HIP 112312 A and B currently have only <0.01. and 0.35-0.62. of their initial Li contents.," If we adopt n(Li)=3.31 as the interstellar Li abundance, then HIP 112312 A and B currently have only <0.01 and 0.35-0.62 of their initial Li contents."
+ It is quite straightforward to calculate luminosities of the HIP 112312 binary system from Hipparcos measured parallax 42.44+3.4ο mas) and measured photometric magnitudes., It is quite straightforward to calculate luminosities of the HIP 112312 binary system from Hipparcos measured parallax 3.4 mas) and measured photometric magnitudes.
+ In calculation of bolometric magnitudes. we use 4—M=1.86 for the distance modulus. =4.75 for the Sun. bolometric correction of 2.73 and 2.77 magnitude at K band for the primary and the secondary. respectively (Bessellοἱo£.," In calculation of bolometric magnitudes, we use m-M=1.86 for the distance modulus, =4.75 for the Sun, bolometric correction of 2.73 and 2.77 magnitude at K band for the primary and the secondary, respectively \citep{BessellBC}."
+1998).. The calculated bolometric magnitudes are 7.78 8.68 and the calculated Juminosities of A B components are ..)=—1.22 and —1.58.. respectively.," The calculated bolometric magnitudes are 7.78 8.68 and the calculated luminosities of A B components are })=-1.22 and -1.58, respectively."
+ Uncertainty in luminosity is estimated to be Alog(Z/L..)=0.160 which is mainly due to the Hipparcos distance uncertainty., Uncertainty in luminosity is estimated to be )=0.160 which is mainly due to the Hipparcos distance uncertainty.
+ Physical companionship of the HIP 112312 components is verified from the common proper motions of the primary and the secondary (Table 2)., Physical companionship of the HIP 112312 components is verified from the common proper motions of the primary and the secondary (Table \ref{data}) ).
+ The proper motions listed in Table 2. are weighted means of SPM 2.0 (Plataisera£.1998) and UCACI (Zachariasetαἱ.2000) catalog values.," The proper motions listed in Table \ref{data}
+ are weighted means of SPM 2.0 \citep{SPM2} and UCAC1 \citep{UCAC1}
+ catalog values."
+ Proper motions of the primary from Hipparcos catalog 2.50 mas/yr. 2.21 mas/yr) are in fair agreemer= with the Table 2 values.," Proper motions of the primary from Hipparcos catalog 2.50 mas/yr, 2.21 mas/yr) are in fair agreement with the Table 2 values."
+ From our echelle spectra. we determined weighted mean radial velocities of the primary and secondary from a line fitting method (see Table 2)).," From our echelle spectra, we determined weighted mean radial velocities of the primary and secondary from a line fitting method (see Table \ref{data}) )."
+ We also estimated the rotational velocity /)) of the primary to be 40. km/sec) but were unable to do so for the secondary because of its low S/N spectrum., We also estimated the rotational velocity ) of the primary to be 40 km/sec) but were unable to do so for the secondary because of its low S/N spectrum.
+ Using the measured radial velocities. Hipparcos distances. and proper motions. we derived galactic space motions (U.V.W)) with respect to the Sun of HIP 112312 by using the formulation of Johnson&Soderblom(1987): k&mi/sec for the primaryV. and11.83: (—13.8=1.5.—19.1]+1.6.—7.22.2)mfsec for the secondary.," Using the measured radial velocities, Hipparcos distances, and proper motions, we derived galactic space motions U,V,W ) with respect to the Sun of HIP 112312 by using the formulation of \citet{Johnson87}; km/sec for the primary and $(-13.8\pm1.5, -19.1\pm1.6,-7.2\pm2.2)\,km/sec$ for the secondary."
+ Based on the similar space motions. proximity to Earth (24 pc). and very young ages 20 Myr. see Figure 3 and discussion in section 2.3)). the HIP 112312 binary components may be members of the Pictoris moving group (Zuckermanetaf.2001).," Based on the similar space motions, proximity to Earth (24 pc), and very young ages 20 Myr, see Figure \ref{HRD} and discussion in section \ref{ages}) ), the HIP 112312 binary components may be members of the Pictoris moving group \citep{bPic2}."
+. The Pictoris group was defined as nearby stars that have UVW velocity each within 2. km/sec of those of Pictoris. (-10.8.—-16.4.—-8.9).. and age consistent with 12 Myr (Zuckermanetal.2001).," The Pictoris group was defined as nearby stars that have UVW velocity each within 2 km/sec of those of Pictoris, (-10.8,-16.4,-8.9), and age consistent with 12 Myr \citep{bPic2}."
+. We compare pre-main sequence (PMS) stellar evolutionary models with our observed HIP 112312 binary data by plotting HIP 112312 A and Bon a versus ..) plot along with three PMS models (Figure 3))., We compare pre-main sequence (PMS) stellar evolutionary models with our observed HIP 112312 binary data by plotting HIP 112312 A and B on a versus ) plot along with three PMS models (Figure \ref{HRD}) ).
+ The high luminosities of A and B clearly indicate their young age (<<20. Myr): however the scatter among the estimated ages from different models is rather large., The high luminosities of A and B clearly indicate their young age <20 Myr); however the scatter among the estimated ages from different models is rather large.
+ The three theoretical models plotted in Figure 3. also predict Li depletion patterns for PMS stars., The three theoretical models plotted in Figure \ref{HRD} also predict Li depletion patterns for PMS stars.
+ Since we know how much Li currently remains in HIP 112312 A and B components contents to the predicted values (Table 3)). Zucker," Since we know how much Li currently remains in HIP 112312 A and B components (section \ref{LiContent}) ), we can directly compare the observed Li contents to the predicted values (Table \ref{LiTable}) )."
+manefaf(2001). identified 30 co-moving stars near Earth. the Pictoris moving group.," \citet{bPic2} identified 30 co-moving stars near Earth, the Pictoris moving group."
+ By plotting late- members on a color magnitude diagram with theoretical isochrones. they estimated an age of the group as 12 Myr.," By plotting late-type members on a color magnitude diagram with theoretical isochrones, they estimated an age of the group as 12 Myr."
+Spectrograph (MagE; ?)) on the 6.5m Clay Telescope at Las Campanas Observatory on 2008 November 26 and 2010 March 9 (see Figure 2)).,Spectrograph (MagE; \citealt{2008SPIE.7014E.169M}) ) on the 6.5m Clay Telescope at Las Campanas Observatory on 2008 November 26 and 2010 March 9 (see Figure \ref{fig:NLTT20346_spectra}) ).
+" MagE is a cross-dispersed optical spectrograph, covering 3,000 to 10,000 at medium resolution (R~ 4,100)."," MagE is a cross–dispersed optical spectrograph, covering 3,000 to 10,000 at medium resolution $R \sim 4,100$ )."
+ Our observations employed a 0.7” slit aligned at the parallactic angle., Our observations employed a $0.7^{\prime\prime}$ slit aligned at the parallactic angle.
+" Observations were made under clear conditions with an average seeing of —0.7"".", Observations were made under clear conditions with an average seeing of $\sim$ $\arcsec$.
+" The separation between the primary and secondary is —2"" therefore minimal contamination (« 1%)) was expected.", The separation between the primary and secondary is $\sim$ $\arcsec$ therefore minimal contamination $<$ ) was expected.
+ A 300s integration was obtained for each component followed by a ThAr lamp spectrum for wavelength calibration., A 300s integration was obtained for each component followed by a ThAr lamp spectrum for wavelength calibration.
+ The spectrophotometric standard GD 108 or EG 274 was observed for flux calibration., The spectrophotometric standard GD 108 or EG 274 was observed for flux calibration.
+ Ten Xe-flash and Quartz lamp flats as well as twilight flats were taken at the start of the evening for pixel response calibration., Ten Xe-flash and Quartz lamp flats as well as twilight flats were taken at the start of the evening for pixel response calibration.
+" The data were reduced using the MagE Spectral Extractor pipeline (MASE; ?)) which incorporates flat fielding, sky subtraction and flux calibration IDL routines."," The data were reduced using the MagE Spectral Extractor pipeline (MASE; \citealt{2009PASP..121.1409B}) ) which incorporates flat fielding, sky subtraction and flux calibration IDL routines."
+ We obtained low-resolution near-infrared (NIR) spectroscopy for both the primary and secondary components of NLTT 20346 using the SpeX spectrograph (?)) mounted on the 3m NASA Infrared Telescope Facility (IRTF) on 2009 December 3 UT (see Figure 3)) ., We obtained low-resolution near-infrared (NIR) spectroscopy for both the primary and secondary components of NLTT 20346 using the SpeX spectrograph \citealt{2003PASP..115..362R}) ) mounted on the 3m NASA Infrared Telescope Facility (IRTF) on 2009 December 3 UT (see Figure \ref{fig:near-IR}) ).
+ We used the spectrograph in prism mode with the 005 slit aligned to the parallactic angle., We used the spectrograph in prism mode with the $\farcs$ 5 slit aligned to the parallactic angle.
+ This resulted in R= A/ AA®& 120 spectral data over the wavelength range in 0.7-2.5 jm. Conditions were clear and the seeing was 0:46 at K., This resulted in $R~\equiv~\lambda$ / $\Delta\lambda~\approx$ 120 spectral data over the wavelength range in 0.7–2.5 $\mu$ m. Conditions were clear and the seeing was $\farcs$ 46 at $K$.
+ The stars were separated sufficiently that we were able to obtain 12 individual exposure times of 30 seconds (total exposure time of 6 minutes) in an ABBA dither pattern along the slit., The stars were separated sufficiently that we were able to obtain 12 individual exposure times of 30 seconds (total exposure time of 6 minutes) in an ABBA dither pattern along the slit.
+ Immediately after the science observations we observed the AOV star HD 43607 at a similar airmass for telluric corrections and flux calibration., Immediately after the science observations we observed the A0V star HD 43607 at a similar airmass for telluric corrections and flux calibration.
+" Internal flat-field and Ar arc lamp exposures were acquired for pixel response and wavelength calibration, respectively."," Internal flat-field and Ar arc lamp exposures were acquired for pixel response and wavelength calibration, respectively."
+" All data were reduced using the SpeXtool package version 3.4 (?,, ?)) using standard settings."," All data were reduced using the SpeXtool package version 3.4 \citealt{2004PASP..116..362C}, \citealt{2003PASP..115..389V}) ) using standard settings."
+" In addition to NIR spectroscopy, we used the SpeX guider camera to image NLTT 20346 and resolve the two components with a S/N of > 50 for each."," In addition to NIR spectroscopy, we used the SpeX guider camera to image NLTT 20346 and resolve the two components with a S/N of $>$ 50 for each."
+" The resolution is 0712 per pixel with a 60"" x 60"" field of view.", The resolution is $\farcs$ 12 per pixel with a $\arcsec$ $\times$ $\arcsec$ field of view.
+ We, We
+into the midplane to ensure adequate vertical resolution (AG=0.3° at the midplane).,into the midplane to ensure adequate vertical resolution $\Delta \theta = 0.3^o$ at the midplane).
+ The o erid is unilormlv spaced., The $\phi$ grid is uniformly spaced.
+" The number of terms in the spherical harmonic expansion Lor the eravitational potential of the disk is Nyy),=32 when δις256. while N44,=48 when Noe=512."," The number of terms in the spherical harmonic expansion for the gravitational potential of the disk is $N_{Ylm} = 32$ when $N_\phi = 256$, while $N_{Ylm} = 48$ when $N_\phi = 512$."
+ The Jeans length criterion (e.g.. Boss et al.," The Jeans length criterion (e.g., Boss et al."
+ 2000) and the Toomre length. criterion (Nelson 2006) are both monitored throughout the evolutions to ensure that any cblumps that might form are not numerical artifacts., 2000) and the Toomre length criterion (Nelson 2006) are both monitored throughout the evolutions to ensure that any clumps that might form are not numerical artifacts.
+" The Jeans length criterion consists of requiring (hat all of the grid spacings in the spherical coordinate grid remain smaller (han 1/4 of the Jeans length À,— Cp where c, is the local sound speed. G the gravitational constant. and p the densitv."," The Jeans length criterion consists of requiring that all of the grid spacings in the spherical coordinate grid remain smaller than 1/4 of the Jeans length $\lambda_J =
+\sqrt{ \pi c_s^2 \over G \rho}$ , where $c_s$ is the local sound speed, $G$ the gravitational constant, and $\rho$ the density."
+ Similarly. the Toomre length criterion consists of requiring that all of the grid spacings remain smaller than 1/4 of the Toomre length Ay=(22/6). where X is the mass surface density.," Similarly, the Toomre length criterion consists of requiring that all of the grid spacings remain smaller than 1/4 of the Toomre length $\lambda_T = (2 c_s^2 / G \Sigma)$, where $\Sigma$ is the mass surface density."
+ Once well«defimed Iragments form. these criteria may be violated at the maximum densities of the chumps. due to the non-adaptive nature of the spherical coordinate eyid. as is expected (o be (he case for sell-gravitating clumps that are trving to contract (o hieher densities on a fixed grid.," Once well-defined fragments form, these criteria may be violated at the maximum densities of the clumps, due to the non-adaptive nature of the spherical coordinate grid, as is expected to be the case for self-gravitating clumps that are trying to contract to higher densities on a fixed grid."
+ ILowever. provided that the Jeans and Toomre constraints are satislied at the (time that well-defined clumps appear. these clumps are expected to be eenuine and not spurious artifacts.," However, provided that the Jeans and Toomre constraints are satisfied at the time that well-defined clumps appear, these clumps are expected to be genuine and not spurious artifacts."
+ The boundary conditions are chosen at both 20 ancl GO AU to absorb radial velocity perturbations. to simulate the continued existence of the disk inside and outside the active numerical grid.," The boundary conditions are chosen at both 20 and 60 AU to absorb radial velocity perturbations, to simulate the continued existence of the disk inside and outside the active numerical grid."
+ As discussed in detail by Boss (1998). the use of such non-reflective boundary conditions should err on the side of caution regarding the growth of perturbations. as found bv Adams. Ruden. Shu (1989).," As discussed in detail by Boss (1998), the use of such non-reflective boundary conditions should err on the side of caution regarding the growth of perturbations, as found by Adams, Ruden, Shu (1989)."
+ Mass and momentum that enters the innermost shell of cells at 20 AU are added to the central protostar. whereas mass or momentum that reaches ihe outermost shell of cells at GO AU remains on the active hydrocdvnamical grid.," Mass and momentum that enters the innermost shell of cells at 20 AU are added to the central protostar, whereas mass or momentum that reaches the outermost shell of cells at 60 AU remains on the active hydrodynamical grid."
+ The controversy over whether or not disk instability can lead to protoplanet formation inside about 20 AU continues unabated (see the recent reviews by Durisen et al., The controversy over whether or not disk instability can lead to protoplanet formation inside about 20 AU continues unabated (see the recent reviews by Durisen et al.
+ 2007 and Mayer. Boss. Nelson 2010).," 2007 and Mayer, Boss, Nelson 2010)."
+ Attempts to find a single reason lor different numerical outcomes for disk instability models have been largely unsuccesslul to date (e.g.. Boss 2007. 2008). implying that the reason cannot be (traced to a single code difference. but rather to the totalitv of differences. such as spatial resolution. gravitational potential accuracy. artificial viscosity. stellar irradiation effects. radiative transler. numerical heating. equations of state. inilial density and temperature profiles. disk surlace boundary condiüons. and time step size. lo name a few.," Attempts to find a single reason for different numerical outcomes for disk instability models have been largely unsuccessful to date (e.g., Boss 2007, 2008), implying that the reason cannot be traced to a single code difference, but rather to the totality of differences, such as spatial resolution, gravitational potential accuracy, artificial viscosity, stellar irradiation effects, radiative transfer, numerical heating, equations of state, initial density and temperature profiles, disk surface boundary conditions, and time step size, to name a few."
+ Comparison calculations on nearly identical disk models have led Boss (2007) ancl Cai et al. (, Comparison calculations on nearly identical disk models have led Boss (2007) and Cai et al. (
+2010) to reach different conclusions.,2010) to reach different conclusions.
+ While Boss (2007) concluded that fragmentation was possible inside 20 AU. Cai et al. (," While Boss (2007) concluded that fragmentation was possible inside 20 AU, Cai et al. ("
+2010) found no evidence for fragmentationin their models.,"2010) found no evidence for fragmentationin their models,"
+or be accreted onto a central black hole.,or be accreted onto a central black hole.
+ A high μμ/Lgcw ratio is still consistent with the measured upper limit. and would indicate that a significant contribution is made to {η from dust heated by the central AGN. rather than by star formation.," A high $L_{\rm FIR}/L_{\rm
+HCN}$ ratio is still consistent with the measured upper limit, and would indicate that a significant contribution is made to $L_{\rm FIR}$ from dust heated by the central AGN, rather than by star formation."
+ A global measure of the HCN emission from a galaxy in which both AGN and starburst activity are observed cannot definitively establish the dominant source of the far-infrared luminosity. as strong HCN emission is seen in both starbursts and local AGN (Seyfert galaxies: Curran et 22001).," A global measure of the HCN emission from a galaxy in which both AGN and starburst activity are observed cannot definitively establish the dominant source of the far-infrared luminosity, as strong HCN emission is seen in both starbursts and local AGN (Seyfert galaxies: Curran et 2001)."
+ Differentiating between the two may be possible with images of both considerably higher angular resolution and of higher HCN transitions., Differentiating between the two may be possible with images of both considerably higher angular resolution and of higher HCN transitions.
+ In this way it should be possible to trace the distribution of the denser gas as well as its excitation., In this way it should be possible to trace the distribution of the denser gas as well as its excitation.
+ KGI acknowledges support of a PPARC fellowship and the Cavendish Astrophysies Group., KGI acknowledges support of a PPARC fellowship and the Cavendish Astrophysics Group.
+ We thank the anonymous referee for comments that helped to clarify points made in this paper., We thank the anonymous referee for comments that helped to clarify points made in this paper.
+ This research has made use of NASA's Astrophysics Data System., This research has made use of NASA's Astrophysics Data System.
+Relevant details of illustrating5 examples1 for three cases of jo-— 0.24. 3=l/s and —lf4 respectively are shown in three panels(a). (b) and (ο) of Fig.,"Relevant details of illustrating examples for three cases of $\beta=-0.24$ , $\beta=-1/8$ and $\beta=1/4$ respectively are shown in three panels (a), (b) and (c) of Fig."
+ 2 and are also summarizedin Table 2, 2 and are also summarized in Table 2.
+" Whenmc 2. we have Ca 0.0κCP,« 1. Ay,>0. By,m0. Hy,c0 inthe open 2 interval 2C(.1/4.17/2)."," When$m\ge 2$ , we have $C_2>0$, $0<{\cal C}{\cal P}_m<1$ , ${\cal A}_m >0$, ${\cal B}_m>0$, ${\cal H}_m>0$ inthe open $\beta$ interval $\beta\in(-1/4,1/2)$."
+ 'herefore in this case of mc2. jj and. yo remain always upper and lower branches. respectively.," Therefore in this case of $m\ge 2$, $y_1$ and $y_2$ remain always upper and lower branches, respectively."
+" By solutions (40)) for y= 1. we have while by solutions (41)) for ap5ox. we have As the limiting situations for η=1 and ἡx well bracket possible ranges of jj, and js branches (see Appenclix 13). it is obvious that the upper yr branch remains always positive and the value of yy increases with increasing ὁ and decreases with increasing either m ory."," By solutions \ref{alisoleta1}) ) for $\eta=1$ , we have while by solutions \ref{alisoletainf}) ) for $\eta\rightarrow\infty$, we have As the limiting situations for $\eta=1$ and $\eta\rightarrow\infty$ well bracket possible ranges of $y_1$ and $y_2$ branches (see Appendix B), it is obvious that the upper $y_1$ branch remains always positive and the value of $y_1$ increases with increasing $\delta$ and decreases with increasing either $m$ or $\eta$."
+ Meanwhile. the lower yo branch has a specific critical value αν of η bevond which the yo solution becomes unphysical for being negative.," Meanwhile, the lower $y_2$ branch has a specific critical value $\eta_c$ of $\eta$ beyond which the $y_2$ solution becomes unphysical for being negative."
+ This critical value g. is given by where we have δέ always greater than m for all m>2 within 3Cή.1/2).," This critical value $\eta_c$ is given by where we have ${\cal H}_m$ always greater than $m$ for all $m\ge 2$ within $\beta\in(-1/4,1/2)$."
+" Such a critical 5, always exists for the yo branch for all values of 9.", Such a critical $\eta_c$ always exists for the $y_2$ branch for all values of $\delta$.
+ In other words. there does not exist a critical value ὃς for 0 as in the aligned m=1 Case.," In other words, there does not exist a critical value $\delta_c$ for $\delta$ as in the aligned $m=1$ case."
+ This critical value g of 5 decreases with increasing 9. much like the case of two coupled SLDs investigated recently (Lou Shen 2003).," This critical value $\eta_c$ of $\eta$ decreases with increasing $\delta$, much like the case of two coupled SIDs investigated recently (Lou Shen 2003)."
+ For physical regimes of gy; and yp. they both decrease monotonically with increasing η. while yy branch increases monotonically and yo branch decreases monotonically with increasing 9.," For physical regimes of $y_1$ and $y_2$, they both decrease monotonically with increasing $\eta$, while $y_1$ branch increases monotonically and $y_2$ branch decreases monotonically with increasing $\delta$."
+ For the phase relationship between fe and f£. it is straightforward. to show that the ratio ἳμη decreases monotonically with increasing 27and thus increases monotonically with increasing 5g for m22.," For the phase relationship between $\mu^g$ and $\mu^s$, it is straightforward to show that the ratio $\mu^g/\mu^s$ decreases monotonically with increasing $D_s^2$and thus increases monotonically with increasing $\eta$ for $m\ge 2$."
+ For the upper yy branch. we obtain where the left-hand bound corresponds toa= Land the right-hand. bound. correspondsto 7 x.," For the upper $y_1$ branch, we obtain where the left-hand bound corresponds to $\eta=1$ and the right-hand bound correspondsto $\eta\rightarrow\infty$ ."
+ Inthe specified range of3C( 4.1/2). the ratio μμ remains alwaysnegative for the upper yr branch. indicating surface mass density perturbations in two clises are out-of-phase.," Inthe specified range of$\beta\in(-1/4,1/2)$ , the ratio $\mu^g/\mu^s$ remains alwaysnegative for the upper $y_1$ branch, indicating surface mass density perturbations in two discs are out-of-phase."
+ For the lower jo branch in parallel. we derive ∖∖⋎↓↥∢⊾↓⋅⋖⋅⇂↓↕⋖⋅↓⋖⊾∐−↓⋯⊔∠⇂∣⋯⊔⊔∠⇂≼∙∪↓⋅↓⋅∢⋅≱∖↓≻∪⊔∠⇂⊳∖⋯↗∣∶↓⋜⋯∠⇂ ↿↓↥∢⊾↓⋰↓⋏∙≟↓∐−↓⋯⊔∠⇂∣⋡∪⊔⊔∠⇂≼∙∪↓⋅↓⋅∢⋅≱∖↓≻∪⊔∠⇂≱∖↥∢≱↿⇂↥∢⋅≼∼↓⋅↕∣↕≼∼⋜↧↓↗∣∡∖∖⋎↓↕⊲⊔∼↓↥ ⊔↓⋜↧↳∢⊾⊳∖∐⋮∶∕∣⊐," For the lower $y_2$ branch in parallel, we derive where the left-hand bound corresponds to $\eta=1$ and the right-hand bound corresponds to the critical $\eta_c$ which makes $D_s^2=y_2=0$ ."
+∶∪⇀∖↓≻↓≻⋜⊔⋅⋖⋅⊔↿↓∙∖⇁⊳⇂↓↥∢⊾↓∪∖∖⊽∢⊾↓⋅∕∣⊐∣⋡," Apparently, the lower $y_2$ branch (if"
+∶∪⇀∖↓≻↓≻⋜⊔⋅⋖⋅⊔↿↓∙∖⇁⊳⇂↓↥∢⊾↓∪∖∖⊽∢⊾↓⋅∕∣⊐∣⋡↓," Apparently, the lower $y_2$ branch (if"
+∶∪⇀∖↓≻↓≻⋜⊔⋅⋖⋅⊔↿↓∙∖⇁⊳⇂↓↥∢⊾↓∪∖∖⊽∢⊾↓⋅∕∣⊐∣⋡↓⋅," Apparently, the lower $y_2$ branch (if"
+∶∪⇀∖↓≻↓≻⋜⊔⋅⋖⋅⊔↿↓∙∖⇁⊳⇂↓↥∢⊾↓∪∖∖⊽∢⊾↓⋅∕∣⊐∣⋡↓⋅⋜," Apparently, the lower $y_2$ branch (if"
+∶∪⇀∖↓≻↓≻⋜⊔⋅⋖⋅⊔↿↓∙∖⇁⊳⇂↓↥∢⊾↓∪∖∖⊽∢⊾↓⋅∕∣⊐∣⋡↓⋅⋜⋯," Apparently, the lower $y_2$ branch (if"
+∶∪⇀∖↓≻↓≻⋜⊔⋅⋖⋅⊔↿↓∙∖⇁⊳⇂↓↥∢⊾↓∪∖∖⊽∢⊾↓⋅∕∣⊐∣⋡↓⋅⋜⋯≼," Apparently, the lower $y_2$ branch (if"
+∶∪⇀∖↓≻↓≻⋜⊔⋅⋖⋅⊔↿↓∙∖⇁⊳⇂↓↥∢⊾↓∪∖∖⊽∢⊾↓⋅∕∣⊐∣⋡↓⋅⋜⋯≼⇍," Apparently, the lower $y_2$ branch (if"
+∶∪⇀∖↓≻↓≻⋜⊔⋅⋖⋅⊔↿↓∙∖⇁⊳⇂↓↥∢⊾↓∪∖∖⊽∢⊾↓⋅∕∣⊐∣⋡↓⋅⋜⋯≼⇍↓," Apparently, the lower $y_2$ branch (if"
+∶∪⇀∖↓≻↓≻⋜⊔⋅⋖⋅⊔↿↓∙∖⇁⊳⇂↓↥∢⊾↓∪∖∖⊽∢⊾↓⋅∕∣⊐∣⋡↓⋅⋜⋯≼⇍↓↕," Apparently, the lower $y_2$ branch (if"
+∶∪⇀∖↓≻↓≻⋜⊔⋅⋖⋅⊔↿↓∙∖⇁⊳⇂↓↥∢⊾↓∪∖∖⊽∢⊾↓⋅∕∣⊐∣⋡↓⋅⋜⋯≼⇍↓↕↿," Apparently, the lower $y_2$ branch (if"
+∶∪⇀∖↓≻↓≻⋜⊔⋅⋖⋅⊔↿↓∙∖⇁⊳⇂↓↥∢⊾↓∪∖∖⊽∢⊾↓⋅∕∣⊐∣⋡↓⋅⋜⋯≼⇍↓↕↿∖," Apparently, the lower $y_2$ branch (if"
+∶∪⇀∖↓≻↓≻⋜⊔⋅⋖⋅⊔↿↓∙∖⇁⊳⇂↓↥∢⊾↓∪∖∖⊽∢⊾↓⋅∕∣⊐∣⋡↓⋅⋜⋯≼⇍↓↕↿∖∐," Apparently, the lower $y_2$ branch (if"
+∶∪⇀∖↓≻↓≻⋜⊔⋅⋖⋅⊔↿↓∙∖⇁⊳⇂↓↥∢⊾↓∪∖∖⊽∢⊾↓⋅∕∣⊐∣⋡↓⋅⋜⋯≼⇍↓↕↿∖∐⋅," Apparently, the lower $y_2$ branch (if"
+is a FRI radio source at redshift 0.109. (Spinrad 11985).,is a FRI radio source at redshift 0.109 (Spinrad 1985).
+ Its 178-MlIz luminosity is 7.6.1077 NC tse t. slightly above the nominal ERI-EBRAE boundary of ~2.1077 Wo tse + (Fanarolf Riley 1974. hereafter EIU).," Its 178-MHz luminosity is $7.6 \times 10^{25}$ W $^{-1}$ $^{-1}$, slightly above the nominal FRI-FRII boundary of $\sim 2
+\times 10^{25}$ W $^{-1}$ $^{-1}$ (Fanaroff Riley 1974, hereafter FR)."
+ Leahy (1985. 1993) and Jàgeers and de αρ (1985) present intermediate-resolution. VLA maps of the central. regions of the source. while Jageers (1983) has a lower-resolution WSR image which shows the whole source and its field: the source extends for ~1.5 Alpe.," Leahy (1985, 1993) and Jäggers and de Grijp (1985) present intermediate-resolution VLA maps of the central regions of the source, while Jäggers (1983) has a lower-resolution WSRT image which shows the whole source and its field; the source extends for $\sim 1.5$ Mpc."
+ Saripalli ((1996) present high-frequency maps made with the Eliclsberg 100-m. telescope., Saripalli (1996) present high-frequency maps made with the Effelsberg 100-m telescope.
+ The host galaxy ds. classed as a DEP by Wenelham (1966) and appears to lie in a cluster. although strong galactic redeenine makes optical identification of the cluster members dillieult.," The host galaxy is classed as a DE2 by Wyndham (1966) and appears to lie in a cluster, although strong galactic reddening makes optical identification of the cluster members difficult."
+ Pho detection of extended: X-ray. emission (Miley. 11983). the nearbby aligneed sources (Jageers 1983). and he many mv racio sources in the field at 1.5 Cllz make it plausible that the object is the dominant member of a arge cluster.," The detection of extended X-ray emission (Miley 1983), the by ed sources (Jäggers 1983) and the many mJy radio sources in the field at 1.5 GHz make it plausible that the object is the dominant member of a large cluster."
+ Leahy (1985) also attempts to constrain the VM distribution of the source. but notes that it depolarizes rapidly (particularly in. the S lobe) so that lew good measurements are available: this could be taken as evidence or à dense magneto-ionic environment for the source LIvedra A. Taylor 11990.," Leahy (1985) also attempts to constrain the RM distribution of the source, but notes that it depolarizes rapidly (particularly in the S lobe) so that few good measurements are available; this could be taken as evidence for a dense magneto-ionic environment for the source Hydra A, Taylor 1990)."
+ L130 is a wide-angle tail (WAT) radio source., 130 is a wide-angle tail (WAT) radio source.
+ The term WAT has been used to describe many different tvpes of object., The term WAT has been used to describe many different types of object.
+ Llere E shall use it to refer to those FRI sources which are associated with central cluster galaxies OOwen Itudnick 1976) and have luminosities comparable o or exceeding the Fanaroll-Riley break between FRI and FRILL, Here I shall use it to refer to those FRI sources which are associated with central cluster galaxies Owen Rudnick 1976) and have luminosities comparable to or exceeding the Fanaroff-Riley break between FRI and FRII.
+ E shall follow Leahy (1993) in using the behaviour of he jets at the base as another defining feature., I shall follow Leahy (1993) in using the behaviour of the jets at the base as another defining feature.
+ At high resolution one or two welbeollimated jets strong-llavour jets. by the classification of Leahy (1993)] are seen OODonoghue. Owen Eilek 1990). extending for some tens of kpe before broadening. often at a bright Dare point. into he characteristic plumes or tails.," At high resolution one or two well-collimated jets [`strong-flavour' jets, by the classification of Leahy (1993)] are seen O'Donoghue, Owen Eilek 1990), extending for some tens of kpc before broadening, often at a bright flare point, into the characteristic plumes or tails."
+" These jets are very similar o the jets seen in PRI radio galaxies. and. quite cilferent rom the behaviour of jets in more tvpical PRIs. where a collimated inner jet. if visible at all. decollimates rapidly (on scales of a few kpe at most) and. comparatively smoothly into a bright ""weak-Iavour jet with a large opening WATS. according to this definition. never have a weak-lavour jet. but make the transition between strong-Iavour jet and. diffuse. bent tail in a single step."," These jets are very similar to the jets seen in FRII radio galaxies, and quite different from the behaviour of jets in more typical FRIs, where a collimated inner jet, if visible at all, decollimates rapidly (on scales of a few kpc at most) and comparatively smoothly into a bright `weak-flavour' jet with a large opening WATs, according to this definition, never have a weak-flavour jet, but make the transition between strong-flavour jet and diffuse, bent tail in a single step."
+ Phe requirement, The requirement
+BSs contribute little to the core population.,BSs contribute little to the core population.
+ A scenario in which all BSs are generated over the entire cluster by mass-transfer in PBs can not be ruled out., A scenario in which all BSs are generated over the entire cluster by mass-transfer in PBs can not be ruled out.
+ We believe however that a blending between collision induced evolution and internal evolution is at play in the GC core to explain internal. BSs., We believe however that a blending between collision induced evolution and internal evolution is at play in the GC core to explain internal BSs.
+ Our main finding is the need of a population of external PBs in order to generate the bimodal distribution of the BSs observed in 47 Tuc., Our main finding is the need of a population of external PBs in order to generate the bimodal distribution of the BSs observed in 47 Tuc.
+ The required fraction (10%)) to fit the data does not seem unreasonable., The required fraction ) to fit the data does not seem unreasonable.
+ More accurate counts of BSs in GCs with widely different properties are about to be collected from high resolution photometry (Ferraro et al., More accurate counts of BSs in GCs with widely different properties are about to be collected from high resolution photometry (Ferraro et al.
+ 2004 in preparation)., 2004 in preparation).
+ These observations may shed light into the nature of BSs and the importance of PBs in GC evolution., These observations may shed light into the nature of BSs and the importance of PBs in GC evolution.
+ In light of these upcoming observations. in paper II (Mapelli et al.," In light of these upcoming observations, in paper II (Mapelli et al."
+ in preparation). we plan to continue our analysis with our simulations over a wide sample of GCs.," in preparation), we plan to continue our analysis with our simulations over a wide sample of GCs."
+ Theoretical studies using N-body (Baumgardt Makino 2003) or Monte Carlo techniques (Ivanova et al., Theoretical studies using N-body (Baumgardt Makino 2003) or Monte Carlo techniques (Ivanova et al.
+ 2004) will eventually become necessary tool for exploring the formation and evolution of BSs in GCs., 2004) will eventually become necessary tool for exploring the formation and evolution of BSs in GCs.
+ The financial support of the Agenzia Spaziale Italiana and of the is kindly acknowledged., The financial support of the Agenzia Spaziale Italiana and of the is kindly acknowledged.
+ RTR is partially supported by STSel grant GO-8709 and NASA LTSA grant NAG 5-6403., RTR is partially supported by STScI grant GO-8709 and NASA LTSA grant NAG 5-6403.
+Clobulu clusters contain many X-raw sources at lower lmminositics. L4x107!cres.,"Globular clusters contain many X-ray sources at lower luminosities, $L_{\rm x}\ltap 10^{34}$."
+. These sources were first discovered with the Eiusteiu satellite (Παν CCaindlay 1983). and many more were found with ROSAT (for a compilation. see Joluston VVerbuut 1996).," These sources were first discovered with the Einstein satellite (Hertz Grindlay 1983), and many more were found with ROSAT (for a compilation, see Johnston Verbunt 1996)."
+ The nature of these low-luuinosity sources is the subject of debate. because various types of objects can enit N-ravs at such luminosities. such as soft X-ray trausicuts in quiescence. «αταπο variables. RS CVn binaries. and recveled neutron stars (see FFig.," The nature of these low-luminosity sources is the subject of debate, because various types of objects can emit X-rays at such luminosities, such as soft X-ray transients in quiescence, cataclysmic variables, RS CVn binaries, and recycled neutron stars (see Fig."
+ δ iu Verbuut et 11997)., 8 in Verbunt et 1997).
+ The most compcling identification of a dim N-vav source with an object observed at other wavelengths is the recveled radio pulsar in M22s8: the X-ray flux varies on the pulse period (Dauner et H1991)., The most compelling identification of a dim X-ray source with an object observed at other wavelengths is the recycled radio pulsar in 28: the X-ray flux varies on the pulse period (Danner et 1994).
+ Plausible ideutifications with cataclysmic variables have been sugeestedl for din N-rav sources in 66307. 66752. 55901 and. TTuc (Cool et 11995b. Caindlay 1995. Hala et 11997. Verbunt UWasinecr 1998).," Plausible identifications with cataclysmic variables have been suggested for dim X-ray sources in 6397, 6752, 5904 and Tuc (Cool et 1995b, Grindlay 1993, Hakala et 1997, Verbunt Hasinger 1998)."
+ These ideutificatious are based ou the xoxiuitv of the N-rav. position to that of a cataclysiic variable. ando thus heir probability depends ou the accuracy of the X-ray. position.," These identifications are based on the proximity of the X-ray position to that of a cataclysmic variable, and thus their probability depends on the accuracy of the X-ray position."
+ The accuracy of the ROSAT position of a detected N-ray source is determined by two factors: the statistical accuracy ofthe position of the source on the detector. aud he accuracy with which the position of the detector as a whole is xojected on the sky.," The accuracy of the ROSAT position of a detected X-ray source is determined by two factors: the statistical accuracy of the position of the source on the detector, and the accuracy with which the position of the detector as a whole is projected on the sky."
+" For a sufficient ΠΟΥ of photons the statistical error is less than an arcsecond. but he projection iu 5general las a typical error of ~5""."," For a sufficient number of photons the statistical error is less than an arcsecond, but the projection in general has a typical error of $\sim5''$."
+ Secure ↕≼∐∖↕↕↕∐↸⊳⋜↧↕∪↕↕∪↕≯⋜↧↴∖↴∪⋃⋅∩∖↕∐↑∐↸∖≼∐∖↑↸∖↸⊳∪↥⋅∱∐∖↕≼↧↥⋅↸∖≼⊔∐⊳↸∖↴∖↴ he error iu the projection to the statistical eror of the ideutified source. provided that the optical (or radio) )osition its better accuracy.," Secure identification of a source in the detector field reduces the error in the projection to the statistical error of the identified source, provided that the optical (or radio) position has better accuracy."
+ Oulv one identification is recessary. because the roll augle of he detector tthe North-South direction) is accurately known: nonetheless. identification of more than one source is preferable to allow checks ou internal cousisteucy.," Only one identification is necessary, because the roll angle of the detector the North-South direction) is accurately known; nonetheless, identification of more than one source is preferable to allow checks on internal consistency."
+ In a elobular cluster the surface density of possible counterparts is so high that chance coincidence usually. caunot be excluded: a secure identification can usually be mace ouly for N-ray. sources detected well outside the cluster., In a globular cluster the surface density of possible counterparts is so high that chance coincidence usually cannot be excluded; a secure identification can usually be made only for X-ray sources detected well outside the cluster.
+" This nethod las been used by Verbuut UWasinecr (1998) to nuprove the positional accuracy of the sources in the core of TTue from 5"" to 2"", whereby the area in which the source is expected to lie is reduced suthicicutly to exclude several proposed identifications. aud increase the probability of others. iucluding two possible cataclysmic variables."," This method has been used by Verbunt Hasinger (1998) to improve the positional accuracy of the sources in the core of Tuc from $5''$ to $2''$, whereby the area in which the source is expected to lie is reduced sufficiently to exclude several proposed identifications, and increase the probability of others, including two possible cataclysmic variables."
+ Tn this paper we investigate three clusters kuown to contain multiple din N-vav sources iu their core. which," In this paper we investigate three clusters known to contain multiple dim X-ray sources in their core, which"
+Near the eut-ofÉ racius. which lies close to #2=1. the mode properties are similar to those seen in Fig. 1..,"Near the cut-off radius, which lies close to $R=1$, the mode properties are similar to those seen in Fig. \ref{circvsomega},"
+ except that the refractive index. now plotted as as a function of radius. is no longer degenerate.," except that the refractive index, now plotted as as a function of radius, is no longer degenerate."
+ The detailed properties of these solutions have been discussed in ?.., The detailed properties of these solutions have been discussed in \cite{kirk10}.
+ llowever. note that these curves do not describe the radial evolution of à wave packet. as erroneously suggested in (hat paper and assumed by ? (who considered linear polarization). but simply specilv the wave modes into which a wind of given po and σ can convert al a given radius.," However, note that these curves do not describe the radial evolution of a wave packet, as erroneously suggested in that paper and assumed by \citet{asseollobetpellat84}
+ (who considered linear polarization), but simply specify the wave modes into which a wind of given $\mu$ and $\sigma$ can convert at a given radius."
+ The difference arises because the racial momentum flix density. 7. is not a conserved quantity in the radial evolution of the superluminal modes (seealso?)..," The difference arises because the radial momentum flux density, $\nu$, is not a conserved quantity in the radial evolution of the superluminal modes \citep[see also][]{kirkmochol11a}."
+ When the phase-averaged magnetic field vanishes. the properties of linearly. polarized modes are similar to those of circularly polarized modes.," When the phase-averaged magnetic field vanishes, the properties of linearly polarized modes are similar to those of circularly polarized modes."
+ This is illustrated in Fig. 3..," This is illustrated in Fig. \ref{comparisonlincirc},"
+ where one sees that the refractive indices are practically identical., where one sees that the refractive indices are practically identical.
+ This ligure also shows the nonlinearity parameter q., This figure also shows the nonlinearity parameter $q$.
+ For linear polarization. q is defined in ternis of quantities measured al phase zero.," For linear polarization, $q$ is defined in terms of quantities measured at phase zero."
+ This point in (he wave is. however. special.," This point in the wave is, however, special."
+ It corresponds to the turning points of the wwavelorm (?).. where the [Iuid velocity either vanishes or lies precisely in (he propagation direction.," It corresponds to the turning points of the waveform \citep{maxperkins71}, where the fluid velocity either vanishes or lies precisely in the propagation direction."
+ These points do not play a large role in determining (he average properties of the wave., These points do not play a large role in determining the average properties of the wave.
+ As a result. q varies substantially. especially close to points where μμ changes sign (log()2z0.4 in the ligure).," As a result, $q$ varies substantially, especially close to points where $u_{x0}$ changes sign $\log(R)\approx 0.4$ in the figure)."
+ One consequence is (hat an expansion of the fluxes using q as a small parameter is inaclequate over most of (he range relevant for pulsars. despite the [act that the waves are highly nonlinear (see the discussion in reldiscussion)).," One consequence is that an expansion of the fluxes using $q$ as a small parameter is inadequate over most of the range relevant for pulsars, despite the fact that the waves are highly nonlinear (see the discussion in \\ref{discussion}) )."
+ hi the case of circular polarization. the quantities (hal enter into the," In the case of circular polarization, the quantities that enter into the"
+and hence its SLI observation was removed from our observations list.,and hence its SL1 observation was removed from our observations list.
+ Unfortunately. the depth of the previous observation is x lower than we needed here. but (his is still considered a duplicate observation (the limit is x).," Unfortunately, the depth of the previous observation is $\times$ lower than we needed here, but this is still considered a duplicate observation (the limit is $\times$ )."
+ The depth of the earlier observation is insullicient for our purposes hence it can serve only (o constrain the continuum levels. but is not useful in looking for absorption features.," The depth of the earlier observation is insufficient for our purposes hence it can serve only to constrain the continuum levels, but is not useful in looking for absorption features."
+ Table 1. gives à summary of the observations covering threeSpitzer observing cvcles (G01.G02. and GOL).," Table \ref{table_obs} gives a summary of the observations covering three observing cycles (G01,G02, and G04)."
+ The full GOL sample and observations are discussed in detail in (2007)., The full GO1 sample and observations are discussed in detail in \citet{yan07}.
+. The full GO2 sample and observations ave discussed in Dasvraetal.(2009)., The full GO2 sample and observations are discussed in \citet{dasyra09}.
+. The GO4 observations (DIDZZ440793) are presented. here for the first time., The GO4 observations 40793) are presented here for the first time.
+ Combining all these programs. each source has twpically ~ hhrs on source of IRS SLI aud LL2 observations. alihough brighter sources have less exposure lime as can be seen in Table 1..," Combining all these programs, each source has typically $\sim$ hrs on source of IRS SL1 and LL2 observations, although brighter sources have less exposure time as can be seen in Table \ref{table_obs}."
+ The SLI observations were mostly in mapping mode in order (o improve the signal-to-noise ratio in these very faint sources., The SL1 observations were mostly in mapping mode in order to improve the signal-to-noise ratio in these very faint sources.
+ For our sources. we adopted the IRAC and/or VLA positions which are the most accurate available and sufficiently. accurate for the SL1 module whieh has the narrowest slit width ccompared with {for LL2).," For our sources, we adopted the IRAC and/or VLA positions which are the most accurate available and sufficiently accurate for the SL1 module which has the narrowest slit width compared with for LL2)."
+ We use IrsLow. a data reduction code for IRS low resolution observations custom-wrillen bv Dario Fadda (see Facda et al.," We use IrsLow, a data reduction code for IRS low resolution observations custom-written by Dario Fadda (see Fadda et al."
+ 2010 in prep., 2010 in prep.
+ for hither details)., for further details).
+ The data reduction steps include bad pixel removal. background subtraction. co-adding different positions (such as the 2 nod positions in stare mode. or the 4 positions we (vpically used in map mode). ancl finally extracting (he spectra using optimal extraction.," The data reduction steps include bad pixel removal, background subtraction, co-adding different positions (such as the 2 nod positions in stare mode, or the 4 positions we typically used in map mode), and finally extracting the spectra using optimal extraction."
+ Dad pixels are determined as outliers across the available frames., Bad pixels are determined as outliers across the available frames.
+ For the SL data for example. fainter sources have 40 frames per source. but the bright sources have only 6 frames per source.," For the SL data for example, fainter sources have 40 frames per source, but the bright sources have only 6 frames per source."
+ Therefore this bad pixel [ageing is less ellicient for our bright sources with very. few individual Irames leading (ο somewhat noisier spectra than the nominal difference in exposure time., Therefore this bad pixel flagging is less efficient for our bright sources with very few individual frames leading to somewhat noisier spectra than the nominal difference in exposure time.
+ The code is iterative in that. after the first backeround subtraction. any sources present are masked. aud a second background subtraction is done.," The code is iterative in that after the first background subtraction, any sources present are masked, and a second background subtraction is done."
+ In addition. we test the reliability of any features we see. by displaviug all points that go into the spectral extraction aud manually clipping anv obvious outliers (on top of the bad pixel flagging of the code) that might bias our results.," In addition, we test the reliability of any features we see, by displaying all points that go into the spectral extraction and manually clipping any obvious outliers (on top of the bad pixel flagging of the code) that might bias our results."
+ A few spurious spikes are removed in (his way. but (his is mostly a concern for narrow features of interest such as (he PPAII feature.," A few spurious spikes are removed in this way, but this is mostly a concern for narrow features of interest such as the PAH feature."
+ In all cases. where we believe that this feature is or might be present we have," In all cases, where we believe that this feature is or might be present we have"
+The galaxy huuinositv function (LE) is oue of the cornerstones iu our Wulerstanding of ealaxy formation aud evolution.,The galaxy luminosity function (LF) is one of the cornerstones in our understanding of galaxy formation and evolution.
+ Since the introduction of a fitting function Or its shape bv SchecLhter(1976).. the origin of the orm of the LF fimetior1i has been a powerful coustraiut ol mode building (ee.Deusouetal.2003:Sammictal. 2007)..," Since the introduction of a fitting function for its shape by \citet{1976ApJ...203..297S}, the origin of the form of the LF function has been a powerful constraint on model building \citep[e.g.,][]{2003ApJ...599...38B,2007MNRAS.377..285S}."
+ While recent work has focused somewhat ou he luminous eund. its evolution with redshift (Brownetal.2007) aud the role of dry imergers (INhochtar&Burkert2003:Naabetal 2006). the £uüut-eud xovides additional imiportanut clues on galaxy formation.," While recent work has focused somewhat on the luminous end, its evolution with redshift \citep{2007ApJ...654..858B} and the role of dry mergers \citep{2003ApJ...597L.117K,2006ApJ...636L..81N}, the faint-end provides additional important clues on galaxy formation."
+ Systematic studies of the fünt-eud slope iu the local universe reveal ciffereuces between high aud low deusity environments (Trentham1998). as well as for ealaxy saples split by 1uorplkAlogies (c.g.Alarzkeetal.1991).," Systematic studies of the faint-end slope in the local universe reveal differences between high and low density environments \citep{1998MNRAS.294..193T}, as well as for galaxy samples split by morphologies \citep[e.g.][]{1994AJ....108..437M}."
+ The nuderlwing oivsica] processes that shape the faint-end of the LF are generally associated with feedback roni superuovae that is effective. in heating eas and driving winds in shallow gravitational poteutias (Dekel&Sills 1986).., The underlying physical processes that shape the faint-end of the LF are generally associated with feedback from supernovae that is effective in heating gas and driving winds in shallow gravitational potentials \citep{1986ApJ...303...39D}.
+ Althoueh the iuplementation of supernova ecdhack in galaxy formion models has been extensively investigated (e.g.Beusoretal.2003) for the local galaxy LF. its iupact on the redshift evolution ou he faint-eud iis not been as well studied.," Although the implementation of supernova feedback in galaxy formation models has been extensively investigated \citep[e.g.][]{2003ApJ...599...38B}
+ for the local galaxy LF, its impact on the redshift evolution on the faint-end has not been as well studied."
+ The recent wealth of LF neasured to very fait magnitudes in the rest-frame B-xad (e.g.Blantoneta.2003:Wolfal.March-exuietal.2007:Ryan2007) and the rest-frame FUV (e.g.Yan&Winehorst200la:Wryeeretal.2005:Dowweusetal.2006) alows us to test models with high accuracy.," The recent wealth of LF measured to very faint magnitudes in the rest-frame B-band \citep[e.g.][]{2003ApJ...592..819B,2003A&A...408..499W,2007ApJ...656...42M,2007astro.ph..3743R} and the rest-frame FUV \citep[e.g.][]{2004ApJ...612L..93Y,2005ApJ...619L..15W,2006ApJ...653...53B}
+ allows us to test models with high accuracy."
+ The purpose of this letter is to investigate the nuderlving driving mecwunisin for the redshift evolution of the faint-cud slope., The purpose of this letter is to investigate the underlying driving mechanism for the redshift evolution of the faint-end slope.
+" Furthermore. we investigate the Hupact of supernova feedback on the rate of star-formation by varving the relevant efficiency. parameter""S."," Furthermore, we investigate the impact of supernova feedback on the rate of star-formation by varying the relevant efficiency parameters."
+ Iu the following we briefly outline our basic iodelliug approach and refer the reader for more details to Isdhochfar&Burkert(2005)... IKhochfar&Silk(2006) alc reference therein.," In the following we briefly outline our basic modelling approach and refer the reader for more details to \citet{2005MNRAS.359.1379K}, \citet{2006MNRAS.370..902K} and reference therein."
+ We ecucrate merecr trees for dark uatter halos usine aNoute-Carlo approach based onu the extended Press-Schechter formalisua (Somerville&att 1999).., We generate merger trees for dark matter halos using a Monte-Carlo approach based on the extended Press-Schechter formalism \citep{1999MNRAS.305....1S}.
+ As we aim to model the faiut-eud of the LF ο high redshifts. we weed to make sure that the maSS resolution iu our sinulatious is sufficicut.," As we aim to model the faint-end of the LF to high redshifts, we need to make sure that the mass resolution in our simulations is sufficient."
+" We ecnerate nerecr trees from dark matter mass functions between > =Oand o6. and find that resolving cach individial norecr""veov treefvoo down- to a; nass ∖resolution1 ofi My:=f=5«|p AD. and ALnuu=1 M. af:x3 and -= l. respectively. gives robust results."," We generate merger trees from dark matter mass functions between $z=0$ and $z=6$, and find that resolving each individual merger tree down to a mass resolution of $M_{\mbox{min}}=5 \times 10^{9}$ $_{\odot}$ and $M_{\mbox{min}}= 10^{8}$ $_{\odot}$ at $z \leq 3$ and $z \geq 4$ , respectively, gives robust results."
+" Once. a tree reaches Minin. We start moving the tree forward in time iucludi""m physical processes associated with the barvous with11 cach dark matter halo that include eas cooling. star-ormation. supernova feedback. reionization aud mere18o of galaxies on a dynamical friction time-scale."," Once a tree reaches $M_{min}$, we start moving the tree forward in time including physical processes associated with the baryons within each dark matter halo that include gas cooling, star-formation, supernova feedback, reionization and merging of galaxies on a dynamical friction time-scale."
+ As 1C focus of this letter is ou the faiut-cud of the Inunositv Muction. we will omit mcludiug xesceriptious for ACN-eedback (c.g.Boweretal.2006) or cnviromuental effects (Ikhochfar&Ostriker2007) that maiuby iuthence the xieht-eud of the Iuninositv function.," As the focus of this letter is on the faint-end of the luminosity function, we will omit including prescriptions for AGN-feedback \citep[e.g.][]{2006MNRAS.370..645B}
+ or environmental effects \citep{env} that mainly influence the bright-end of the luminosity function."
+ The largest impact on the slope at the füut-eud couCR roni star-formiation and associated supernova feedback (Dekel&Silk 1986)., The largest impact on the slope at the faint-end comes from star-formation and associated supernova feedback \citep{1986ApJ...303...39D}. .
+ Faint galaxies ecnerally occu]Dy sanall dark iatter halos. with shallow poteutial wells. hatallow effective reheating of cold eas in the ISM 1ἂν coedback frou supernovac.," Faint galaxies generally occupy small dark matter halos with shallow potential wells, thatallow effective reheating of cold gas in the ISM by feedback from supernovae."
+ We model star-formatiou iu, We model star-formation in
+coupled with the elfects of relativistic motion can produce a wide range of arm-leneth and. brightness. asvniumetrics (e.g. Jevakumar et al.,coupled with the effects of relativistic motion can produce a wide range of arm-length and brightness asymmetries (e.g. Jeyakumar et al.
+ 2005)., 2005).
+ For a lobe approaching us and propagating through a dense medium. the oppositely directed. lobes could be somewhat svimmetrically located with [large asymmetries in brightness or Lux density ratios.," For a lobe approaching us and propagating through a dense medium, the oppositely directed lobes could be somewhat symmetrically located with large asymmetries in brightness or flux density ratios."
+ In this case. the SW. lobe would be approaching us. which would be inconsistent with the inner double.," In this case, the $_{\rm out}$ lobe would be approaching us, which would be inconsistent with the inner double."
+ On the other hand. if the NEG lobe is approaching us. one would expect it to be brighter due to the ellects of relativistic motion as well as a denser environment. inferred. from. its closer distance to the radio nucleus.," On the other hand, if the $_{\rm out}$ lobe is approaching us, one would expect it to be brighter due to the effects of relativistic motion as well as a denser environment, inferred from its closer distance to the radio nucleus."
+ Phese suggest that the radio jets may have been intrinsically asymmetric during this cycle of activity., These suggest that the radio jets may have been intrinsically asymmetric during this cycle of activity.
+ The possibility of intrinsically asymmetric jets has been suggested. earlier. for example. for the highly asvmmetrie double-lobed radio source 120500|630 (Saikia et al.," The possibility of intrinsically asymmetric jets has been suggested earlier, for example, for the highly asymmetric double-lobed radio source B0500+630 (Saikia et al."
+ 1996). weak-cored one-sided. sources (Saikia et al.," 1996), weak-cored one-sided sources (Saikia et al."
+ 1989. 1990). the one-sided. radio emission in 3€273 which has been imaged with a dvnamic range o£ 10111 (Davis. Muxlow Conway 1985: Conway Davis 1994). the inner jets of the radio galaxies MIST (Ixovalev ct al.," 1989, 1990), the one-sided radio emission in 3C273 which has been imaged with a dynamic range of $^4$ :1 (Davis, Muxlow Conway 1985; Conway Davis 1994), the inner jets of the radio galaxies M87 (Kovalev et al."
+ 2007) and NGC6251 (Jones 1986: Jones Wehrle 2002). and the optical counter jet in the radio galaxy ος0012 (rraix-Burnet 1997).," 2007) and NGC6251 (Jones 1986; Jones Wehrle 2002), and the optical counter jet in the radio galaxy 3C66B (Fraix-Burnet 1997)."
+ On the theoretical side. Icke (1983) has suggested a hvedrocdynamical model for gas Hows in the nuclear region that can alfect the de Laval nozzle causing one-sidedjets. while Wang. Sulkanen Lovelace (1992) have explored models where the ratio of jet Iuminosities depends directly on the degree of asymmetry of the magnetic field in the nuclear region.," On the theoretical side, Icke (1983) has suggested a hydrodynamical model for gas flows in the nuclear region that can affect the de Laval nozzle causing one-sided jets, while Wang, Sulkanen Lovelace (1992) have explored models where the ratio of jet luminosities depends directly on the degree of asymmetry of the magnetic field in the nuclear region."
+ To explain the extreme asvnimetrics onc might also consider the possibility that the orientation of the jet night change during the cillerent eveles of activity: in the extreme cases the approaching jet during one evele may be the receding one in the next or vice versa., To explain the extreme asymmetries one might also consider the possibility that the orientation of the jet might change during the different cycles of activity; in the extreme cases the approaching jet during one cycle may be the receding one in the next or vice versa.
+ However. the alignment of all the four components in 4C02.27 suggests that this is not a likely scenario for this source.," However, the alignment of all the four components in 4C02.27 suggests that this is not a likely scenario for this source."
+ Another interesting aspect on the episodic nature of 4C02.27 is the presence of a hotspot in the SW. lobe., Another interesting aspect on the episodic nature of 4C02.27 is the presence of a hotspot in the $_{\rm out}$ lobe.
+ This is not unique to 4C02.27., This is not unique to 4C02.27.
+ While the outer doubles are often diffuse as for example in 1453|3308 (Schocnmakers ct al., While the outer doubles are often diffuse as for example in J1453+3308 (Schoenmakers et al.
+ 2000a: Ixonar et al., 2000a; Konar et al.
+ 2006). hotspots are also sometimes seen in the outer lobes as for example in the northern lobe of D1834|620 (Schoenmakers et al.," 2006), hotspots are also sometimes seen in the outer lobes as for example in the northern lobe of B1834+620 (Schoenmakers et al."
+ 20000)., 2000b).
+ This can be used to estimate the time scale of interruption of energy. supply., This can be used to estimate the time scale of interruption of energy supply.
+ For twpical sizes of hotspots of a few to 10 kpc in large racio sources (cf., For typical sizes of hotspots of a few to $\sim$ 10 kpc in large radio sources (cf.
+ Jevakumar Saikia 2000). the hotspots are expected to fade in about M10) 107 ve (e.g. Kaiser. Schocnmakers Rotttecring 2000).," Jeyakumar Saikia 2000), the hotspots are expected to fade in about $\sim$ $^4$$-$ $^5$ yr (e.g. Kaiser, Schoenmakers Rötttgering 2000)."
+ This is a small [fraction of the time it takes for the material in the jets of large raclio sources to reach the hotspots from the radio core., This is a small fraction of the time it takes for the material in the jets of large radio sources to reach the hotspots from the radio core.
+ Phorefore it is reasonable to assume that the hotspot fades soon after the last jet material passes through them., Therefore it is reasonable to assume that the hotspot fades soon after the last jet material passes through them.
+ The presence of a hotspot in the SW. lobe implies that it still. receives jet material., The presence of a hotspot in the $_{\rm out}$ lobe implies that it still receives jet material.
+ Phe travel time of the jet material from the core to the hotspot. 4. is given by {δι νι where Dy. is the physical distance of the hotspot [rom the core and View ds Che velocity of the jet.," The travel time of the jet material from the core to the hotspot, $t_{\rm j}$, is given by $D_{\rm hs}$ $_{\rm {jet}}$, where $D_{\rm hs}$ is the physical distance of the hotspot from the core and $_{\rm {jet}}$ is the velocity of the jet."
+ In our case. /; is 22.9. 10° vr for an inclination angle. ó40. and a jet. velocity of 0.5c.," In our case, $t_{\rm j}$ is $\sim$ $\times$ $^6$ yr for an inclination angle, $\phi\sim40^\circ$, and a jet velocity of 0.5c."
+ Llowever. this estimate will be allected by light. travel time ellects due to the orientation of the source axis.," However, this estimate will be affected by light travel time effects due to the orientation of the source axis."
+ For example. if the SW. lobe is on the receding side. it will take longer for the information to reach us compared with it being on the approaching side.," For example, if the $_{\rm out}$ lobe is on the receding side, it will take longer for the information to reach us compared with it being on the approaching side."
+ The observed time dillerence. fans. between the ejection of the last jet material ancl its arrival at the hotspot is ~1.S and 3.9 Myr depending on the orientation. of the source.," The observed time difference, $t_{\rm {obs}}$, between the ejection of the last jet material and its arrival at the hotspot is $\sim$ 1.8 and 3.9 Myr depending on the orientation of the source."
+ I£ the time scale of interruption of the jet is larger than £i. the hotspot in the SW lobe and the inner structure cannot be observed simultaneously.," If the time scale of interruption of the jet is larger than $t_{\rm {obs}}$, the hotspot in the $_{\rm out}$ lobe and the inner structure cannot be observed simultaneously."
+ Therefore. the interruption of jet activity must be less than Las.," Therefore, the interruption of jet activity must be less than $t_{\rm {obs}}$."
+" Also. within this time period the inner double forms. so that the time scale for interruption is less than £i, "," Also, within this time period the inner double forms, so that the time scale for interruption is less than $t_{\rm {obs}}$."
+I is also of interest to note that the time scale of interruption [or this lobe with a hotspot is much smaller than for sav J1453|3308 which has cdilfuse outer lobes., It is also of interest to note that the time scale of interruption for this lobe with a hotspot is much smaller than for say J1453+3308 which has diffuse outer lobes.
+ The dynamical ancl spectral ages of the diffuse outer lobes of 1453|3308 are 7215 and 50 Myr. while that of the inner double is only 2 Myr. suggesting a much longer time scale of interruption (Ixaiser et al.," The dynamical and spectral ages of the diffuse outer lobes of J1453+3308 are $\sim$ 215 and 50 Myr, while that of the inner double is only $\sim$ 2 Myr, suggesting a much longer time scale of interruption (Kaiser et al."
+ 2000: Ixonar et al., 2000; Konar et al.
+ 2006)., 2006).
+ The quasar 4002.27 with an overall linear size of ~470 kpe appears to exhibit signs of episodic activity and we classify it as a DDRQ., The quasar 4C02.27 with an overall linear size of $\sim$ 470 kpc appears to exhibit signs of episodic activity and we classify it as a DDRQ.
+ The source also exhibits evidence of an intrinsic asvmmetry of the oppositely-cirecteck jets., The source also exhibits evidence of an intrinsic asymmetry of the oppositely-directed jets.
+ Although most DORGs appear to be associated with large radio galaxies. often with sizes larger than approximately 1 Alpe. signs of episodic activity are seen in smaller sources as well (Schoenmakers οἱ al.," Although most DDRGs appear to be associated with large radio galaxies, often with sizes larger than approximately 1 Mpc, signs of episodic activity are seen in smaller sources as well (Schoenmakers et al."
+ 2000a.b: Saikia et al.," 2000a,b; Saikia et al."
+ 2006 and references therein)., 2006 and references therein).
+ Assuming an inclination angle of 40° to the line of sight. the intrinsic size of 4602.27. would. be ~730 kpe. comparable to some of the known or candidate DDRGs (Jamrozy et al.," Assuming an inclination angle of $^\circ$ to the line of sight, the intrinsic size of 4C02.27 would be $\sim$ 730 kpc, comparable to some of the known or candidate DDRGs (Jamrozy et al."
+ 2009)., 2009).
+ Considering the tendency for DDRGs to often have large sizes. we examined the structures of large quasars with sizes larger than ~S00 kpe (Table 1). including the quasar type object. | 6541 (Lara οἱ al.," Considering the tendency for DDRGs to often have large sizes, we examined the structures of large quasars with sizes larger than $\sim$ 800 kpc (Table 1), including the quasar type object $+$ 6541 (Lara et al."
+ 2001)., 2001).
+ These do not show evidence of episodie activity. which along with our search of the literature suggests that such objects are not common.," These do not show evidence of episodic activity, which along with our search of the literature suggests that such objects are not common."
+ A deep low-frequency search, A deep low-frequency search
+We iterate the last two steps until the number of particles with negative energy is constant.,We iterate the last two steps until the number of particles with negative energy is constant.
+ This final value gives the sell-bound mass of the satellite míutr)., This final value gives the self-bound mass of the satellite $m_{sb}(\tau)$.
+ In Figure 6. we compare the evolution of the sell-bound mass with the evolution of the mass within the tidal racius., In Figure \ref{fig:mloss} we compare the evolution of the self-bound mass with the evolution of the mass within the tidal radius.
+ 3old lines show median values. while shaded areas show the first ancl third quartile range of the distribution at each time T.," Bold lines show median values, while shaded areas show the first and third quartile range of the distribution at each time $\tau$."
+ Vhe trend of the two estimates is similar. but the mass within rz; is always smaller than the self-bouncl mass.," The trend of the two estimates is similar, but the mass within $r_t$ is always smaller than the self-bound mass."
+ Lt is not surprising that the two definitions do not agree., It is not surprising that the two definitions do not agree.
+ In fact. they use rather different criteria to identifv the satellites.," In fact, they use rather different criteria to identify the satellites."
+ The mass within the tidal radius is defined by the spatial clistribution of particles alone. Ίσα. (6)):," The mass within the tidal radius is defined by the spatial distribution of particles alone, Eq. \ref{eq:trad}) );"
+ the self-bound mass also uses information coming from the particle velocities., the self-bound mass also uses information coming from the particle velocities.
+ Because particles in the outskirt ofa satellite will take roughly a frec-Fall time to be accelerated to the cluster velocity dispersion. they will stick to the satellite for some time after they fall outside the tidal racüus. ancl will be counted in the self-bounc mass.," Because particles in the outskirt of a satellite will take roughly a free-fall time to be accelerated to the cluster velocity dispersion, they will stick to the satellite for some time after they fall outside the tidal radius, and will be counted in the self-bound mass."
+ Thus. the estimate based on tidal (ασ is actually a conservatively small measure of the mass associated with the satellite.," Thus, the estimate based on tidal radii is actually a conservatively small measure of the mass associated with the satellite."
+ We use the evolution of the mass within à as our second estimate of survival times., We use the evolution of the mass within $r_t$ as our second estimate of survival times.
+ Figure ο shows that survival times. defined by the condition: mtr;Tan)=0. we Torr=75. 1.5. 4 and 2.5 Garr [rom small to large masses.," Figure \ref{fig:mloss} shows that survival times, defined by the condition: $m(r_t; \tau_{sur}) = 0$, are $\tau_{sur} = 7.5,$ $7.5$, $4$ and $2.5$ Gyr from small to large masses."
+ These survival times are slightly [larger than the survival times derived. with Equation (4))., These survival times are slightly larger than the survival times derived with Equation \ref{eq:tsur}) ).
+ However. they are consistent with those survival times. because both distributions have large scatter.," However, they are consistent with those survival times, because both distributions have large scatter."
+ The longer survival time of small satellites is due to the combination of two ellects: (1) smaller. satellites are more compact (Lormen 1997: see also Section. ??)): (2) dynamical friction is less effective on smaller satellites: in fact. smaller satellites have larger distances from the cluster center (Figure 3)). and sulfer a weaker globaltice.," The longer survival time of small satellites is due to the combination of two effects: (1) smaller satellites are more compact (Tormen 1997; see also Section \ref{sec:instr}) ); (2) dynamical friction is less effective on smaller satellites; in fact, smaller satellites have larger distances from the cluster center (Figure \ref{fig:orbits}) ), and suffer a weaker globaltide."
+ These results suggest that the high force. and. mass, These results suggest that the high force and mass
+three canonical values of e.,three canonical values of $\epsilon$.
+ Lo additio:. the values of the correlation length extrapolated to z=0 ra(0)) are tabulated in Table { or three dillerent. values of Q4.," In addition, the values of the correlation length extrapolated to $z = 0$ $r_0(0)$ ) are tabulated in Table \ref{r0} for three different values of $\Omega_0$."
+ Tie strong dependeuce of the correlation scale length ou the valie of the evolutionary parameter (e )ls ac lirect result of the Cosmological term (G(z)) in Equation ].., The strong dependence of the correlation scale length on the value of the evolutionary parameter $\epsilon$ ) is a direct result of the Cosmological term $G(z)$ ) in Equation \ref{limberE}. .
+ Relative to the precictious of litear theory (e=0.8). the results for fixed. clustering iu co-moviiο coordinates (e=—1.2). are supoessed by an additional [actor of (1+2). or roughly a factor of 1.572 at the redshifts of inteest (recall that ," Relative to the predictions of linear theory $\epsilon = 0.8$ ), the results for fixed clustering in co-moving coordinates $\epsilon = -1.2$ ), are suppressed by an additional factor of $\sim (1 + z)$, or roughly a factor of $1.5$ $2$ at the redshifts of interest (recall that $r_0(z) \propto G(z)^{-\gamma}$ )."
+When comparing our results to previous spectroscopic results. we cleary show excellent agreement with recent measurements (Carlbergefaf1999:Small1999) when the predictions of linear theory are used t« quautifv the evolution of clustering.," When comparing our results to previous spectroscopic results, we clearly show excellent agreement with recent measurements \citep{carlberg99,small99} when the predictions of linear theory are used to quantify the evolution of clustering."
+ This is extreijely. encouraging [or our technique. as we uniformly sample a larger redshift rauge. showing. Lor he first tine withiu the same dataset. the slight decrease in the correlation strength for z«l.as] oedicted by mocels of galaxy formation (Baughetal.1999).," This is extremely encouraging for our technique, as we uniformly sample a larger redshift range, showing, for the first time within the same dataset, the slight decrease in the correlation strength for $z < 1$, as predicted by semi-analytic models of galaxy formation \citep{baugh99}."
+. On the other hand. the measurement oL the correlation scale length from the CFRS data (LeFévreefal.1996) are ouly in agreement with our results for fixed clustering in «‘O-noving coordinates. which disagrees with the hierarchical erowth of dark matter halos (Baughetal.1999).," On the other hand, the measurement of the correlation scale length from the CFRS data \citep{leFevre96} are only in agreement with our results for fixed clustering in co-moving coordinates, which disagrees with the hierarchical growth of dark matter halos \citep{baugh99}."
+. The discrepancy between he CFRS and other nieasuremeuts is most likely due to their relatively sinall sample size. their sinall fields. aud their ueplec X the redshift evolution of the Lumiuosity [uuctiou (Carlbergefed.1999).," The discrepancy between the CFRS and other measurements is most likely due to their relatively small sample size, their small fields, and their neglect of the redshift evolution of the Luminosity function \citep{carlberg99}."
+. While important. the evolution of the angular correlation fuuctiou with redshilt s1uootlies over the galaxy luminosity Duuction. ignoring variations in clustering between galaxies of different intrinsic luminosity.," While important, the evolution of the angular correlation function with redshift smoothes over the galaxy luminosity function, ignoring variations in clustering between galaxies of different intrinsic luminosity."
+" As a result. we subdivided our sample into three redshift intervals 0.6. 0.[Xz<0.8. 0.6€z< 1.0). aud measured the angular correlation function as a function of tli C aud B absolute magnitude iu intervals of 2.0. from 22"" to 12""."," As a result, we subdivided our sample into three redshift intervals $0.2 \leq z \leq 0.6$ , $0.4 \leq z \leq 0.8$, $0.6 \leq z
+\leq 1.0$ ), and measured the angular correlation function as a function of both $U$ and $B$ absolute magnitude in intervals of $2.0^m$, from $22^m$ to $12^m$."
+" In order to improve the uunber of galaxies in the faint euc of our analysis. we rebinued the data so that the faint. bin was our imagulitudes wide 16-12"")."," In order to improve the number of galaxies in the faint end of our analysis, we rebinned the data so that the faint bin was four magnitudes wide $16^m$ $12^m$ )."
+ In the eid. we obtained twelve different. tmeastuements of he multivariate augular correlation fuction Qe(0.z. M)) asa function of both C and B absolute nagenituces.," In the end, we obtained twelve different measurements of the multivariate angular correlation function $w(\theta, z, M)$ ) as a function of both $U$ and $B$ absolute magnitudes."
+ The munber of subdivisions used iu this particular analysis reintrodiced οιie of the problems our new technique was desigued to avoid. namely the effects of siuall saiuple size.," The number of subdivisions used in this particular analysis reintroduced one of the problems our new technique was designed to avoid, namely the effects of small sample size."
+ We. therefore. ouly used a joint redshift-absolute maguLE.=de bin when the number of objects τι1 the bin exceeded one hundred galaxies.," We, therefore, only used a joint redshift-absolute magnitude bin when the number of objects in the bin exceeded one hundred galaxies."
+ From Figures T.. 8 (U aud B respectively). it is clear that here is no obvious evolution within a given redshift interval with intrinsic luminosity (althoug1 this«οι] be a result of too few galaxies).," From Figures \ref{umz}, \ref{bmz} $U$ and $B$ respectively), it is clear that there is no obvious evolution within a given redshift interval with intrinsic luminosity (although thiscould be a result of too few galaxies)."
+" Between redshift intervals. however. there is strong evollon i1 the amplitucle of the angular correlation function (A,.). whicl. given the wider redshift bins. is corpletely consistent"," Between redshift intervals, however, there is strong evolution in the amplitude of the angular correlation function $A_w$ ), which, given the wider redshift bins, is completely consistent"
+also produces spiral structure.,also produces spiral structure.
+ Iu this run the azimuth averaged temperatures were similar to fig. 2..," In this run the azimuth averaged temperatures were similar to fig. \ref{fig:temp-prof},"
+ but the temperature iu the sliocked regious are not as extreme., but the temperature in the shocked regions are not as extreme.
+ Water ice would still be vaporized everywhere but more refractory species iuay not be., Water ice would still be vaporized everywhere but more refractory species may not be.
+ In the outer disk. egrains less than 1 uuu in size which pass throughOm such a shock aud which contain water ice will be vaporized ou a timescale of 101—10° s depending on the temperature (Eneel.Luuine&Lewis1990:Leuzuni.GailHeunine1995) and the remaining more inert species may cdissagregate.," In the outer disk, grains less than $\sim 1$ mm in size which pass through such a shock and which contain water ice will be vaporized on a timescale of $\lesssim10^4-10^5$ s depending on the temperature \citep{ELL90,LGH95} and the remaining more inert species may dissagregate."
+ Iu this region. passage through the warmest part of a spiral arm requires ~]—2 wr. so sufficient time exists to return gralus of this size to the gas pliase.," In this region, passage through the warmest part of a spiral arm requires $\sim 1-2$ yr, so sufficient time exists to return grains of this size to the gas phase."
+ Grain growth may still occur between the spiral aruis aud. when the spiral arms have decayed. but must begiu with gaseous material each time. so growth of solid material iuto larger entities will be suppressed.," Grain growth may still occur between the spiral arms and when the spiral arms have decayed, but must begin with gaseous material each time, so growth of solid material into larger entities will be suppressed."
+ Temperatures at high altitudes are lower. but graius which form there will teuc to sink to the midplane as they grow larger aud also be destroyed.," Temperatures at high altitudes are lower, but grains which form there will tend to sink to the midplane as they grow larger and also be destroyed."
+ Therefore Jovian plauet formation by tlie core accretion mechanisin will occur much more slowly. if at all in this system.," Therefore Jovian planet formation by the core accretion mechanism will occur much more slowly, if at all in this system."
+ The weak link remaining in the argument agalust the core accretiou mechanism is the lack of kuowledege of the micropliysics importaut for dust coagulation., The weak link remaining in the argument against the core accretion mechanism is the lack of knowledge of the microphysics important for dust coagulation.
+ For example. one could iuiagiue tliat erowth of silicate and iron grains is catalyzed by temporarily enliauced. cross sections as mantles of more volatile material [orm on their surfaces aud the gravitational torques produced by the ünary interaction enhauce mixine throughout the disks.," For example, one could imagine that growth of silicate and iron grains is catalyzed by temporarily enhanced cross sections as mantles of more volatile material form on their surfaces and the gravitational torques produced by the binary interaction enhance mixing throughout the disks."
+ Even if this type of interaction takes R9]ace. the eventual formation of planet sized objects remains in doubt.," Even if this type of interaction takes place, the eventual formation of planet sized objects remains in doubt."
+" As the rocky aggregates erow. conditions appropriate for dust coagulation (e.g. ""perfect sticking) break down aud collisions jetweeu particles become Increasiugly disruptive due to the finite strength of the aggregates."," As the rocky aggregates grow, conditions appropriate for dust coagulation (e.g. `perfect sticking') break down and collisions between particles become increasingly disruptive due to the finite strength of the aggregates."
+ The disruption of solid bodies depeuds strougly on the relative velocity of the inipactor ancl arget particles. with disruption occurring lor velocities 21-3 kim/s for planetesimal sized targets (<1 kim) (Benz&Asphaug1999).," The disruption of solid bodies depends strongly on the relative velocity of the impactor and target particles, with disruption occurring for velocities $\gtrsim$ 1-3 km/s for planetesimal sized targets $\lesssim 1$ km) \citep{BA99}."
+. On average in au accretion disk. the relative velocity of dlanetesimals will be proportional to their eccentricities. tre2tyr.," On average in an accretion disk, the relative velocity of planetesimals will be proportional to their eccentricities, $v_{\rm rel}\approx ev_{\rm orb}$."
+ For our model. a relative velocity of 1 kua/s corresponds to an eccentricity ofe~0.05 at 1 AU. or e~0.15 at 10 AU.," For our model, a relative velocity of 1 km/s corresponds to an eccentricity of $e\sim 0.05$ at 1 AU, or $e\sim 0.15$ at 10 AU."
+ We have seen that gravitational torques are stroug enough to generate large amplitucle spiral structure as they drive the gas outo eccentric orbits., We have seen that gravitational torques are strong enough to generate large amplitude spiral structure as they drive the gas onto eccentric orbits.
+ Gas eccentricities are quickly damped due to shock clissipation. but plauetesimals are ouly weakly coupled to the gas aud will have time to encounter other objects aud collide or to increase their eccentricities still further as the evolution proceeds.," Gas eccentricities are quickly damped due to shock dissipation, but planetesimals are only weakly coupled to the gas and will have time to encounter other objects and collide or to increase their eccentricities still further as the evolution proceeds."
+ Π particle eccentricities cau grow to €0.1. we expect that the growth of kilometer sized bodies will be suppressed tn binary systems such as the oue modeled here.," If particle eccentricities can grow to $e\sim0.1$, we expect that the growth of kilometer sized bodies will be suppressed in binary systems such as the one modeled here."
+ However. a more detailed analysis must be done iu order to coustrain this possibility.," However, a more detailed analysis must be done in order to constrain this possibility."
+ The couclusious regarding planet formation will remain valid as long as the temperatures determiued from the model are lower limits on the temperatures present in real systems., The conclusions regarding planet formation will remain valid as long as the temperatures determined from the model are lower limits on the temperatures present in real systems.
+ If the, If the
+and global parameters of the true. physical properties are reproduced well by the parametric fit. with svstematic error «lO pper cent at 7200.,"and global parameters of the true physical properties are reproduced well by the parametric fit, with systematic error $< 10$ per cent at $r_{200}$."
+ This result is fairly consistent with 7.. who find that the hyclrostatic mass estimated. from a combination of «3 model fits to the X-ray surface brightness data and spatial temperature information agrees to within LOpper cent (rZ rsoo) for the simulated: clusters which contain feedback.," This result is fairly consistent with \cite{Kay:2004b}, who find that the hydrostatic mass estimated from a combination of $\beta$ model fits to the X-ray surface brightness data and spatial temperature information agrees to within $10$ per cent $r \la
+r_{500}$ ) for the simulated clusters which contain feedback."
+ They. find that the presence of feedback ooduces a higher degree of thermalisation than if the cluster were simply ceseribec by ài non-raciative model and so he estimated mass is close to the true value., They find that the presence of feedback produces a higher degree of thermalisation than if the cluster were simply described by a non-radiative model and so the estimated mass is close to the true value.
+ In the low signal-to-noise regime the increase in noise leads to much arecr widths in the probability clistributions of the eluster quantities and. while shifting the peaks of the distributions. ave still consistent with the known simulation values within he errors.," In the low signal-to-noise regime the increase in noise leads to much larger widths in the probability distributions of the cluster quantities and, while shifting the peaks of the distributions, are still consistent with the known simulation values within the errors."
+" reftableevidence,eluesgivesthenaturallogarilhinoflh", \\ref{table:evidence_values} gives the natural logarithm of the evidence for different models given the data.
+beevidence feWablU axe Ομ igh noiseregimessupporllheinclusionof hyper, The evidence values for both low and high signal-to-noise regimes support the inclusion of hyper-parameters.
+ heratiosofbest f HvaluesforthesSZand rayhyper peremclersineachregimearel land3.5: l(highsigneal noisc).and2.5: landl.l: lowsignaltlonoise). fork’ Bland B3respectively. indicalinglhalthehyper paramelersweightthetikelihoodin favourof," The ratios of best fit values for the SZ and X-ray hyper-parameters in each regime are 15:1 and 3.5:1 (high signal-to-noise), and 2.5:1 and 1.1:1 (low signal-to-noise), for FB1 and FB3 respectively, indicating that the hyper-parameters weight the likelihood in favour of the SZ data."
+ihesZdata M hile hec Πα”... raydalaaremuch moresensitivelovariationsoverthisrange," While the CBI2 SZ data are insensitive to variation on scales smaller than kpc, the X-ray data are much more sensitive to variations over this range."
+ Lhencdtvicentaepees /dness hesesni ragdelta. whileslillproducingarelativelggood fillothes Zdata," The nature of the parametric model is such that it is inherently smooth on these smaller scales; hence clumping of the gas will deviate the model from the X-ray data, while still producing a relatively good fit to the SZ data."
+ Habart ab orbe raysurfacebrightness, Therefore the weights will typically favour the SZ data relative to the X-ray surface brightness.
+ Lithecascof hel B3data. therelalirewesgralvgpyrr brrone pila raydataarcacloscrmalchlothemodelovertheconsidercdradiusran," In the case of the FB3 data, the relative weightings are more similar since the X-ray data are a closer match to the model over the considered radius range."
+ ως noiscleadstoadecreascintheraltiosofhyper paramcltersforcacheluster.andsobothlhes Zand X ragdatasclsprovideamorcconsislentmalehtolhemodel.," A reduction in signal-to-noise leads to a decrease in the ratios of hyper-parameters for each cluster, and so both the SZ and X-ray data sets provide a more consistent match to the model."
+ Even though the simulated. clusters are clearly not isothermal. for the purposes of comparison to previous work and to demonstrate the svstematic cilferences in. derived cluster properties using the two different models. we perform joint fits using the single isothermal 2 model.," Even though the simulated clusters are clearly not isothermal, for the purposes of comparison to previous work and to demonstrate the systematic differences in derived cluster properties using the two different models, we perform joint fits using the single isothermal $\beta$ model."
+ The electron density and temperature are then given by where poo and Zig are the central electron density and temperature respectively. roore is the core radius. and 3 isa parameter that determines the large scale behaviour of the electron density.," The electron density and temperature are then given by where $n_\rmn{e0}$ and $T_\rmn{e0}$ are the central electron density and temperature respectively, $r_\rmn{core}$ is the core radius, and $\beta$ is a parameter that determines the large scale behaviour of the electron density."
+ Figures Sand 9— show the resulting posterior distributions for the model parameters. the profiles. and elobal values within rooo.," Figures \ref{figure:kay_simulations_isothermal_beta_params} and \ref{figure:kay_simulations_isothermal_beta_results} show the resulting posterior distributions for the model parameters, the profiles, and global values within $r_{200}$."
+ There is a strong systematic dillerence between the estimated and true simulation values when using this model. resulting from over-estimates of the electron. temperature and. SZ decrement.," There is a strong systematic difference between the estimated and true simulation values when using this model, resulting from over-estimates of the electron temperature and SZ decrement."
+ The result is a systematic over-estimate of the total mass of ~ pper cent at σο�� and an under-estimate of the total mass ator«rsoo., The result is a systematic over-estimate of the total mass of $\sim$ per cent at $r_{200}$ and an under-estimate of the total mass at $r < r_{500}$.
+ These results are consistent with the findings by 7. who measure similar errors in estimating the total mass based upon the isothermal mocdel., These results are consistent with the findings by \cite{Kay:2004b} who measure similar errors in estimating the total mass based upon the isothermal model.
+ Phe svstematie error in derived PY ας he: bighneada lows to-noise scenarios. and verelore the introduction of both thermal noise and intrinsic CALB anisotropy is not enough to dominate over the ellects of the intrinsic isothermal +7 model discrepancy. with the simulations.," The systematic error in derived cluster properties is seen in both the high and low signal-to-noise scenarios, and therefore the introduction of both thermal noise and intrinsic CMB anisotropy is not enough to dominate over the effects of the intrinsic isothermal $\beta$ model discrepancy with the simulations."
+ lt is informative to compare the relative quality of the fit that the cntropy-basecl ancl single isothermal ΙΟ ΠΩ ΑΛ ΠΟ ο nla alwogodcad ts she o," It is informative to compare the relative quality of the fit that the entropy-based and single isothermal $\beta$ models give, based upon their respective logarithmic evidence values."
+ iyadkondtylow bbihon Theisqudyannaeey qnse Hbeygipalégdatlsidewets αλα ob sibbel be Xxception of the low signal-φομ inddomwiskgnahe FBS simulated cluster.," In almost all the cases the evidence for the isothermal $\beta$ model is significantly lower than that of the entropy-based model, with the exception of the low signal-to-noise scenario for the FB3 simulated cluster."
+ This is due o the [lat core nature of the X-ray surface brightness profile or this cluster. as a result of the displacement of the gas »ealkk in the central region.," This is due to the flat core nature of the X-ray surface brightness profile for this cluster, as a result of the displacement of the gas peak in the central region."
+ As such when significantly lower signal-to-noise data is used the isothermal 3 mocoel. with its lat core behaviour at small racii. provides an equally good it GP not slightly. better) than the entropy-based. mocel.," As such when significantly lower signal-to-noise data is used the isothermal $\beta$ model, with its flat core behaviour at small radii, provides an equally good fit (if not slightly better) than the entropy-based model."
+ llowever in the case of the FBI simulated cluster. the X-ray ποιο is strongly peakecl in the central region ancl therefore he isothermal 3 moclel provides a significantly poorer Lit o both high and low signal-to-noise data.," However in the case of the FB1 simulated cluster, the X-ray profile is strongly peaked in the central region and therefore the isothermal $\beta$ model provides a significantly poorer fit to both high and low signal-to-noise data."
+" This can be seen upon visual inspection of the electron. density. profiles. [or FBI in refheure:kaymudlations;solhermal,eta,esults. wheretheisothermats model fails to provide a suitable fit at both small and large raclil."," This can be seen upon visual inspection of the electron density profiles for FB1 in \\ref{figure:kay_simulations_isothermal_beta_results}, where the isothermal $\beta$ model fails to provide a suitable fit at both small and large radii."
+" The evidence values in reflable:evicence,.aluessupporttheinelusiono fh paramelersinbolhthehighandlouwsignal yperlo noisescenariosforlheisolhermal3 model.", The evidence values in \\ref{table:evidence_values} support the inclusion of hyper-parameters in both the high and low signal-to-noise scenarios for the isothermal $\beta$ model.
+ The relative weightings of each the SZ ancl X-ray.| cata sets are— to 30:1 (FBI) and 3:1 (FBS) for the high signal-to-noise data. and 4:1 and 1.1: for the corresponding low signal-to-noise data.," The relative weightings of each the SZ and X-ray data sets are to 30:1 (FB1) and 3:1 (FB3) for the high signal-to-noise data, and 4:1 and 1.1:1 for the corresponding low signal-to-noise data."
+ The likelihood calculation is therefore weighted in favour of the SZ data. except in the case of the low signal-to-noise 173 data where both are almost equally. favoured.," The likelihood calculation is therefore weighted in favour of the SZ data, except in the case of the low signal-to-noise FB3 data where both are almost equally favoured."
+ We have developed a parametric mocel for the gas in galaxy clusters. based. on three physical assumptions: the dark matter follows an NEW profile. the gas entropy. can be deseribed by a power law with a Uattened core. and the," We have developed a parametric model for the gas in galaxy clusters, based on three physical assumptions: the dark matter follows an NFW profile, the gas entropy can be described by a power law with a flattened core, and the"
+"planetary nature of the candidate must be established statistically, as in the latter two cases.","planetary nature of the candidate must be established statistically, as in the latter two cases."
+ The sinallest planet discovered to date. Nepler-LObb. was announced recently by Batalhactal.(2011).. ud is the \T\Uissiow’s first rocky planet.," The smallest planet discovered to date, b, was announced recently by \cite{Batalha:11}, and is the Mission's first rocky planet."
+ It has a measured radius of 1.116!UnRy and a mass of {οD3Mi. leading to a mean density of SS!S4ecm? that iuplies a significant iron nias fraction (Datalhactal. 2011).," It has a measured radius of $1.416^{+0.033}_{-0.036}
+\,R_{\earth}$ and a mass of $4.6^{+1.2}_{-1.3}\,M_{\earth}$, leading to a mean density of $8.8^{+2.1}_{-2.9}\,{\rm g~cm}^{-3}$ that implies a significant iron mass fraction \citep{Batalha:11}."
+. Its parent star. IWepler-lO 111901151. 1119021305|5011286). is relatively bright among the ttargets uunaegnitude Ap= 10.96) and displays periodic sienals with periods of 0.81 days aud. 15.3 davs. aud fiux decrements (enorme Π darkeniuz) of 152+4l| ppm and 376E9 ppui respectively (Batalhaetal.2011).," Its parent star, Kepler-10 11904151, 119024305+5014286), is relatively bright among the targets magnitude $K\!p = 10.96$ ) and displays periodic signals with periods of 0.84 days and 45.3 days, and flux decrements (ignoring limb darkening) of $152 \pm
+4$ ppm and $376 \pm 9$ ppm, respectively \citep{Batalha:11}."
+. The exteusive observations that followed the detection of these signals are οσο im detail by those authors. and include the difficult 1ieasureinentτος of the refiex racdial-velocity motion of the star with a scimi-unplitude of ouly 23g aand a period that is consistent with the shorter signal.," The extensive observations that followed the detection of these signals are documented in detail by those authors, and include the difficult measurement of the reflex radial-velocity motion of the star with a semi-amplitude of only $3.3^{+0.8}_{-1.0}$ and a period that is consistent with the shorter signal."
+ As is customary also iu ground-based searches for transiting plaucts. the shapes of the spectral lines were exinuiued carefully to rule out changes of simular amplitude correlating with orbital phase that nuelt indicate a false positive. such as a background echpsing binary (EB) blended with the target. or an ED physically associated with it.," As is customary also in ground-based searches for transiting planets, the shapes of the spectral lines were examined carefully to rule out changes of similar amplitude correlating with orbital phase that might indicate a false positive, such as a background eclipsing binary (EB) blended with the target, or an EB physically associated with it."
+ Towever the precision of the measurements (bisector spaus) compared to the ill BV amplitude did not allow such changes to be ruled out unauubieuouslv., However the precision of the measurements (bisector spans) compared to the small RV amplitude did not allow such changes to be ruled out unambiguously.
+ False positive scenarios were explored with the aid ofBLENDER.. a technique that models the transit light. curves to test a wide range of blend coufiguratious (Torresetal.2011).. and it was found that the overwhcluine majority of them can be rejected.," False positive scenarios were explored with the aid of, a technique that models the transit light curves to test a wide range of blend configurations \citep{Torres:11}, and it was found that the overwhelming majority of them can be rejected."
+ This aud other evidence presented by etal.(2011). allowed the plauctary nature of LOD) to be established with very high coufidence., This and other evidence presented by \cite{Batalha:11} allowed the planetary nature of b to be established with very high confidence.
+ This was uot the case. however. for the L5-day period signal referred to as OOhject of Tuterest 72.02). which is the subject of this paper.," This was not the case, however, for the 45-day period signal referred to as Object of Interest 72.02), which is the subject of this paper."
+ No significant RV signal was detected at this period. aud oulv an upper linüt on its amplitude could be placed.," No significant RV signal was detected at this period, and only an upper limit on its amplitude could be placed."
+ UsingBLENDER.. Batalhaetal.(2011) were able to rule out a Lluge fraction of the bleud scenarios involving circular orbits Guchiding hierarchical triples). but eccentric orbits were not explored because of the increased complexity of the problem aud the unch larger space of parameters for false positives.," Using, \cite{Batalha:11} were able to rule out a large fraction of the blend scenarios involving circular orbits (including hierarchical triples), but eccentric orbits were not explored because of the increased complexity of the problem and the much larger space of parameters for false positives."
+ While circular orbits are a reasonable assumption for ορίιο-100) because of the stroug effects of tidal forces at close range. this is not true for oon account of it wich longer orbital period (see.Alazeh 2008): eccentric orbits cau not be ruled out.," While circular orbits are a reasonable assumption for b because of the strong effects of tidal forces at close range, this is not true for on account of its much longer orbital period \citep[see, e.g.,][]{Mazeh:08}; eccentric orbits can not be ruled out."
+ This provides the motivation for the present work. iu which we set out to examine all viable astroplivsical false positive scenarios for wwith the eoal of validating it as à bona-fide planet.," This provides the motivation for the present work, in which we set out to examine all viable astrophysical false positive scenarios for with the goal of validating it as a bona-fide planet."
+ Iu addition to improvements in tle nunodeling. wei bring to bear new near-infrared observations obtained with the SSpace Telescope in which the transits are clearly detected. as well as the complete arscual of follow-up observations gathered by the tte. including high-resolution adaptive optics iniaeiug id speckle interferometry. high-resolution spectroscopy. id an analysis based on the oobservations themselves of the difference nuages in id out of transit for positional displacements (ceutroid motion).," In addition to improvements in the modeling, we bring to bear new near-infrared observations obtained with the Space Telescope in which the transits are clearly detected, as well as the complete arsenal of follow-up observations gathered by the team, including high-resolution adaptive optics imaging and speckle interferometry, high-resolution spectroscopy, and an analysis based on the observations themselves of the difference images in and out of transit for positional displacements (centroid motion)."
+ All of these observations combined with the strong constraints provided bx ssigmificautly limit the kinds of blends that remain possible. auc as we describe below they allow us to claim with very Lieh coufidence that lis indecd a planct.," All of these observations combined with the strong constraints provided by significantly limit the kinds of blends that remain possible, and as we describe below they allow us to claim with very high confidence that is indeed a planet."
+ Its estimated radius is approximately of that of Neptune., Its estimated radius is approximately of that of Neptune.
+ With this. Iepler-10 becomes he Missions third confirmed nmulti-plauet svstem (afteretal.200101a) containing a transiting super-Earth-size ylanet aud at least one larger planet that also transits.," With this, Kepler-10 becomes the Mission's third confirmed multi-planet system \citep[after
+Kepler-9 and Kepler-11;][]{Holman:10, Lissauer:11a} containing a transiting super-Earth-size planet and at least one larger planet that also transits."
+ We beginran with a brief recapitulation of the technique. including recent improvements.," We begin with a brief recapitulation of the technique, including recent improvements."
+ We then oxeseut the oobservations at that help rule out mauyv blends. aud we sumnunanze additional constraints available from other observations.," We then present the observations at that help rule out many blends, and we summarize additional constraints available from other observations."
+ This is followed by the application of to Hin order to ideutifv all blends scenarios that can niunmic he transit heht curve., This is followed by the application of to in order to identify all blends scenarios that can mimic the transit light curve.
+ Next we combine this information with the other coustraimts and carry out a statistical assessnient of the false aluiu rate for the planuct ivpothesis. leading to the validation of the candidate as Ixepler-10 We couclude with a discussion of the οκ]ολο constitutioncc. of the new planet iu the ποτ of current models and the sienificauce of this type of validation.," Next we combine this information with the other constraints and carry out a statistical assessment of the false alarm rate for the planet hypothesis, leading to the validation of the candidate as c. We conclude with a discussion of the possible constitution of the new planet in the light of current models, and the significance of this type of validation."
+ The detailed morphology of a transit Πο curve (leugth of imeress/ceress. total duration) contains Huportant information that can be used to reject many false positive scenarios producing brightness variations that do not quite have the right shape. even though they may well match the observed trausit depth (sec.c.g.Suellenetal. 2009).," The detailed morphology of a transit light curve (length of ingress/egress, total duration) contains important information that can be used to reject many false positive scenarios producing brightness variations that do not quite have the right shape, even though they may well match the observed transit depth \citep[see,
+e.g.,][]{Snellen:09}."
+. citepTLTorres:01.Torres:11 takes advantage of this to explore a. verv large range— of scenarios. including backeround or foreground eclipsing binaries blended with the target. as well as cclipsing binarics plysically associated with the target iu a hierarchical triple configuration.," \\citep{Torres:04, Torres:11} takes advantage of this to explore a very large range of scenarios, including background or foreground eclipsing binaries blended with the target, as well as eclipsing binaries physically associated with the target in a hierarchical triple configuration."
+" Following the uotation introduced by Torresetal.(2011).. the objects composing the binary are referred to as the “secondary” and ""tertianv. audthe candidate is the “primar."," Following the notation introduced by \cite{Torres:11}, the objects composing the binary are referred to as the “secondary” and “tertiary”, and the candidate is the “primary”."
+ The tertiary can be either a star Guelucding a white dwarf) or a planet. aud the secondary can be a niain-sequeuce star or a (backeround) edant.," The tertiary can be either a star (including a white dwarf) or a planet, and the secondary can be a main-sequence star or a (background) giant."
+ With the help of model isochrones to set the stellar properties.," With the help of model isochrones to set the stellar properties,"
+(VIMS) onboard the Cassini Spacecrall have already provided. constraints on the eeonmeltrv of sell-gravitv wakes in the A and D rings (lechnanεἰaf2007:NicholsonandΠοια2010) and the architecture of the Cassini Division (Hedmane£af.2010).,"(VIMS) onboard the Cassini Spacecraft have already provided constraints on the geometry of self-gravity wakes in the A and B rings \citep{Hedman07,NH10} and the architecture of the Cassini Division \citep{Hedman10}."
+. llowever. (hese analvses only used a [fraction of the information returned by VIMS. because thev were based on light curves derived Irom a single spectral channel.," However, these analyses only used a fraction of the information returned by VIMS, because they were based on light curves derived from a single spectral channel."
+ During each occultation. VIAIS simultaneously measures (he opacity of (he rings over a range of wavelengths from 0.85 to 5.0 jan. which includes the strong water-ice absorption band at 3.1 jan. Thus each stellar occultation can in principle provide high-spatial- (ransmüssion specira of the rings.," During each occultation, VIMS simultaneously measures the opacity of the rings over a range of wavelengths from 0.85 to 5.0 $\mu$ m, which includes the strong water-ice absorption band at 3.1 $\mu$ m. Thus each stellar occultation can in principle provide high-spatial-resolution transmission spectra of the rings."
+ In. practice. the optical depth: of most regions in Saturns main rings does not vary with wavelength because nearly all of the particles in (he main rings are much Luger than the near-infrared. wavelengths observed.," In practice, the optical depth of most regions in Saturn's main rings does not vary with wavelength because nearly all of the particles in the main rings are much larger than the near-infrared wavelengths observed."
+ In (his geometrical optics limit (he transmission is essentially independent ol wavelength., In this geometrical optics limit the transmission is essentially independent of wavelength.
+ llowever. (he (ransmission can vary wilh wavelength when the particles are comparable in size to the observing wavelength.," However, the transmission can vary with wavelength when the particles are comparable in size to the observing wavelength."
+ Several features in Saturns rings are strongly. forward scattering in the visible and near-infrared. indicating that thev are composed primarily of micron-sized grains (Lloranviefa£2009) ancl (his has been confirmed [or the F-ring by detailed spectrophotometric analyses (Showalterefal.1992:Vahidiniaaf 2011).," Several features in Saturn's rings are strongly forward scattering in the visible and near-infrared, indicating that they are composed primarily of micron-sized grains \citep{Horanyi09} and this has been confirmed for the F-ring by detailed spectrophotometric analyses \citep{Showalter92, Vahidinia11}."
+. Searches for transmission variations using occullations of the F ring did not reveal any statistically significant. trends at ultra-violet or visible wavelengths (Boshefa£.2002).. but VIMS occeultations by the F ring and other similarly dusty ringlets in (he A-rine’s Encke Gap and (the Cassini Divisions Laplace Gap have revealed a narrow opacity dip in the Granusnission spectra near 2.87 jm. As discussed in detail below. this feature provides novel constraints on the composition and structure of these dusty rings.," Searches for transmission variations using Earth-based occultations of the F ring did not reveal any statistically significant trends at ultra-violet or visible wavelengths \citep{Bosh02}, but VIMS occultations by the F ring and other similarly dusty ringlets in the A-ring's Encke Gap and the Cassini Division's Laplace Gap have revealed a narrow opacity dip in the transmission spectra near 2.87 $\mu$ m. As discussed in detail below, this feature provides novel constraints on the composition and structure of these dusty rings."
+ Of particular interest is the ability of near-infrared stellar occultations to discern variations in the rings! particle size distribution on finer spatial scales than otherwise possible., Of particular interest is the ability of near-infrared stellar occultations to discern variations in the rings' particle size distribution on finer spatial scales than otherwise possible.
+ Our analvsis begins by describing the relevant observations and how thev were processed to obtain light curves., Our analysis begins by describing the relevant observations and how they were processed to obtain light curves.
+ Second. we examine an illustrative example of the óransmission spectra and demonstrate how the observed feature can be explained in terms of the Christiansen Effect associated with the strong water-ice absorption band centered al μην. We then discuss how the strength of this feature relates to the local particle size distribution.," Second, we examine an illustrative example of the transmission spectra and demonstrate how the observed feature can be explained in terms of the Christiansen Effect associated with the strong water-ice absorption band centered at $\mu$ m. We then discuss how the strength of this feature relates to the local particle size distribution."
+ In the future. we expect that combining these óransmission spectra with relevant reflectance spectra and phase curves will place light constraints on the size distribution. but such a complete photometric analvsis is bevond the scope of (his paper.," In the future, we expect that combining these transmission spectra with relevant reflectance spectra and phase curves will place tight constraints on the size distribution, but such a complete photometric analysis is beyond the scope of this paper."
+" Instead. we turn our attention to the variations in the (transmission spectra, which allow us to discern trends in (the particle size distribution."," Instead, we turn our attention to the variations in the transmission spectra, which allow us to discern trends in the particle size distribution."
+ We find that the strength of the opacity dip varies svstematically among the dilferent ringlets. demonstrating (hat these dusty rings do have somewhat different particle size distributions.," We find that the strength of the opacity dip varies systematically among the different ringlets, demonstrating that these dusty rings do have somewhat different particle size distributions."
+ In particular. we explore the spectral variations within the F ring itself.," In particular, we explore the spectral variations within the F ring itself,"
+not reach zero intensity in their bottoms. which suggests that the cloud in. which the lines originate does not obscure the entire continuum emitting source.,"not reach zero intensity in their bottoms, which suggests that the cloud in which the lines originate does not obscure the entire continuum emitting source."
+ That in turn suggests that this absorption system Is intrinsic to the quasar., That in turn suggests that this absorption system is intrinsic to the quasar.
+ The lower ejection velocity systems in the complex also do not reach zero intensity. which makes it reasonable to assume that they are also associated with the quasar (even though we cannot unambiguously establish whether they are saturated).," The lower ejection velocity systems in the complex also do not reach zero intensity, which makes it reasonable to assume that they are also associated with the quasar (even though we cannot unambiguously establish whether they are saturated)."
+ There are three (or four. if one includes the uncertain lab.=tam System) other heavy element systems with ejection velocities lower than the complex.," There are three (or four, if one includes the uncertain $z_{\rm
+abs}=z_{\rm em}$ system) other heavy element systems with ejection velocities lower than the complex."
+ Two of these systems are clearly infalling: even if the emission redshift of the quasar is indeed underestimated by as much as 1000 +. these systems still have τρις>Dag their infall velocities are ~2000 and ~3600 ! with respect to adopted emission redshift of the quasar.," Two of these systems are clearly infalling: even if the emission redshift of the quasar is indeed underestimated by as much as 1000 $^{-1}$, these systems still have $z_{\rm abs}>z_{\rm em}$; their infall velocities are $\sim$ 2000 and $\sim$ 3600 $^{-1}$ with respect to adopted emission redshift of the quasar."
+ Arguments presented above suggest that all eight (or nine) systems are physically associated with the quasar., Arguments presented above suggest that all eight (or nine) systems are physically associated with the quasar.
+ If this is the case. then all systems need to be considered as one large associated absorption complex. with dispersion of —4300 !.," If this is the case, then all systems need to be considered as one large associated absorption complex, with dispersion of $\sim$ 4300 $^{-1}$."
+ It appears that the absorption originates in clouds m the immediate vicinity of the quasar., It appears that the absorption originates in clouds in the immediate vicinity of the quasar.
+ The ejection velocities and the velocity dispersion of the complex (both of the order of thousands of +) appear to be too high to interpret the absorbers as originating in the halos of galaxies in the quasar host cluster., The ejection velocities and the velocity dispersion of the complex (both of the order of thousands of $^{-1}$ ) appear to be too high to interpret the absorbers as originating in the halos of galaxies in the quasar host cluster.
+ The combination of high redshift. brightness. richness of the heavy element absorption — and associated absorption in particular — in the spectrum of 1160343820 ts truly unique.," The combination of high redshift, brightness, richness of the heavy element absorption – and associated absorption in particular – in the spectrum of 1603+3820 is truly unique."
+ Eleven confirmed systems in the spectrum of 1160343820 ranks among the richest known., Eleven confirmed systems in the spectrum of 1603+3820 ranks among the richest known.
+ The York catalog (York et citeyor199].. Richards et citeric1999)) contains only 10 other quasars with à greater number of absorbers.," The York catalog (York et \\cite{yor1991}, Richards et \\cite{ric1999}) ) contains only 10 other quasars with a greater number of absorbers."
+ Of these quasars. only two (QO0958+455| and QI2254317) are of comparable brightness. but both are at considerably lower redshifts.," Of these quasars, only two (Q0958+551 and Q1225+317) are of comparable brightness, but both are at considerably lower redshifts."
+ 1160343820 is even more spectacular when associated absorption spectrum is concerned., 1603+3820 is even more spectacular when associated absorption spectrum is concerned.
+ York’s catalog contains only three quasars which have eight or more systems within 8000 ! of the emission redshift (QI037—270. QIS114+09]1 and Q1556+335).," York's catalog contains only three quasars which have eight or more systems within 8000 $^{-1}$ of the emission redshift (Q1037–270, Q1511+091 and Q1556+335)."
+ All three are much fainter: among QSOs brighter than 16 mag no object comes close to 1160343820., All three are much fainter; among QSOs brighter than 16 mag no object comes close to 1603+3820.
+ We can hypothesize that since none of the absorbers in the complex appears to be drastically different than the other ones they all may have been produced by a single event in the quasar’s past. and that they may represent the velocity dispersion of shells of matter ejected during this event.," We can hypothesize that since none of the absorbers in the complex appears to be drastically different than the other ones they all may have been produced by a single event in the quasar's past, and that they may represent the velocity dispersion of shells of matter ejected during this event."
+ On the other hand. the spectrum also contains an absorber which is much stronger than the systems from the complex. as well as two infalling absorbers. which indicate that the environment of 1160343820 is more complicated.," On the other hand, the spectrum also contains an absorber which is much stronger than the systems from the complex, as well as two infalling absorbers, which indicate that the environment of 1603+3820 is more complicated."
+ We stress that this paper presents an analysis based on the observations in relatively low resolution., We stress that this paper presents an analysis based on the observations in relatively low resolution.
+ Our conclusions are by necessity mostly qualitative since the resolution is inadequate. for performing detailed studies of chemical composition or velocity structure of the individual systems., Our conclusions are by necessity mostly qualitative since the resolution is inadequate for performing detailed studies of chemical composition or velocity structure of the individual systems.
+ High resolution studies of the complex will be of special interest. since it is heavily blended in our data and it is very likely that high resolution spectrum will reveal more absorption systems.," High resolution studies of the complex will be of special interest, since it is heavily blended in our data and it is very likely that high resolution spectrum will reveal more absorption systems."
+ 1160343820 ts very bright for a. =2.51 quasar and is therefore an excellent candidate for such observations., 1603+3820 is very bright for a $z=2.51$ quasar and is therefore an excellent candidate for such observations.
+ We would like ο thank T. Alderoft. J. Bechtold. D. Dobrzycka. M. Elvis. S. Mathur. J. Scott. and A. Siemiginowska for helpful discussions. A. Milone for assistance at the MMT. and F. Drake for providing good working environment for this project.," We would like to thank T. Aldcroft, J. Bechtold, D. Dobrzycka, M. Elvis, S. Mathur, J. Scott, and A. Siemiginowska for helpful discussions, A. Milone for assistance at the MMT, and F. Drake for providing good working environment for this project."
+ T. Alderoft and 1. Bechtold wrote computer codes that were used in the analysis., T. Aldcroft and J. Bechtold wrote computer codes that were used in the analysis.
+ A.D. acknowledges support from NASA Contract No., A.D. acknowledges support from NASA Contract No.
+ NAS8-39073 (Chandra X-Ray Observatory Center)., NAS8-39073 (Chandra X-Ray Observatory Center).
+ The Hamburg Quasar Survey is supported by the DFG through grants Re 353/11 and Re 353/22., The Hamburg Quasar Survey is supported by the DFG through grants Re 353/11 and Re 353/22.
+For the 3 hot post-AGB candidates with uceligible circumstellar extinction. IRASI7203-1531. TRASIL7L60-3111 (SAO 209306) and IRASISS279-1707 (LSS 5112). we modeled the spectra using solar iietallicity Iurucz (1991) model atiiosphneres (Fig.,"For the 3 hot post-AGB candidates with negligible circumstellar extinction, IRAS17203-1534, IRAS17460-3114 (SAO 209306) and IRAS18379-1707 (LSS 5112), we modelled the spectra using solar metallicity Kurucz (1994) model atmospheres (Fig."
+ 7) aud derived the effecIve temperatures and eravities of these stars., 7) and derived the effective temperatures and gravities of these stars.
+" The V-baud fluxes of the stars were corrected for extinction assuniug the normal interstellar extiuction aw. ie. Ap, = 3.1 ον V) (seo eg."," The V-band fluxes of the stars were corrected for extinction assuming the normal interstellar extinction law, i.e. $_{v}$ = 3.1 $\times$ $-$ $_{\rm 2200\AA~}$ (see eg."
+ Seaton. 1979).," Seaton, 1979)."
+ The stellar flux cistiibu22004ions normalised to the corrected V-biud flux of cach star were compared with I&urucz models normalised to the respective model's flux at., The stellar flux distributions normalised to the corrected V-band flux of each star were compared with Kurucz models normalised to the respective model's flux at.
+ Loe & in the case of the OSTIT star. IRÀASITT160-3111 (SAO 209206) was estimated to be L0.," Log g in the case of the O8III star, IRAS17460-3114 (SAO 209206) was estimated to be 4.0."
+ This value is uncertain due to the nou-availability of Ίππο models of lower eravitv at the high temperature (Tg =3500018) of the sar., This value is uncertain due to the non-availability of Kurucz models of lower gravity at the high temperature $_{\rm eff}$ =35000K) of the star.
+ The star may also be slightly metal deficient as discussed in Sec., The star may also be slightly metal deficient as discussed in Sec.
+ 3., 3.
+ For the remaining stars. we obtained the effective temperatures and eravities based ou their optical spectral types (Lane. 1992).," For the remaining stars, we obtained the effective temperatures and gravities based on their optical spectral types (Lang, 1992)."
+ Table 5 lists the adopted Tig aud log ο values., Table 5 lists the adopted $_{\rm eff}$ and log g values.
+ The Tor an log ο values estimated frou the optical spectral types of the stars may be uncertain by  dx 1000f aud + 1.0 respectively., The $_{\rm eff}$ and log g values estimated from the optical spectral types of the stars may be uncertain by $\sim$ $\pm$ 1000K and $\pm$ 1.0 respectively.
+ High resolution optical spectra of these stars are required for an accurate cleternunation of T. aud log ooe. The stars were placed on Schónuberber's (1982. 1987) post-AGB evolutionary tracks for core masses (AL) of 0.516. 0.565. 0.598 and O.6LL AL. (Fig.," High resolution optical spectra of these stars are required for an accurate determination of $_{\rm eff}$ and log g. The stars were placed on Schönnberber's (1983, 1987) post-AGB evolutionary tracks for core masses $_{c}$ ) of 0.546, 0.565, 0.598 and 0.644 $_{\odot}$ (Fig."
+ 8)., 8).
+ The far-IR &ux distributions of the hot post-ACD candidates iuust necessarily be due to flux from the iof central stars absorbed aud re-radiated by the cold cimennstellar dust envelopes which are a remnant of nass-loss on the ACB phase of the stir., The far-IR flux distributions of the hot post-AGB candidates must necessarily be due to flux from the hot central stars absorbed and re-radiated by the cold circumstellar dust envelopes which are a remnant of mass-loss on the AGB phase of the star.
+ Heuce. the integrated far-IR fux (Fg) must be comparable to or ess than the iutegrated stellar flux (Ea).," Hence, the integrated far-IR flux $_{\rm fir}$ ) must be comparable to or less than the integrated stellar flux $_{\rm star}$ )."
+" We estimated Fs, from the 124 to 100μ IRAS flux distributious or the stars.", We estimated $_{\rm fir}$ from the $\mu$ to $\mu$ IRAS flux distributions for the stars.
+ The inteerated stellar flux (Fa) frou to was estimated by combining the IUE spectra with the U.D.V magnitudes of the stars from iterature.," The integrated stellar flux $_{\rm star}$ ) from to was estimated by combining the IUE spectra with the U,B,V magnitudes of the stars from literature."
+ The IUE spectra and the U.D.V. magnitudes were corrected for interstellar extinction derived frou he feature.," The IUE spectra and the U,B,V magnitudes were corrected for interstellar extinction derived from the feature."
+ TRASIL7123-1755 was not detected iu the UV., IRAS17423-1755 was not detected in the UV.
+ Its U.BWVRAT maeudes (Cauba oet al..," Its U,B,V,R,I magnitudes (Gauba et al.,"
+ 2003) were corrected for interstellar extinction using the Stauard extinction law (Ricke Lehofsla. 1985) aud the integrated stellar flux from (U-baud) to (I-band) was estimated.," 2003) were corrected for interstellar extinction using the standard extinction law (Rieke Lebofsky, 1985) and the integrated stellar flux from (U-band) to (I-band) was estimated."
+" In the three cases with uceligible circumstellar extinction (IRAS 17203-1531. TRASIL7L60-BILL and IRASIS379-1707). Fe, was conrparable to or significantly less than Foray."," In the three cases with negligible circumstellar extinction (IRAS 17203-1534, IRAS17460-3114 and IRAS18379-1707), $_{\rm fir}$ was comparable to or significantly less than $_{\rm star}$."
+ Iu coutras stars with circtunstellar extinction in the UV. eg. IRASII331-6135.," In contrast stars with circumstellar extinction in the UV, eg. IRAS14331-6435,"
+" IRASL7TO7Lists. IRASLITA3II-1921. IRASITI23-1755. IRAS15062|2110 and IRAS22023|529) showed a hnieh ratio of Fg, Έντα Πιοισαας partial obscuration of the ceutral stars."," IRAS17074-1845, IRAS17311-4924, IRAS17423-1755, IRAS18062+2410 and IRAS22023+5249) showed a high ratio of $_{\rm fir}$ $_{\rm star}$ indicating partial obscuration of the central stars."
+ These stars may have dusty circumstellar clisks., These stars may have dusty circumstellar disks.
+ The wavelength a which the shortward edee of the CIV. (or NV) absorptio1 profile intersects the stellay continu is usually used to estimate the terminal wind velocities (v4) of hot stars from) low resolution IUE spectra., The wavelength at which the shortward edge of the CIV (or NV) absorption profile intersects the stellar continuum is usually used to estimate the terminal wind velocities $_{\infty}$ ) of hot stars from low resolution IUE spectra.
+ Iowever. this vahe ds usually au upper nuit to the true vy.," However, this value is usually an upper limit to the true $_{\infty}$."
+ Alternatively. the difference between tle absorption and emission line ceuters (Av) are used to estimate v4 from low resolution IUE spectra (Prinja. 1991).," Alternatively, the difference between the absorption and emission line centers $\Delta$ v) are used to estimate $_{\infty}$ from low resolution IUE spectra (Prinja, 1994)."
+ But high dispersion studies (see eg., But high dispersion studies (see eg.
+ Perinotto et al.," Perinotto et al.,"
+ 1982) lave shown that Av is ess than vy., 1982) have shown that $\Delta$ v is less than $_{\infty}$.
+ Considerable circumstellar extinction iu 10 of the 15 stars discussed in this paper. has siguificautly distorted the continu fiux distributions aud line iuteusities of the CTV profiles.," Considerable circumstellar extinction in 10 of the 15 stars discussed in this paper, has significantly distorted the continuum flux distributions and line intensities of the CIV profiles."
+ ITeuce. the point where the due absorption edge of the CIV profile iutersects the UV. contiuuua is not reliable.," Hence, the point where the blue absorption edge of the CIV profile intersects the UV continuum is not reliable."
+ These factors also coutribue to the fact that the CTV P-Cveni profiles are not. clearly observe in the low resolution spectra of these stars., These factors also contribute to the fact that the CIV P-Cygni profiles are not clearly observed in the low resolution spectra of these stars.
+ The cussion peaks may ο too weak or lost in the reddened continuum., The emission peaks may be too weak or lost in the reddened continuum.
+ ence it is not possible to estimate νς from the shortward edee of he CIV absorption profile or from the difference between he absorption and cussion line ceuters., Hence it is not possible to estimate $_{\infty}$ from the shortward edge of the CIV absorption profile or from the difference between the absorption and emission line centers.
+ The CIV resonance doublet ) ds rot resolved iu the IUE low resolution spectra.," The CIV resonance doublet, ) is not resolved in the IUE low resolution spectra."
+ Stellar wind παν be assumed to be abseut in a BOW: star., Stellar wind may be assumed to be absent in a B0V star.
+ We ameasured the absorption niüuunaiun of the CIV απο in the UV(IUE) spectra of a BOW standard star. HD36512 1A} from the atlas by Ieck et al. (," We measured the absorption minimum of the CIV feature in the UV(IUE) spectrum of a B0V standard star, HD36512 ) from the atlas by Heck et al. ("
+1981).,1984).
+ The ditterence between this waveleneth aud he absorption minima] of the CIV feaure in the IUE spectra of our ho post-AGD candidaOR Was used to estimate the wind velocities in these stars., The difference between this wavelength and the absorption minima of the CIV feature in the IUE spectra of our hot post-AGB candidates was used to estimate the wind velocities in these stars.
+ Tf iuste. we had used the mean laboratory waveleneth of the CTV resonance doublet ). the esnuated terminal wind velocites would have been ereater bv 1," If instead, we had used the mean laboratory wavelength of the CIV resonance doublet ), the estimated terminal wind velocites would have been greater by $^{-1}$."
+" The NY resonance douJet ) is also ""resolved in low resolution IUE spectra."," The NV resonance doublet, ) is also unresolved in low resolution IUE spectra."
+ Moreover. the NV cature is often contaminated by eeocoroual Lyman a.," Moreover, the NV feature is often contaminated by geocoronal Lyman $\alpha$."
+ From he standard star atlas (Ileck et aL.," From the standard star atlas (Heck et al.,"
+ 1981) we were unable to find a star in which the NV liue is distinguishable from Lyman a aud unaffected. by stellar wind., 1984) we were unable to find a star in which the NV line is distinguishable from Lyman $\alpha$ and unaffected by stellar wind.
+ Deuce. iu he case of IRAS1258LIS37. we estimated the wind velocity frou the difference between the NV absorption nuüninuiuu ae the mean laboratory wavelcneth of the line ).," Hence, in the case of IRAS12584-4837, we estimated the wind velocity from the difference between the NV absorption minimum and the mean laboratory wavelength of the line )."
+ The estimated wind velocitics aro listed in Table 5., The estimated wind velocities are listed in Table 5.
+ No shift was observed iu the case of IRASITOTI-I815. IRASIS062|2110. IRASIS23TO- aud IRAS22195|5131.," No shift was observed in the case of IRAS17074-1845, IRAS18062+2410, IRAS18379-1707 and IRAS22495+5134."
+ CIV (or NV) lines were not observed in IBAST1331-6135 and IRAS16206-5958., CIV (or NV) lines were not observed in IRAS14331-6435 and IRAS16206-5958.
+ The SUV (391A) absorption line in TRASL1331-6135 did not show a waveleneth shift due to stellar wind., The SiIV ) absorption line in IRAS14331-6435 did not show a wavelength shift due to stellar wind.
+ For low, For low
+"large extent, there still exists qualitative differences between the different scenarios.","large extent, there still exists qualitative differences between the different scenarios."
+ This of course is a reflection of the vastly contrasting models of reionization compared in this paper and difference will disappear if the models investigated are more similar to each another., This of course is a reflection of the vastly contrasting models of reionization compared in this paper and difference will disappear if the models investigated are more similar to each another.
+" The previous section provided a qualitative description of the differences in the 21-cm ὁΤι, fluctuations for two cases i.e., with and without the assumption that T;=Τις."," The previous section provided a qualitative description of the differences in the 21-cm $\delta \mathrm{T_b}$ fluctuations for two cases i.e., with and without the assumption that $\mathrm{T_s =
+ T_k}$."
+" In this section, a number of statistical tests are performed to quantify these differences."," In this section, a number of statistical tests are performed to quantify these differences."
+ For this purpose we focus on the scenario with the miniqsos., For this purpose we focus on the scenario with the miniqsos.
+ This exercise can be repeated for the other two scenarios as well., This exercise can be repeated for the other two scenarios as well.
+" Firstly, Fig."," Firstly, Fig."
+" 14 shows the 3-D power spectra of dT, at redshift 12, 10, 8 and 6."," \ref{fig:powspec} shows the 3-D power spectra of $\delta
+\mathrm{T_b}$ at redshift 12, 10, 8 and 6."
+" In our model, reionization begins at around redshift 12.7 and thus z=12 can be considered the onset of reionization."," In our model, reionization begins at around redshift 12.7 and thus z=12 can be considered the onset of reionization."
+" The dashed curve reflects the assumption that T,=Γι while the solid curve takes into account the evolution of T, self-consistently.", The dashed curve reflects the assumption that $\mathrm{T_s} = \mathrm{T_k}$ while the solid curve takes into account the evolution of $\mathrm{T_s}$ self-consistently.
+" At the initial stages of reionization, only regions at close proximity to the source have been ionized and heated, and the volume filling by ""spheres of ionization"" has just begun, leaving large portions of the Universe neutral and cooler than the CMB."," At the initial stages of reionization, only regions at close proximity to the source have been ionized and heated, and the volume filling by “spheres of ionization” has just begun, leaving large portions of the Universe neutral and cooler than the CMB."
+" Here we assume that after the epoch of recombination the Universe cools adiabatically, ie, Tko1/(14z)? and Tompος1/(1+z), which results in Τι<<Town at the onset of reionization."," Here we assume that after the epoch of recombination the Universe cools adiabatically, i.e., $\mathrm{T_k}
+\propto 1/(1+z)^2$ and $\mathrm{T_{CMB}} \propto 1/(1+z)$, which results in $\mathrm{T_k} << \mathrm{T_{CMB}}$ at the onset of reionization."
+" Thus, we have a situation in which we have assumed T,=Ty, with Τι Tomp, resulting in an decoupling of T; and Tomp, and is manifested as increased 21-cm power."," Thus, we have a situation in which we have assumed $\mathrm{T_s} = \mathrm{T_k}$, with $\mathrm{T_k} <<
+\mathrm{T_{CMB}}$ , resulting in an decoupling of $\mathrm{T_s}$ and $\mathrm{T_{CMB}}$, and is manifested as increased 21-cm power."
+" The discrepancy quickly (by z= 10) disappears as the ""spheres of Lya “ overlap to volume-fill the simulation box and hence increase the T, — Τι coupling efficiency.", The discrepancy quickly (by $z = 10$ ) disappears as the “spheres of $\alpha$ “ overlap to volume-fill the simulation box and hence increase the $\mathrm{T_s}$ — $\mathrm{T_k}$ coupling efficiency.
+" Therefore, in our model, it is safe to make the assumption that Τς follows Τι after about z=10."," Therefore, in our model, it is safe to make the assumption that $\mathrm{T_s}$ follows $\mathrm{T_k}$ after about $z = 10$."
+ From Eq., From Eq.
+" 6 it follows that dT), attains large negative values when two conditions are satisfied, Ts;<<Tcomp and a high neutral hydrogen overdensity."," \ref{eq:dtb} it follows that $\delta \mathrm{T_b}$ attains large negative values when two conditions are satisfied, $\mathrm{T_s} << \mathrm{T_{CMB}}$ and a high neutral hydrogen overdensity."
+" We therefore plot the mean, minimum and maximum values of dT, as a function of redshift in Fig. 15.."," We therefore plot the mean, minimum and maximum values of $\delta \mathrm{T_b}$ as a function of redshift in Fig. \ref{fig:dtbmmm}."
+" As expected the mean and minimum values of JT} are much lower for the case in which the spin temperature has been set to the kinetic temperature, because the Universe has cooled substantially below the CMB and coupling Τι everywhere in the Universe to T'; enforces the conditions required for large negative dT» as stated above."," As expected the mean and minimum values of $\delta \mathrm{T_b}$ are much lower for the case in which the spin temperature has been set to the kinetic temperature, because the Universe has cooled substantially below the CMB and coupling $\mathrm{T_k}$ everywhere in the Universe to $\mathrm{T_s}$ enforces the conditions required for large negative $\delta \mathrm{T_b}$ as stated above."
+" The maximum positive values though do not differ because a positive dT, indicates that T;>Tcmp and hence has to be heated by the radiating source."," The maximum positive values though do not differ because a positive $\delta \mathrm{T_b}$ indicates that $\mathrm{T_s}
+> \mathrm{T_{CMB}}$ and hence has to be heated by the radiating source."
+" And in regions where the source has heated the IGM, the presence of Lyo flux and colliosional coupling almost always guarantees the coupling of Ts to Γι. anyway."," And in regions where the source has heated the IGM, the presence of $\alpha$ flux and colliosional coupling almost always guarantees the coupling of $\mathrm{T_s}$ to $\mathrm{T_k}$ anyway."
+" As another simple diagnostic of the efficiency of T, Τι coupling we plot, (i) the differential temperature, i.e., Tcms/Ts, (ii) their ratio, T;/Ts, as a function of redshift (Fig. 16))."," As another simple diagnostic of the efficiency of $\mathrm{T_s}$ --- $\mathrm{T_k}$ coupling we plot, (i) the differential temperature, i.e., $1 - \mathrm{T_{CMB}/T_{s}}$ , (ii) their ratio, $\mathrm{T_k}/\mathrm{T_s}$, as a function of redshift (Fig. \ref{fig:meantstk}) )."
+" Note here that if the spin temperature was artificially set to the kinetic temperature, the differential temperature would be saturated at unity."," Note here that if the spin temperature was artificially set to the kinetic temperature, the differential temperature would be saturated at unity."
+ But the coupling ensures that the spin temperature remains below that of the CMB till about a redshift of 10., But the coupling ensures that the spin temperature remains below that of the CMB till about a redshift of 10.
+" Although significantly different at the early phases of reionization, in the"," Although significantly different at the early phases of reionization, in the"
+surface brightness maps and to construct total radio luminosities by integrating these quantities across the entire projected cluster surface out to the virial radius.,surface brightness maps and to construct total radio luminosities by integrating these quantities across the entire projected cluster surface out to the virial radius.
+" The integrated projected density is of course M, and we will designate the integrated turbulent pressure as Ty=miv?.", The integrated projected density is of course $M_v$ and we will designate the integrated turbulent pressure as $\Gamma_v = \sum_{i} m_i v_i^2$.
+" Therefore0, we can create simulated radio surface brightness maps for our clusters by normalizing maps of projected turbulent pressure and projected mass using assumed radio luminosities and rest-frame spectra.", Therefore we can create simulated radio surface brightness maps for our clusters by normalizing maps of projected turbulent pressure and projected mass using assumed radio luminosities and rest-frame spectra.
+" Because we may not fully resolve intracluster turbulence, the total amount of turbulent pressure in our clusters may be lower than the ~10% of hydrostatic pressure seen in high-resolution simulations (???).."," Because we may not fully resolve intracluster turbulence, the total amount of turbulent pressure in our clusters may be lower than the $\sim 10\%$ of hydrostatic pressure seen in high-resolution simulations \citep{Ricker2001,Ritchie2002,Mitchell2009}."
+" However, because the normalization of the radio power is supplied independently (see below), all we require of the turbulent pressure maps is that they be sufficiently diffuse and representative in spatial extent of clusters containing radio halos."," However, because the normalization of the radio power is supplied independently (see below), all we require of the turbulent pressure maps is that they be sufficiently diffuse and representative in spatial extent of clusters containing radio halos."
+ The detailed structure of the maps should not be regarded as realistic., The detailed structure of the maps should not be regarded as realistic.
+" Because of the beam smearing described below, this fact does not significantly affect our analysis."," Because of the beam smearing described below, this fact does not significantly affect our analysis."
+" Also, while we cannot depend on these simulated clusters to provide correct high-resolution X-ray and radio surface brightness maps, we can still use them to identify broad features, such large-scale shocks and the relative radial dependence of turbulence."," Also, while we cannot depend on these simulated clusters to provide correct high-resolution X-ray and radio surface brightness maps, we can still use them to identify broad features, such large-scale shocks and the relative radial dependence of turbulence."
+" Using these integrated quantities, we construct a rest frame 1.4 GHz radio power via where C, is a scaling constant, M, is the virial mass, T', is the virial turbulent pressure, and D, is the magnetic field parameter: where B(M,)=(B)(M,/(M))° and Bomp=3.2(1+2)3wG is the equivalent magnetic field strength of the cosmic microwave background."," Using these integrated quantities, we construct a rest frame $1.4$ GHz radio power via where $C_s$ is a scaling constant, $M_v$ is the virial mass, $\Gamma_v$ is the virial turbulent pressure, and $B_s$ is the magnetic field parameter: where $B(M_v) \equiv \aveb (M_v/\langle M \rangle)^b$ and $B_{\rm CMB} \equiv 3.2(1+z)^2 \mg$ is the equivalent magnetic field strength of the cosmic microwave background."
+" This formulation separates physical processes that generate CRs (M, and D,) from those that contribute to radio emission (B in the numerator) and CR losses due to emission (B? in the denominator) and inverse Compton scattering (Bcmp).", This formulation separates physical processes that generate CRs $M_v$ and $\Gamma_v$ ) from those that contribute to radio emission $B$ in the numerator) and CR losses due to emission $B^2$ in the denominator) and inverse Compton scattering $B_{\rm CMB}$ ).
+ The losses enter into this equation because they limit the maximum CR energy., The losses enter into this equation because they limit the maximum CR energy.
+" In this formalism, M, measures the total cluster mass and thus should scale with the dependence of CR generation on hadronic secondary processes, whereas I, measures the total cluster turbulence and thus should provide a measure of the reacceleration of CR. electrons by that turbulence."," In this formalism, $M_v$ measures the total cluster mass and thus should scale with the dependence of CR generation on hadronic secondary processes, whereas $\Gamma_v$ measures the total cluster turbulence and thus should provide a measure of the reacceleration of CR electrons by that turbulence."
+ We willset (M)=1.5x101°Mo., We willset $\langle M \rangle = 1.5 \times 10^{15} \msol$.
+" There are thus five independent parameters: the average magnetic field (B), the scaling of magnetic field with cluster mass, b, the scaling of radio power with virial mass, a, the scaling of radio power with turbulent pressure, c, and an overall scaling parameter C,."," There are thus five independent parameters: the average magnetic field $\aveb$, the scaling of magnetic field with cluster mass, $b$ , the scaling of radio power with virial mass, $a$, the scaling of radio power with turbulent pressure, $c$, and an overall scaling parameter $C_s$."
+ A summary of our model parameters is given in1., A summary of our model parameters is given in.
+. 'This model is a generalization of the one derived in ?.., This model is a generalization of the one derived in \citet{Cassano2005a}.
+" Note that the analysis of hadronic secondary models of ? identified the functional form of the magnetic field as B=B(M,)?/(B(M,)?+ i.e., the denominator is not squared."," Note that the analysis of hadronic secondary models of \citet{Dolag2000} identified the functional form of the magnetic field as $B
+= B(M_v)^2 / (B(M_v)^2 + B_{\rm CMB}^2)$: i.e., the denominator is not squared."
+" Our model Bà):easily accommodates this scenario: when Bowp dominates, this will appear as a constant factor folded into C,, and when B(M,) dominates this will simply adjust the mass scaling factor a."," Our model easily accommodates this scenario: when $B_{\rm CMB}$ dominates, this will appear as a constant factor folded into $C_s$, and when $B(M_v)$ dominates this will simply adjust the mass scaling factor $a$."
+" Provided that the necessary adjustments to C, and a are made, our model holds even for intermediate cases where B~Bomp, since we are largely in the regime where M,«(M)."," Provided that the necessary adjustments to $C_s$ and $a$ are made, our model holds even for intermediate cases where $B \sim B_{\rm CMB}$, since we are largely in the regime where $M_v < \langle M \rangle$."
+ We are fixing the form of the magnetic field dependence since the radio synchrotron power will always depend on magnetic field pressure (B?) independently of the CR generation and acceleration mechanisms (see ? for a discussion)., We are fixing the form of the magnetic field dependence since the radio synchrotron power will always depend on magnetic field pressure $B^2$ ) independently of the CR generation and acceleration mechanisms (see \citet{Cassano2005a} for a discussion).
+" We stress that although this model is relatively simple, it allows us to explore a range of plausible acceleration mechanisms and examine relative changes to luminosity functions, scaling relations, and other radio properties."," We stress that although this model is relatively simple, it allows us to explore a range of plausible acceleration mechanisms and examine relative changes to luminosity functions, scaling relations, and other radio properties."
+ This model allows us to explore both CR generation mechanisms or a mixture of both., This model allows us to explore both CR generation mechanisms or a mixture of both.
+" For example, the hadronic secondaries model should predict radio power which scales with cluster mass, so c=0 in this case."," For example, the hadronic secondaries model should predict radio power which scales with cluster mass, so $c=0$ in this case."
+" A reacceleration model is proportional to turbulent pressure, so a would be 0."," A reacceleration model is proportional to turbulent pressure, so $a$ would be $0$."
+" Note that the model of CBS06 is based on reacceleration, but only scales with cluster mass."," Note that the model of CBS06 is based on reacceleration, but only scales with cluster mass."
+" This is because I'; roughly scales with M, with a logarithmic slope of 1.7, as shown in2."," This is because $\Gamma_v$ roughly scales with $M_v$ with a logarithmic slope of $1.7$, as shown in."
+". Note that our scaling relation is roughly consistent with those derived from other simulations (e.g.,?).."," Note that our scaling relation is roughly consistent with those derived from other simulations \citep[e.g.,][]{Vazza2006}."
+" There is some scatter in this relation due to the merger history of a particular cluster: recent mergers produce stronger levels of turbulence, which tend to scatter the cluster higher in this relation."," There is some scatter in this relation due to the merger history of a particular cluster: recent mergers produce stronger levels of turbulence, which tend to scatter the cluster higher in this relation."
+ Note that the model of CBS06 corresponds here to a=4/3 and c— 0., Note that the model of CBS06 corresponds here to $a=4/3$ and $c=0$ .
+" A degeneracy existsfor calculations of total radio luminosity between models that scale with turbulent pressure and those that scale with mass, since we may freely exchange c for 1.7a and vice-versa."," A degeneracy existsfor calculations of total radio luminosity between models that scale with turbulent pressure and those that scale with mass, since we may freely exchange $c$ for $1.7a$ and vice-versa."
+" However,"," However,"
+The Meijer G-function is a very general. analytical function introduced by ? which includes most of the special functions as specific cases.,"The Meijer G-function is a very general, analytical function introduced by \citet{1936NAvW.18..10} which includes most of the special functions as specific cases."
+ It is defined in terms of the inverse Mellin transform (?) by: where C is à contour in the complex plane. IL(5) is theGamma function and a and b are vectors of dimension p and q. respectively.," It is defined in terms of the inverse Mellin transform \citep{9780070195493} by: where $\mathit{C}$ is a contour in the complex plane, $\Gamma(s)$ is theGamma function and $\mathbf{a}$ and $\mathbf{b}$ are vectors of dimension $p$ and $q$, respectively."
+ The basic properties of the Meter G-function are too numerous to be mention here., The basic properties of the Meijer G-function are too numerous to be mention here.
+ We only provide a short list of the most relevant properties for this work., We only provide a short list of the most relevant properties for this work.
+ A Meyer G-function with p>q can be transformed to another G-function with p<q: Other property is that 1f one the parameters of a and b appears in both the numerator and denominator of the integrand. the order of the Meijer G-function may decrease and the fraction simplified.," A Meijer G-function with $p>q$ can be transformed to another G-function with $p<q$ : Other property is that if one the parameters of $\mathbf{a}$ and $\mathbf{b}$ appears in both the numerator and denominator of the integrand, the order of the Meijer G-function may decrease and the fraction simplified."
+ The positions of the parameters dictates which order a or 2 will decrease., The positions of the parameters dictates which order $m$ or $n$ will decrease.
+" For example ifa,=bj for some &=1.2....n and j2m l.m+ 2.....g. the orders p.g and à of the Meijer G-function will decrease: In the other case if a;=b; forsome k=ΗΕ1.n-2...p and j= 1.2...../. the orders p.g and mof the Metjer G-function will decrease: The order reduction formulafor the Meijer) G-function 15:"," For example if $a_{k}=b_{j}$ for some $k=1,2,...,n$ and $j=m+1,m+2,...,q$ , the orders $p,q$ and $n$ of the Meijer G-function will decrease: In the other case if $a_{k}=b_{j}$ for some $k=n+1,n+2,...,p$ and $j=1,2,...,m$ , the orders $p,q$ and $m$of the Meijer G-function will decrease: The order reduction formulafor the Meijer G-function is:"
+We have analyzed the kinematic properties of all S remnants.,We have analyzed the kinematic properties of all 8 remnants.
+ We computed the rotation curves of the old. anc voung populations. for all the remnants.," We computed the rotation curves of the old and young populations, for all the remnants."
+ The top right. panels of Figs. 6-, The top right panels of Figs. \ref{M12_kin}-
+ 9. show some examples of these curves., \ref{M13_kin} show some examples of these curves.
+ For simulations M12. Ml12orb. MI290. rMI2. MIS. and MIIO. the old. populations have a lower mean rotational velocity than their vounger counterparts.," For simulations M12, M12orb, M1290, rM12, M13, and M110, the old populations have a lower mean rotational velocity than their younger counterparts."
+ The dillerence is about 200 300kms+ for MI2. MI20rb. ancl MI290 (mass ratio 2-1). 150kms.+ for MIA (mass ratio 3:1) and 50kms.+ for MIIO (mass ratio 10:1).," The difference is about $200-300\,{\rm km\,s^{-1}}$ for M12, M12orb, and M1290 (mass ratio 2:1), $150\,{\rm km\,s^{-1}}$ for M13 (mass ratio 3:1) and $50\,{\rm km\,s^{-1}}$ for M110 (mass ratio 10:1)."
+ Hence. it decreases with increasing miss ratio of the progenitors.," Hence, it decreases with increasing mass ratio of the progenitors."
+ Comparing MI2 with rMI2. we also find that reversing the rotation of CGal2 (so that," Comparing M12 with rM12, we also find that reversing the rotation of Gal2 (so that"
+We now briefly consider Pop.,We now briefly consider Pop.
+l massive stars of inilial masses in the range 16—64M... typically. SN progenitors.,"I massive stars of initial masses in the range $16-64\ \Msun$, typically, SN progenitors."
+ Since nucleosvnthesis caleulations are limited in our code. we cannot follow the evolution all the wav to the collapse of an iron core.," Since nucleosynthesis calculations are limited in our code, we cannot follow the evolution all the way to the collapse of an iron core."
+ Llowever. we come quite close to it.," However, we come quite close to it."
+ These massive stus go through. advanced nuclear burning stages. until a core composed of the enc-product of our nuclear reactions network is obtained.," These massive stars go through advanced nuclear burning stages, until a core composed of the end-product of our nuclear reactions network is obtained."
+ Core Masses range monotonically [rom 2.44. for the GLA. initial mass and 1.7. for the 16M. initial mass., Core masses range monotonically from $\Msun$ for the $\Msun$ initial mass and $\Msun$ for the $\Msun$ initial mass.
+ The core is enveloped by lavers of different composition. the outermost being predominantly helium.," The core is enveloped by layers of different composition, the outermost being predominantly helium."
+ Envelope masses depend stronelv on the mass loss law assume., Envelope masses depend strongly on the mass loss law assumed.
+ The core contracts. becoming degenerate and uistable. since its mass exceeds the Chancirasekhar limit.," The core contracts, becoming degenerate and unstable, since its mass exceeds the Chandrasekhar limit."
+" As contraction accelerates. temperatures rise (oa [ew 10! I, where electron-positron pairs are created. which enhances (he instabilitv. lowering the adiabatic exponent."," As contraction accelerates, temperatures rise to a few $10^{10}$ K, where electron-positron pairs are created, which enhances the instability, lowering the adiabatic exponent."
+ Pair production replaces iron photodisintegration as the mechanism leading (o core collapse., Pair production replaces iron photodisintegration as the mechanism leading to core collapse.
+ Density profiles throughout the stars are shown in Fig. 12.., Density profiles throughout the stars are shown in Fig. \ref{fig:rhoprof}.
+ The code crashes when the collapse approaches free-fall. with the adiabatic exponent verv close to 4/3 throughout the core.," The code crashes when the collapse approaches free-fall, with the adiabatic exponent very close to 4/3 throughout the core."
+ Since (his point is somewhat arbitrary. (he curves representing stars of cillerent initial masses do not exhibit a perlectly regular (monotonic) behaviour: (his is sometimes (the case for evolutionary tracks or characterises of massive stars in the late stages (?.. 2)). resulting from the complexity of the processes (aking part in Chem. and the related parameters ancl thresholds.," Since this point is somewhat arbitrary, the curves representing stars of different initial masses do not exhibit a perfectly regular (monotonic) behaviour; this is sometimes the case for evolutionary tracks or characteristics of massive stars in the late stages \citealt{1996snai.book.....A}, \citealt{2008ApJ...673.1014U}) ), resulting from the complexity of the processes taking part in them, and the related parameters and thresholds."
+ We do not claim that these calculations shed lisht on pre-supernova evolution: rather. we mention them here as an example of the robustness of the code. which is capable οἱ dealing with complex processes under critical conditions without Iailing.," We do not claim that these calculations shed light on pre-supernova evolution; rather, we mention them here as an example of the robustness of the code, which is capable of dealing with complex processes under critical conditions without failing."
+ Finally. adding (he results obtained for lower masses of Pop.," Finally, adding the results obtained for lower masses of Pop."
+lI. described in Section 4.3.. we show in Fig.,"I, described in Section \ref{canon}, we show in Fig."
+ 13. evolutionary tracks of the stellar central points in the (log7.logp) plane. exhibiting the branching off between stars that end their lives as WDs. and stars that go through advanced nuclear burning stages. ending their lives in dvnamie core collapse.," \ref{fig:rhoctc} evolutionary tracks of the stellar central points in the $(\log T,\log\rho)$ plane, exhibiting the branching off between stars that end their lives as WDs, and stars that go through advanced nuclear burning stages, ending their lives in dynamic core collapse."
+ The term ‘non-canonical refers to stars of unusual internal structure and composition., The term `non-canonical' refers to stars of unusual internal structure and composition.
+ Such configurations may result [rom stellar mergers. where the merging stars may be AIS stars. ejants. compact stars or any combination of different tvpes.," Such configurations may result from stellar mergers, where the merging stars may be MS stars, giants, compact stars or any combination of different types."
+ Stellar mergers are probably, Stellar mergers are probably
+in. they displace the dark matter from the centre.,"in, they displace the dark matter from the centre."
+ Dark matter particles move out of the nucleus., Dark matter particles move out of the nucleus.
+ The density of the dark matter in the centre drops by more than an order of magnitude. an effect that does not happen in cored haloes.," The density of the dark matter in the centre drops by more than an order of magnitude, an effect that does not happen in cored haloes."
+ We note that several authors have recently. studied: this »ocess in detail., We note that several authors have recently studied this process in detail.
+ For example study the change in an initial cuspy density. profile due to he frictional effects of sinking objects., For example study the change in an initial cuspy density profile due to the frictional effects of sinking objects.
+ This process allows initially cusped density profiles to be transformed into nearly iumonic cores as the globulars fall in., This process allows initially cusped density profiles to be transformed into nearly harmonic cores as the globulars fall in.
+ For the case of the five elobulars in the Fornax halo. the maximum initial central density slope for which this can occur is approximately 0.5.," For the case of the five globulars in the Fornax halo, the maximum initial central density slope for which this can occur is approximately 0.5."
+" ""hus in order for these data to be consistent with a CDAL halo. the central cusp must have been modified and Lattened to a slope shallower than this value."," Thus in order for these data to be consistent with a CDM halo, the central cusp must have been modified and flattened to a slope shallower than this value."
+ Exotic scenarios which might achieve this include infall and subsequent blowout of massive gas clouds or star clusters 2005)., Exotic scenarios which might achieve this include infall and subsequent blowout of massive gas clouds or star clusters .
+. The Fornax dwarf spheroidal has five globular clusters orbiting at à projected radius of  Lkkpe from its centre., The Fornax dwarf spheroidal has five globular clusters orbiting at a projected radius of $\sim$ kpc from its centre.
+ Using a cuspy CDM potential with central slope steeper than 1.5 and normalised to match the inferred properties from the kinematical data. we find that these elobulars would all sink o the halo centre within 5 Cvrs.," Using a cuspy CDM potential with central slope steeper than 0.5 and normalised to match the inferred properties from the kinematical data, we find that these globulars would all sink to the halo centre within 5 Gyrs."
+ ὃν contrast. we showed that if Fornax has a constant densitv central core then the dynamical [friction time comes. arbitrarily long - the globulars stop sinking at he edee of the core. thus the present. position of theinmnermmosh elobular gives a lower limit of the core radius of the dark matter distribution.," By contrast, we showed that if Fornax has a constant density central core then the dynamical friction time becomes arbitrarily long - the globulars stop sinking at the edge of the core, thus the present position of the globular gives a lower limit of the core radius of the dark matter distribution."
+ Since CDM uniquely predicts that all haloes are cusped. this suggests that the dark matter distribution within Fornax is not cold. but may be warm dark matter or some other candidate.," Since CDM uniquely predicts that all haloes are cusped, this suggests that the dark matter distribution within Fornax is not cold, but may be warm dark matter or some other candidate."
+ Alternatively. the mass distribution has been modified through some exotic dynamical phenomenon such as rapid mass loss2005).," Alternatively, the mass distribution has been modified through some exotic dynamical phenomenon such as rapid mass loss."
+". If the phase space density. of the core. measured as Q~LO 7M. ""(kms 4) 7 in our model. is primordial then it implies a warm clark matter particle of mass 0.5 kkeV2001).. This is consistent with that required to largely solve the substructure problem19992)."," If the phase space density of the core, measured as $Q\sim 10^{-5}$ $_\odot$ $^{-3}$ $^{-1}$ $^{-3}$ in our model, is primordial then it implies a warm dark matter particle of mass $\sim 0.5$ keV. This is consistent with that required to largely solve the substructure problem."
+.. Although the dwarf spheroidals around the Milkv. Way do not contain prominent nuclei. about 30 percent of dwarl spheroidals (dks) in clusters are nucleated.," Although the dwarf spheroidals around the Milky Way do not contain prominent nuclei, about 30 percent of dwarf spheroidals (dE's) in clusters are nucleated."
+ LH these nuclei form by the merging of star clusters then we must conclude that these galaxies have euspy mass distributions., If these nuclei form by the merging of star clusters then we must conclude that these galaxies have cuspy mass distributions.
+ This could. be due to the fact that transformation to dl via galaxy harassment gives an exponential distribution of stars which usually dominate the central mass cüstribution., This could be due to the fact that transformation to dE via galaxy harassment gives an exponential distribution of stars which usually dominate the central mass distribution.
+ Therefore elobulars could sink via friction against the stellar |ackeround., Therefore globulars could sink via friction against the stellar background.
+ Most nucleated cles are near the cluster centre where harassment is important which supports this idea., Most nucleated dE's are near the cluster centre where harassment is important which supports this idea.
+ Our numerical simulations show that the standard Chandreskhar estimate of sinking timescales works well for cuspy cores but fails completely for cored mass distributions., Our numerical simulations show that the standard Chandreskhar estimate of sinking timescales works well for cuspy cores but fails completely for cored mass distributions.
+" In addition we showed that over LO"" particles are required within the core region to suppress heating from. particle noise1998).", In addition we showed that over $10^6$ particles are required within the core region to suppress heating from particle noise.
+. This is particularly important. in a cored mass distribution where the potential minimum is quite shallow., This is particularly important in a cored mass distribution where the potential minimum is quite shallow.
+ The fact that cored mass distributions do not give rise to dynamical friction is an important result., The fact that cored mass distributions do not give rise to dynamical friction is an important result.
+ We, We
+reconstruction techuique described iu Ἱνταιις&Iilleu-brand(2007b).. which iteratively fits a template PSF to the primary and then subtracts the secondary to fit an nuproved estimate of the primary.,"reconstruction technique described in \citet{Kraus07b}, which iteratively fits a template PSF to the primary and then subtracts the secondary to fit an improved estimate of the primary."
+ This routine was nmaplemieuted using the ALLSTAR routine in the IRAF package DAOPITOT (Stetson 1987). and based on previous experience with NIRC2 imaging. we chose to model the PSF using a Gaussian core with Lorentzian wines.," This routine was implemented using the ALLSTAR routine in the IRAF package DAOPHOT (Stetson 1987), and based on previous experience with NIRC2 imaging, we chose to model the PSF using a Gaussian core with Lorentzian wings."
+ This method is illustrated in Figure 3.., This method is illustrated in Figure \ref{figIm}.
+ In addition to this method. we used a recursive-phase aleoritli (6.8Lolunaunetal.1983). to recoustruct Fourier Phase from Llispeetrum phase. aud then fitted a binary model to the reconstructed phase of the inage.," In addition to this method, we used a recursive-phase algorithm \citep[e.g][]{Lohmann83} to reconstruct Fourier Phase from bispectrum phase, and then fitted a binary model to the reconstructed phase of the image."
+ Tn Table 1.. we sununarize the relative astrometry aud photometry that we measured from our direct dmaeing observations.," In Table \ref{tabImaging}, we summarize the relative astrometry and photometry that we measured from our direct imaging observations."
+ The consistencv in separation aud position anele at the level iu IT aud K bands aud the level iu J baud despite very different. PSFs and diffraction effects demoustrates the reliability of our results., The consistency in separation and position angle at the level in H and K bands and the level in J band despite very different PSFs and diffraction effects demonstrates the reliability of our results.
+ However. the suall reported statistical exrors in Table 1 are not eutirelv consistent internally. requiring additional svstematic errors to be added.," However, the small reported statistical errors in Table \ref{tabImaging} are not entirely consistent internally, requiring additional systematic errors to be added."
+ Therefore. we assign a weighted mean solution of separation 53.6+0.5 1uumas and a position angele of 306.140.6 ddeerces.," Therefore, we assign a weighted mean solution of separation $53.6\pm0.5$ mas and a position angle of $306.4\pm0.6$ degrees."
+ The coutrast ratio in J. IT and Iv is consistent with there beiug no significaut difference iu reddening between the coupoucuts.," The contrast ratio in J, H and K is consistent with there being no significant difference in reddening between the components."
+ Both compoucuts have near infrared colors that are consistent with other vouug Taurus icubers of spectral type  MI-M2. sueeesting that they have similar cluperatures aud that neither has a significant A band excess due to a hot iunuer disk.," Both components have near infrared colors that are consistent with other young Taurus members of spectral type $\sim$ M1-M2, suggesting that they have similar temperatures and that neither has a significant $K$ band excess due to a hot inner disk."
+ In order to model the stars aud the disk. it is iniportaut to know the bolometric ininositv of the svsteni. and hence the reddening.," In order to model the stars and the disk, it is important to know the bolometric luminosity of the system, and hence the reddening."
+" To estimate the reddening. we used the following data: intrinsic colors reported i Dessell&Brett(1988) and Dessell(1990) interpolating between eiauts (lose(g)~ 1) and dsvarfs (loe(g) 5) at au assined log(g) of 3.7. JI. photometry from 2MASS. V-baud photometry averaged tween Cohen&Ίνα(1979). and Hausonetal.(2001) with an assumed πας overall uncertainty. aud au ΑΠΟ spectral classification frou, Collen&Iulia(1979) with an assumed d sub-class uncertainty."," To estimate the reddening, we used the following data: intrinsic colors reported in \citet{Bessell88} and \citet{Bessell90} interpolating between giants $(g) \sim 1$ ) and dwarfs $(g) \sim 5$ ) at an assumed $(g)$ of 3.7, JHK photometry from 2MASS, V-band photometry averaged between \citet{Cohen79} and \citet{Hanson04} with an assumed mag overall uncertainty, and an M1.5 spectral classification from \citet{Cohen79} with an assumed 1 sub-class uncertainty."
+ We arrive at Ey—2.3d02 aud E;gy=0.58+0.0L., We arrive at $E_{V-K}=2.3\pm0.3$ and $E_{J-K}=0.58\pm0.04$.
+ The ratio Eyκy=LUd0.6 is significantly lower han the 5.6 typical of Taurus (Whittetetal.2001).. sugecsting a reddening law more typical of a deuse core.," The ratio $E_{V-K}/E_{J-K} =
+4.0\pm0.6$ is significantly lower than the 5.6 typical of Taurus \citep{Whittet01}, suggesting a reddening law more typical of a dense core."
+ Based on £yyg. the reddening ly is 2.520.3 using the relationship zdzmLALLYκ from Whittetetal.(2001). ot given the uncertainties in the reddening law. ly=3 as used by D'Alessioetal.(2005) is certainly consistent.," Based on $E_{V-K}$, the reddening $A_V$ is $2.5\pm0.3$ using the relationship $A_V \approx 1.1 E_{V-K}$ from \citet{Whittet01}, but given the uncertainties in the reddening law, $A_V=3$ as used by \citet{DAlessio05} is certainly consistent."
+ Absolute stellar properties like masses can be difficult o estimate because pre-inaiu-sequence evolutionary uodels are uot welbcalibrated. but based onu the sescription described in Άννας&IDilleubrand(2007 th. components are κο to have masses of ~0.A46 ALL.," Absolute stellar properties like masses can be difficult to estimate because pre-main-sequence evolutionary models are not well-calibrated, but based on the prescription described in \citet{Kraus07a}, both components are likely to have masses of $\sim$ 0.5-0.6 $M_{\sun}$."
+ Using a L1l15ppe distance to Taurus. the system has an absolute IS inaguitude of 2.55. which iuplies absolute I& magnitudes of 3.2 and 3.1 for the componcuts.," Using a pc distance to Taurus, the system has an absolute K magnitude of 2.55, which implies absolute K magnitudes of 3.2 and 3.4 for the components."
+ According to the models of Baraffectal. (1998).. assunine au effective temperature of ~3600 Να applicable to MI.5 stars (πα1999).. the stars are LAITY? old aud are of mass —0.52 and 40.62 ML...," According to the models of \citet{Baraffe98}, assuming an effective temperature of $\sim$ K applicable to M1.5 stars \citep{Luhman99}, the stars are MYr old and are of mass $\sim$ 0.52 and $\sim$ $M_\sun$."
+ If the effective temperature is closer to IIs. the stars would be of sliehtlv lower mass and <3 MM old. cousistent with the canonical age of Taurus.," If the effective temperature is closer to K, the stars would be of slightly lower mass and $<3$ Myr old, consistent with the canonical age of Taurus."
+ The observational scatter in the Taurus IIR diagram (c.gLuhuiusn2001) ids consistent with Colu Tau/1 being either sinele as assumed by previous authors or an equal brightness double., The observational scatter in the Taurus HR diagram \citep[e.g][]{Luhman04} is consistent with CoKu Tau/4 being either single as assumed by previous authors or an equal brightness double.
+ The mass ratio. however. should be more robust than the masses or age (as it is only weakly dependent ou reddening). so we cau assert with ereater confidence that Ap/A4=O.8540.05.," The mass ratio, however, should be more robust than the masses or age (as it is only weakly dependent on reddening), so we can assert with greater confidence that $M_B/M_A=$ $\pm$ 0.05."
+ The nini possible semi-nmuajor axis (6) for the orbit is the apparent separation divided by 6). which is το ΔΑΙ: based on au assumed distance of Lippe for Taurus. giving an orbital period 10 vears.," The minimum possible semi-major axis $a$ ) for the orbit is the apparent separation divided by $e$ ), which is $e$ AU based on an assumed distance of pc for Taurus, giving an orbital period $>$ 10 years."
+ Based ou considerations of dvuamical disk tzuucation iu Section 3.2.. it is likely that the system was observed near anim Clougation.," Based on considerations of dynamical disk truncation in Section \ref{sectModel}, it is likely that the system was observed near maximum elongation."
+ We expect the inner radius of the Colww Tau/1 disk ο be set by tidal trumeation., We expect the inner radius of the CoKu Tau/4 disk to be set by tidal truncation.
+ To simplify our discussion. we will asstune that the disk is roughly co-planar with he binary auc (as described in Section 3.1)) that the ünarv coniponeuts have similar masses.," To simplify our discussion, we will assume that the disk is roughly co-planar with the binary and (as described in Section \ref{sectProperties}) ) that the binary components have similar masses."
+ In this geometry. Artvinowicz&Lubow(1991) calculate that the 3:1 orbital resonance at Ry~2.080 is opened for all mit immprobably simall ecceutricities (ο20.03).," In this geometry, \citet{Artymowicz94} calculate that the 3:1 orbital resonance at $R_D\sim2.08a$ is opened for all but improbably small eccentricities $e\ga0.03$ )."
+ Beust&Dutrev(2005). predict a gap size of >2.60 for eccentrieifies 2O.1 for CeCe Tau. which has a similar nass ratio to Col&u Tau/1 (aud find this not to be large enough for that system).," \citet{Beust05} predict a gap size of $>2.6a$ for eccentricities $>0.1$ for GG Tau, which has a similar mass ratio to CoKu Tau/4 (and find this not to be large enough for that system)."
+ For very large eccentricities (6m 0.6). larger resonances open and predicted gap sizes Toni Artviuowicz&Lubow(1991) are in excess of Ap—3a but are viscosity-depeudent.," For very large eccentricities $e\ga0.6$ ), larger resonances open and predicted gap sizes from \citet{Artymowicz94} are in excess of $R_D\sim3a$ but are viscosity-dependent."
+ The closer near-equal spectroscopic binary ΠΟ 98800D. (Akesonetal.2007) Is an example of a svstem with high eccentricity that should have opened lage resonances iu the disk: but the radiative transfer modeling has too mauy uncertainties to clearly sav which resonance has opened., The closer near-equal spectroscopic binary HD 98800B \citep{Akeson07} is an example of a system with high eccentricity that should have opened large resonances in the disk: but the radiative transfer modeling has too many uncertainties to clearly say which resonance has opened.
+ Based oon the imünimüuuni semicnajor axis of L6) AAU from Section 2.. we arrive at nimium disk truucation radius of ~16 AAU if Colku Tau/ lis a ow (6~ 0.1) eccentricity binary. or 13 AAU if Colku Taufl is a high (οz 0.6) ecceutricity binary.," Based on the minimum semi-major axis of $e$ AU from Section \ref{sectObservations}, we arrive at minimum disk truncation radius of $\sim$ AU if CoKu Tau/4 is a low $e\sim0.1$ ) eccentricity binary, or $\sim$ AU if CoKu Tau/4 is a high $e\ga0.6$ ) eccentricity binary."
+ These jiunbers are discrepant from the AAU truncation radius predicted by the detailed spectral modchueg of D'Alessioetal.(2005).. who examined radiative trauster uodels consisting of a geometrical thin wall aud a range of dust types.," These numbers are discrepant from the AU truncation radius predicted by the detailed spectral modeling of \citet{DAlessio05}, who examined radiative transfer models consisting of a geometrically thin wall and a range of dust types."
+ We will examine several wavs to resolve this discrepancy iu tuu the differing radiation oivttern of a binary versus a single star. errors in the stellar huninosity. aud the possibility of a substantially different radiative transfer iocdel.," We will examine several ways to resolve this discrepancy in turn: the differing radiation pattern of a binary versus a single star, errors in the stellar luminosity, and the possibility of a substantially different radiative transfer model."
+ A binary star has a different radiation pattern to a suele star., A binary star has a different radiation pattern to a single star.
+ Iu particular. on a circular ring of radius Ry ceeutered on the binary ceuter of mass. flux is ereatest at points on this rine aligned with the axis of the binary.," In particular, on a circular ring of radius $R_d$ centered on the binary center of mass, flux is greatest at points on this ring aligned with the axis of the binary."
+ For an SAAU binary (each componcut LAAT from the center of mass) aud a LLAAT rine. this geometry means that the fux from the binary is between 0.92 aud 1.28 times the flux from a sinele star with the same total huuinositv.," For an AU binary (each component AU from the center of mass) and a AU ring, this geometry means that the flux from the binary is between 0.92 and 1.28 times the flux from a single star with the same total luminosity."
+ We will asstune that the 115 Kk modeled by D'Alessioetal.(2005). represeuts, We will assume that the $\sim$ K modeled by \citet{DAlessio05} represents
+"Y. even when Y. corresponds to a density 97. lower: in fact when f,=Py the main contribution IC luminosity comes [rom region A and Eq 18. and Eq 10 hold: only Ys depends on densiy and o but the dependence are very weak ∪↾∖⊔↖-xn""τ−-- ","$Y$, even when $Y$ corresponds to a density $\delta^2$ lower; in fact when $\tilde{F}_p=F_p$, the main contribution to IC luminosity comes from region A and Eq \ref{eq:yf}
+ and Eq \ref{eq:ys1} hold: only $\tilde{Y}_S$ depends on density and $\delta$ but the dependence are very weak $\tilde{Y}_S\,(1+\tilde{Y}_S)^{(3-p)}\propto
+\tilde{n}^{(p-2)/2}\times \delta^{(p-2)}$ )."
+These illustrative examples o “how the model works should now be followed bv à modeling of real data., These illustrative examples of how the model works should now be followed by a modeling of real data.
+ In fact the general agreement between the two spectral shapes suggests that data. could also be we| reproduced. in. this framework. with possibly dillerent/ values. for the break frequencies and. peak [Lux respect tohe standard model.," In fact the general agreement between the two spectral shapes suggests that data could also be well reproduced in this framework, with possibly different values for the break frequencies and peak flux respect to the standard model."
+ As regard the corresponding estimated tdensity compared to the standard model. there are two competitive facts: the peak flux is brought to observec ranges by an higher density but the absorption frequency cannot be too higho (see fig.," As regard the corresponding estimated density compared to the standard model, there are two competitive facts: the peak flux is brought to observed ranges by an higher density but the absorption frequency cannot be too high (see fig."
+e 2 and 3) (note that. unlike the standard model a high. density. << 1070m7 does no produce an extreme inverse Compton component. which is not observed): Therefore. even if there is the possibility tha a higher density is required. the suitable value can be derivec within the errors only by a broadband modeling of data. with a Consistent. variations of all the parameters.," 2 and 3) (note that unlike the standard model a high density, $\gsim 10^{2}$ $^{-3}$ does not produce an extreme inverse Compton component, which is not observed); Therefore, even if there is the possibility that a higher density is required, the suitable value can be derived within the errors only by a broadband modeling of data, with a consistent variations of all the parameters."
+ Similarly to the standard. model there are univoca correlations between the estimated spectral parameters anc the derived. fireball energy. external density anc shock parameters.," Similarly to the standard model there are univocal correlations between the estimated spectral parameters and the derived fireball energy, external density and shock parameters."
+" They areMGUIl/de5/68:1/27!"" 3/20,:25/120,:3/42,:25/68:1/211/2:5,:11/125:1/40,:5/60:1/21.bj, in slow cooling5 regimeLes, while for a racdiativoe regime5PuA2S5/60:3/20,."," They are, in slow cooling regime while for a radiative regime."
+":17/65,:25/60:1/21771/26,:1/20,:2/30,:5/60:1/27οο... Some constraints must be added to this set of equations.", Some constraints must be added to this set of equations.
+" Jocatise €,Leg|ey=1. wheree, is the fractional energy that goes to protons. we most likely expect ο|CH*0.5."," Because $\epsilon_e + \epsilon_B + \epsilon_p=1$, where$\epsilon_p$ is the fractional energy that goes to protons, we most likely expect $\epsilon_e+ \epsilon_B<0.5$."
+ ‘The kinetic energy in the alterelow should not be much lower then the energy released in the ταν phase. otherwise very large elliciencies of. radiation. incompatible with internal shock models. would be required (Lazzati. Cihisellini οσο 1999).," The kinetic energy in the afterglow should not be much lower then the energy released in the $\gamma$ -ray phase, otherwise very large efficiencies of radiation, incompatible with internal shock models, would be required (Lazzati, Ghisellini Celotti 1999)."
+ Finally the apparent source size. the rate of expansion and the energy. in cach altcrelow should be consistent with observations in the radio wavebancds of interstcllar scintillation and. its. quenching (Waxman. Wulkarni Frail 1998).," Finally the apparent source size, the rate of expansion and the energy in each afterglow should be consistent with observations in the radio wavebands of interstellar scintillation and its quenching (Waxman, Kulkarni Frail 1998)."
+ In principle these requirements together with observations of the spectral parameters. at a particular time would set limits on 9 and therefore on rw structure of the magnetic field. generated in. shocks., In principle these requirements together with observations of the spectral parameters at a particular time would set limits on $\delta$ and therefore on the structure of the magnetic field generated in shocks.
+" ""nfortunatelvy the current data co not. allow a moctel independent estimation of spectra parameters.", Unfortunately the current data do not allow a model independent estimation of spectral parameters.
+ Actually the spectral shape. and in particular the break frequencies ancl the peak flux. are not unequivocaIv constrained by data in even the best studied. afterglows (e.g. GRB 970508. GRB YIN).," Actually the spectral shape, and in particular the break frequencies and the peak flux, are not unequivocally constrained by data in even the best studied afterglows (e.g. GRB 970508, GRB 000418)."
+ These spectral features are derived. by modcling the sparse measurement collected. in. various wavebancds at cilferent times: very often dillerent. versions of the standard niocel can fit the data. consequently finding. dillerent spectrals parameters (see for example Berger et al 2000. for GhD 000418).," These spectral features are derived by modeling the sparse measurement collected in various wavebands at different times: very often different versions of the standard model can fit the data, consequently finding different spectral parameters (see for example Berger et al 2000, for GRB 000418)."
+ Moreover these model dependent: quantities (above all 7) enter with higher dependences than 6 in the estimation of the fireball parameters (eq. 40--47)):, Moreover these model dependent quantities (above all $\tilde{\nu_a}$ ) enter with higher dependences than $\delta$ in the estimation of the fireball parameters (Eq. \ref{deltaE}- \ref{deltaebf}) );
+ therefore he break frequencies ancd peak Lux. clerivecl under the assumption that 6=1. cannot obviously be used to figure out the value of6.," therefore the break frequencies and peak flux, derived under the assumption that $\delta=1$, cannot obviously be used to figure out the value of $\delta$."
+ These equations and the physical limits on he fireball parameters should be used instead to constrain he overall modeling of data., These equations and the physical limits on the fireball parameters should be used instead to constrain the overall modeling of data.
+ The results would actually Head O very interesting conclusions., The results would actually lead to very interesting conclusions.
+ It the modeling succeeded in reproducing the data with oκ loa dillerent value for the external density. energy and shock parameters would. possibly be derived.," If the modeling succeeded in reproducing the data with $\delta<1$, a different value for the external density, energy and shock parameters would possibly be derived."
+ Hit indicates a denser environment. this could be more compatible with the strong observational evidence that would place GRBs explosions in a star forming region.," If it indicates a denser environment, this could be more compatible with the strong observational evidence that would place GRBs explosions in a star forming region."
+" Llowever a short scale magnetic field (sav 03;10. 1) would be probably ruled out. since extremely high. densities would. be required. in order to get [luxes in the observed. ranges and he racio H,) would. probably be too heavily absorbed: (in acüabatic) ancl fy.xὃ1yes (in radiative regime)."," However a short scale magnetic field (say $\delta\lsim 10^{-4}$ ) would be probably ruled out, since extremely high densities would be required in order to get fluxes in the observed ranges and the radio flux $\nu<\tilde{\nu_a}$ ) would probably be too heavily absorbed: $\tilde{F}_{\nu}\propto \tilde{n}^{-1/2}$ (in adiabatic) and $\tilde{F}_{\nu}\propto \delta^{-1}\,\tilde{n}^{-5/3}$ (in radiative regime)."
+ An other implication of the possibility o fit the very same data with models based on cülferent shoc‘k physics assumptions is that parameter estimations grealv suller from shock theory uncertainties., An other implication of the possibility to fit the very same data with models based on different shock physics assumptions is that parameter estimations greatly suffer from shock theory uncertainties.
+" In this work we ive only explored the effect of a short-scale magnetic field. but other reasonable variations of the standard model (such as time dependent eg.e, and p) may allect the parameter estimates well."," In this work we have only explored the effect of a short-scale magnetic field, but other reasonable variations of the standard model (such as time dependent $\epsilon_B,
+\epsilon_e $ and $p$ ) may affect the parameter estimates as well."
+ Phis would call for alternative model independent ernal density and energy estimates in order to constrain. together with the afterglow moceling. the shock physics.," This would call for alternative model independent external density and energy estimates in order to constrain, together with the afterglow modeling, the shock physics."
+ On the other hand if this model can reproduce data only with à.=l1. it could be either. evidence in. favor of an extremely stable Ποιά or the indication that some other modifications of the theory must. be taken into account.," On the other hand if this model can reproduce data only with $\delta=1$, it could be either evidence in favor of an extremely stable field or the indication that some other modifications of the theory must be taken into account."
+ For example. as already mentioned. ὁ can vary with time: for instance if the Bavleigh-Tavlor. instability occurs at the contact discontinuitv. the lenethscale woul erow as the fireball evolves and consequently late multibancds observations would be more consistent with àl. while carly measurements would. require à1.," For example, as already mentioned, $\delta$ can vary with time: for instance if the Rayleigh-Taylor instability occurs at the contact discontinuity, the lengthscale would grow as the fireball evolves and consequently late multibands observations would be more consistent with $\delta=1$, while early measurements would require $\delta\ll1$."
+ Observationa tests of this possibility can be performed. using the theory developed in this work: in fact we can look for a coheren changing in time of 6. modeling instantaneous spectra a different times. with all the remaining parameters. kep constant.," Observational tests of this possibility can be performed using the theory developed in this work; in fact we can look for a coherent changing in time of $\delta$ , modeling instantaneous spectra at different times, with all the remaining parameters kept constant."
+ OF course. witha very good dataset. this approach," Of, course, with a very good dataset, this approach"
+corresponding (very eccentric and. inclined) particle orbits also intersect the disc near to their. pericentres. (down to ~50au)) possibly accounting for the inferred inner warp of the dise. which is aligned with the outer [lared envelope in the SW.,"corresponding (very eccentric and inclined) particle orbits also intersect the disc near to their pericentres (down to $\sim$ ) possibly accounting for the inferred inner warp of the disc, which is aligned with the outer flared envelope in the SW."
+ In summary. there are three recognizable particle eroupings that can be related to the morphology of the + Pic disc.," In summary, there are three recognizable particle groupings that can be related to the morphology of the $\beta$ Pic disc."
+" ""μονο are: the highly. eccentric and inclined. particles that reach apocentre in the SW. (extension. B particles). the moderately eccentric and weakly inclined. particles that reach apocentre in the NIS (extension A particles). and the relatively unperturbed particles inside 7200 radius (Figs."," These are: the highly eccentric and inclined particles that reach apocentre in the SW (extension B particles), the moderately eccentric and weakly inclined particles that reach apocentre in the NE (extension A particles), and the relatively unperturbed particles inside $\sim$ radius (Figs."
+ 12 and 13)., $12$ and $13$ ).
+ We are grateful to the University of Hawaii for supporting the observations., We are grateful to the University of Hawaii for supporting the observations.
+ This work was funded in part bv a grant from NASA held by D. Jewitt., This work was funded in part by a grant from NASA held by D. Jewitt.
+ We also thank .J. Gracie. D. Zuckerman. and E. Becklin for the use of their coronagraph.," We also thank J. Gradie, B. Zuckerman, and E. Becklin for the use of their coronagraph."
+ The clarity of this paper was improved by comments from S. Ica., The clarity of this paper was improved by comments from S. Ida.
+ P. Ixalas acknowledges financial support and hospitality given by the Astronomy Unit. Queen Mary. Westfield College. during his visit there.," P. Kalas acknowledges financial support and hospitality given by the Astronomy Unit, Queen Mary Westfield College, during his visit there."
+flows aud the ecometry of field lines.,flows and the geometry of field lines.
+ So. for the outflow. the fast-maegnetosouic surface docs not need to cover all solid angles.," So, for the outflow, the fast-magnetosonic surface does not need to cover all solid angles."
+ We cau expect that the traus-fast MITD outflow is realized at least iu some part of the magnetic field lines. aud the distribution of the fast-iuagnetosouic surface would plav a very important part in explaining the ecucration and collimation of a highly accelerated jet or πα.," We can expect that the trans-fast MHD outflow is realized at least in some part of the magnetic field lines, and the distribution of the fast-magnetosonic surface would play a very important part in explaining the generation and collimation of a highly accelerated jet or wind."
+" We are grateful to Akira Tomimatsu. Sachiko Tsuruta aud Vasily S. Beskin for useful discussions aud to an anouvmous referee for constitutive criticism, which helped us to improve the paper."," We are grateful to Akira Tomimatsu, Sachiko Tsuruta and Vasily S. Beskin for useful discussions and to an anonymous referee for constitutive criticism which helped us to improve the paper."
+discussion was purely geometrical and did not need the description of an explicit six dimensional effective field theory.,discussion was purely geometrical and did not need the description of an explicit six dimensional effective field theory.
+ In the following discussion however it becomes essential for the branes to be BPS states., In the following discussion however it becomes essential for the branes to be BPS states.
+" When originally proposed[10],, the authors realized that their new mechanism (the enhancoon) could resolve a naked timelike singularity produced by a Dp-brane wrapped in a 2-cycle of the K3 surface which was being calledrepulson'""."," When originally \cite{JPP}, the authors realized that their new mechanism (the enhançoon) could resolve a naked timelike singularity produced by a $p$ -brane wrapped in a 2-cycle of the K3 surface which was being called."
+". In order to investigate it, they used the 10-dimensional supergravity of the system D2-D6 on K3 model although they showed that the same conclusions hold for any Dp-D(p +4) on K3model""."," In order to investigate it, they used the 10-dimensional supergravity of the system D2-D6 on K3 model although they showed that the same conclusions hold for any $p$ $(p+4)$ on K3."
+. The geometric locus of the enhangoon is independent of the model., The geometric locus of the enhançoon is independent of the model.
+" Let us see how the enhancoon comes up in our model, for the case p—0."," Let us see how the enhançoon comes up in our model, for the case $p=0$."
+" We will use the DO/D4-branes on the K3 surface, that is, without including so the gauge theory is abelian in the low energy approximation that follows."," We will use the D0/D4-branes on the K3 surface, that is, without including D2-branes, so the gauge theory is abelian in the low energy approximation that follows."
+" D2-branes are not relevant for the phenomenon we are describing, they do not ""see"" the enhangoon shell."," D2-branes are not relevant for the phenomenon we are describing, they do not “see” the enhançoon shell."
+ This fact allows us to simplify the computation but still trust the result as extended for the full model., This fact allows us to simplify the computation but still trust the result as extended for the full model.
+ The simplest static supergravity solution consistently truncated to its bosonic part can be written as: The line element corresponds to the string frame., The simplest static supergravity solution consistently truncated to its bosonic part can be written as: The line element corresponds to the string frame.
+" 457. is the metric of the K3 surface of unit volume, and εκ. is its corresponding volume form."," $dS_{K3}^2$ is the metric of the K3 surface of unit volume, and $\epsilon_{K3}$ is its corresponding volume form."
+" Providing that the solution is asymptotically flat, the harmonic functions are: where V is the volume of K3 at »=x and V,=(4a/,)! “."," Providing that the solution is asymptotically flat, the harmonic functions are: where $V$ is the volume of K3 at $r=\infty$ and $V_{*}=(4\pi l_s)^4=\frac{\mu_0}{\mu_4}$ ."
+ The volume of K3 at arbitrary r can be read offfrom (4.1)):, The volume of K3 at arbitrary $r$ can be read offfrom \ref{sugra}) ):
+Dwarf novae (DNe) are asubclass of cataclysmic variable (CV) svstems. in which a white dwarl (WD. the primary) aceretes hydrogen-rich matter from a low-mass main star (the secondary) filling its Roche lobe.,"Dwarf novae (DNe) are a subclass of cataclysmic variable (CV) systems, in which a white dwarf (WD, the primary) accretes hydrogen-rich matter from a low-mass main sequence-like star (the secondary) filling its Roche lobe."
+ In these svstems. the transferred gas forms an accretion disk around the WD.," In these systems, the transferred gas forms an accretion disk around the WD."
+ It is believed (hat the accretion disk is subject to a thermal instability Chat eauses evelie changes of the accretion rate., It is believed that the accretion disk is subject to a thermal instability that causes cyclic changes of the accretion rate.
+" A low rate of accretion (z10MAL. 1) quiescent stage is followed everv few weeks to months by a high rate of accretion (z2LO“AL. "" outburst stage of days to weeks.", A low rate of accretion $\approx 10^{-11} M_{\odot}$ $^{-1}$ ) quiescent stage is followed every few weeks to months by a high rate of accretion $\approx 10^{-8} M_{\odot}$ $^{-1}$ ) outburst stage of days to weeks.
+ These outbursts (dwarf nova - DN accretion event or nova-like high state). are believed to be punctuated every few thousand vears or more by a (thermonuclear runaway. (TNR) explosion: the classical nova 1993).," These outbursts (dwarf nova - DN accretion event or nova-like high state), are believed to be punctuated every few thousand years or more by a thermonuclear runaway (TNR) explosion: the classical nova \citep{hac93}."
+ The WD dominates the far uliraviolet (FUV) in many. and probably. most. DNe in quiescence.," The WD dominates the far ultraviolet (FUV) in many, and probably, most, DNe in quiescence."
+ As a consequence. quiescent DNe provide a unique laboratory not only for understanding the physics of accretion but also for understanding physical processes of WDs.," As a consequence, quiescent DNe provide a unique laboratory not only for understanding the physics of accretion but also for understanding physical processes of WDs."
+ As a result.Telescope (LST) and other UV observatories have been used to directly observe the effects of accretion on the WDs of these svstems.," As a result, ) and other UV observatories have been used to directly observe the effects of accretion on the WDs of these systems."
+ These studies have vielded determinations of effective temperature. (he rotation rates. photospheric abundances. cooling rates following outburst. and dviamical masses in a number of DNe 1999)..," These studies have yielded determinations of effective temperature, the rotation rates, photospheric abundances, cooling rates following outburst, and dynamical masses in a number of DNe \citep{sio99,gan99}."
+ llowever. there are some indications of an additional component besiles the white dwarf in the spectra of some DNe in quiescence. such as the presence of emission lines and the bottoms of Lyman alpha profiles which do not eo to zero as in a pure white dwarl.," However, there are some indications of an additional component besides the white dwarf in the spectra of some DNe in quiescence, such as the presence of emission lines and the bottoms of Lyman alpha profiles which do not go to zero as in a pure white dwarf."
+ Possible locations of this additional component could be: (1) a heated reeion of the WD: (2) direct emission from the boundary laver: (3) an opticalle thick region of the disk: or (4) a corona/chromosphere above a cool disk (Ixo&Ixallman&AMever-Hofneister 1997).," Possible locations of this additional component could be: (1) a heated region of the WD; (2) direct emission from the boundary layer; (3) an optically thick region of the disk; or (4) a corona/chromosphere above a cool disk \citep{kok89,kok92,mey89,ko96,liu97,lad97}."
+. VW [Isi is a Κον svstem for understanding DNe in general., VW Hyi is a key system for understanding DNe in general.
+ It is the closest. pe) and brightest example of an SU UMa-tvpe DN ancl it lies along a line of sight with an exceptionally low interstellar column," It is the closest \citep[placed it at 65 pc]{war87}
+ and brightest example of an SU UMa-type DN and it lies along a line of sight with an exceptionally low interstellar column"
+and the planet must be accounted for to produce accurate diffusion coefficieuts (22)... (?)..,"and the planet must be accounted for to produce accurate diffusion coefficients \citep{MT99,PS04}. \citep{ZSZ02}."
+is the square of the adiabatic sound. speed.,is the square of the adiabatic sound speed.
+ As in compressible MIID. the WIND dispersion relation has 6 requency eigenvalues o for any choice of the real wavevector Kk.," As in compressible MHD, the WKB dispersion relation has 6 frequency eigenvalues $\omega$ for any choice of the real wavevector $\bk$."
+ M any of the eigenvalues has a non-zero imaginary part. he system experiences a local instability.," If any of the eigenvalues has a non-zero imaginary part, the system experiences a local instability."
+ Otherwise the eigenvalues are all real ancl the group velocities Ow/Ok; are independent. of ΚΙ. indicating that the wave propagation in his limit is anisotropic but non-dispersive.," Otherwise the eigenvalues are all real and the group velocities $\partial\omega/\partial k_i$ are independent of $|\bk|$, indicating that the wave propagation in this limit is anisotropic but non-dispersive."
+ The system of equations is then formally hyperbolic., The system of equations is then formally hyperbolic.
+ A full investigation of the dispersion relation is cillicul out. the following observations may be mace., A full investigation of the dispersion relation is difficult but the following observations may be made.
+" First. when AR),""=M;;-—0 the dispersion relation is the standard one of compressible hydrodynamics. and information is propagate at the sound speed relative to the mean flow."," First, when $R_{ij}=M_{ij}=0$, the dispersion relation is the standard one of compressible hydrodynamics, and information is propagated at the sound speed relative to the mean flow."
+ Second. when Ry=0 but Aj;=0. the system can. be shown to be hyperbolic.," Second, when $R_{ij}=0$ but $M_{ij}\ne0$, the system can be shown to be hyperbolic."
+ The dispersion. relation is related. to that. of compressible MED., The dispersion relation is related to that of compressible MHD.
+ Third. in an incompressible Iuid. modes with &;u40 are eliminated as their speed of propagation is infinite.," Third, in an incompressible fluid, modes with $k_iu_i'\ne0$ are eliminated as their speed of propagation is infinite."
+ The remaining part of the dispersion. relation. given by the first. line of equation (46)). describes. the propagation of transverse modes. similar to Alfyvenn waves. at finite speed.," The remaining part of the dispersion relation, given by the first line of equation \ref{dispersion}) ), describes the propagation of transverse modes, similar to Alfvénn waves, at finite speed."
+ Finally. there are circumstances in which the system fails to be hyperbolic because the dispersion relation indicates a local instability. but this appears to be atypical.," Finally, there are circumstances in which the system fails to be hyperbolic because the dispersion relation indicates a local instability, but this appears to be atypical."
+ ln a stratified shearing sheet. magnetorotational turbulence would. develop. according to these equations just as in an incompressible shearing sheet. except that all stress components would be proportional to p/O7 rather than £7.," In a stratified shearing sheet, magnetorotational turbulence would develop according to these equations just as in an incompressible shearing sheet, except that all stress components would be proportional to $p/\Omega_z^2$ rather than $L^2$."
+" The total shear stress A4,—[2,4 can be compared with that corresponding to an elfective viscosity (i;=op/O.", The total shear stress $M_{xy}-R_{xy}$ can be compared with that corresponding to an effective viscosity $\mu=\alpha p/\Omega$.
+ For the fiducial parameters in a Ixeplerian disc. this gives àzz0.051.," For the fiducial parameters in a Keplerian disc, this gives $\alpha\approx0.081$ ."
+ fall the parameters C are scaled by a constant factor. this value of a scales as 1/67.," If all the parameters $C_i$ are scaled by a constant factor, this value of $\alpha$ scales as $1/C_i^2$."
+" When the Rossby number Ro=ο of the shearing sheet is varied in the range 0«Ro<1I in which the svstem is Ravleigh-stable but magnetorotationally unstable. the total shear stress Ay—[6,4 in the steady turbulent state is not simply proportional to the shear rate 2:1. at. fixecl angular velocity."," When the Rossby number ${\rm Ro}=A/\Omega$ of the shearing sheet is varied in the range $0<{\rm Ro}<1$ in which the system is Rayleigh-stable but magnetorotationally unstable, the total shear stress $M_{xy}-R_{xy}$ in the steady turbulent state is not simply proportional to the shear rate $2A$, at fixed angular velocity."
+ This is shown in Lig., This is shown in Fig.
+ 6. and demonstrates one aspect in which the present model dilfers significantly from a viscous representation of the turbulence.," 6, and demonstrates one aspect in which the present model differs significantly from a viscous representation of the turbulence."
+ The local dispersion relation for axisvnunetric density waves in à two-dimensional disc mocoel is where & ds the radial wavenumber (Goldreich ‘Tremaine 1979)., The local dispersion relation for axisymmetric density waves in a two-dimensional disc model is where $k$ is the radial wavenumber (Goldreich Tremaine 1979).
+ The same mode can be identified: within the present model if one considers ao three-dimensiona compressible shearing sheet without vertical gravity ane stratification., The same mode can be identified within the present model if one considers a three-dimensional compressible shearing sheet without vertical gravity and stratification.
+ “To avoid complications it is convenient to replace the thermal energy equation by the isotherma condilion p=ορ. ce=constant.," To avoid complications it is convenient to replace the thermal energy equation by the isothermal condition $p=c_{\rm s}^2\rho$, $c_{\rm s}={\rm constant}$ ."
+ Equations (40)) (43)). linearized around ai basic state consisting of steady. homogeneous magnetorotational turbulence. then admit solutions proportional to explikweiat) with no »rturbations of (i.oiIusMILM).," Equations \ref{compressible_rho}) \ref{compressible_mij}) ), linearized around a basic state consisting of steady, homogeneous magnetorotational turbulence, then admit solutions proportional to $\exp({\rm i}kx-{\rm
+ i}\omega t)$, with no perturbations of $(u_z,R_{xz},R_{yz},M_{xz},M_{yz})$."
+ In the absence of urbulent stresses. the dispersion relation (47)) is recovered.," In the absence of turbulent stresses, the dispersion relation \ref{gt}) ) is recovered."
+ A numerical solution of the dispersion relation in. the oesence of turbulent stresses. indicates that the two-dimensional wave is camped in a Weplerian disc when the iducial parameters €;=1 are adopted (Fig., A numerical solution of the dispersion relation in the presence of turbulent stresses indicates that the two-dimensional wave is damped in a Keplerian disc when the fiducial parameters $C_i=1$ are adopted (Fig.
+ 7)., 7).
+" The real mut of the frequency. is somewhat greater than predicted w equation (47)). and agrees much better if e; is replaced w the ""magnetosonie speed? (e|Mp)7."," The real part of the frequency is somewhat greater than predicted by equation \ref{gt}) ), and agrees much better if $v_{\rm
+ s}$ is replaced by the `magnetosonic speed' $(c_{\rm
+ s}+M/\rho)^{1/2}$."
+ Interestingly. he imaginary part corresponds quite closely with a viscous damping rate (A7 if the ellective viscosity is”=acl/O with a0.06.," Interestingly, the imaginary part corresponds quite closely with a viscous damping rate $\nu k^2$ if the effective viscosity is $\nu=\alpha c_{\rm s}^2/\Omega$ with $\alpha\approx0.06$."
+ This is remarkable when it is considered that there are no viscous or dilfusive terms in the governing equations., This is remarkable when it is considered that there are no viscous or diffusive terms in the governing equations.
+ When a stratified disc is warped. horizontal pressure eraclicnts are introduced that drive epievclic motions with an amplitude proportional to the distance z above the mic-plane (Papaloizou Pringle 1983).," When a stratified disc is warped, horizontal pressure gradients are introduced that drive epicyclic motions with an amplitude proportional to the distance $z$ above the mid-plane (Papaloizou Pringle 1983)."
+ Phe amplitude: and phase of these oscillations are critical to the propagation ancl damping of the warp., The amplitude and phase of these oscillations are critical to the propagation and damping of the warp.
+ As discussed by Torkelsson ct al. (, As discussed by Torkelsson et al. (
+2000). the dywnamies of these shearing epicyvelicmotions. and in particular their damping time-scale. can be studiedwithin the shearing sheet.,"2000), the dynamics of these shearing epicyclicmotions, and in particular their damping time-scale, can be studiedwithin the shearing sheet."
+ Consider equations (40)) (43)). linearized around abasic state consisting οἱ steacky. homogeneous," Consider equations \ref{compressible_rho}) \ref{compressible_mij}) ), linearized around a basic state consisting of steady, homogeneous"
+ Consider equations (40)) (43)). linearized around abasic state consisting οἱ steacky. homogeneous|," Consider equations \ref{compressible_rho}) \ref{compressible_mij}) ), linearized around a basic state consisting of steady, homogeneous"
+ejants are rare.,giants are rare.
+ Tables 1 and 2 give the atmospheric parameters (Zi. logg. and. [Fe/IH]) derived from iron lines in the high-resolution spectra.," Tables 1 and 2 give the atmospheric parameters $T_{\rm eff}$ , $\log g$, and [Fe/H]) derived from iron lines in the high-resolution spectra."
+ A majority of the entries for Table 2 are taken from the papers reporting the Li abundances., A majority of the entries for Table 2 are taken from the papers reporting the Li abundances.
+ The LTE and Non-LTE Li abunelances are tabulated) where (he latter are based on (he recipe given in Lindetal.(2009)., The LTE and Non-LTE Li abundances are tabulated where the latter are based on the recipe given in \citet{lind2009}.
+. In the Tables. we also present C/C values for the Li-rich IN. giants.," In the Tables, we also present $^{12}$ $^{13}$ C values for the Li-rich K giants."
+ The values of C/C measurements are based on molecular lines in the rreeion using a similar procedure (to that. described in Kumar&Reedy(2009)., The values of $^{12}$ $^{13}$ C measurements are based on molecular lines in the region using a similar procedure to that described in \citet{bharat2009}.
+.. For a lew giants carbon ratios couldu't be determined owing to their large rotation and as a result broad spectral features., For a few giants carbon ratios couldn't be determined owing to their large rotation and as a result broad spectral features.
+ Giants [rom our and Brown et al., Giants from our and Brown et al.
+'s surveys are shown on the IIR. diagram in Fig.,'s surveys are shown on the HR diagram in Fig.
+ 2 along with evolutionary. tracks (Bertellietal.2008) computed for solar metallicity. |Fe/11] = 0.0. and stellar masses ranging [rom 0.544. to 3M...," 2 along with evolutionary tracks \citep{bertelli2008} computed for solar metallicity, [Fe/H] = 0.0, and stellar masses ranging from $M_{\odot}$ to $M_{\odot}$."
+ All stars observed only at low resolution are asstuned {ο have a solar metallicity., All stars observed only at low resolution are assumed to have a solar metallicity.
+ For Li-rich. IX giants actual metallicities measured [rom the high resolution spectra are adopted aud most are close to solar values., For Li-rich K giants actual metallicities measured from the high resolution spectra are adopted and most are close to solar values.
+ Errors derived from the quoted uncertainties in the parallaxes ancl in the derivation of temperatures for the Li-rich IX giants are marked in Fig., Errors derived from the quoted uncertainties in the parallaxes and in the derivation of temperatures for the Li-rich K giants are marked in Fig.
+ 2., 2.
+ All but a very few stars in the two samples have evolved bevond the base of the RGB (the red dashed line in Fig., All but a very few stars in the two samples have evolved beyond the base of the RGB (the red dashed line in Fig.
+ 2) and. therefore. their surface Li abundance has been greatly diluted.," 2) and, therefore, their surface Li abundance has been greatly diluted."
+ In panel a of Fig., In panel a of Fig.
+ 2. there is (he expected concentration of stars will Iuminosities of about log1.7.," 2, there is the expected concentration of stars with luminosities of about $\log \simeq 1.7$."
+ This population in the IIR. diagram in the main are red clump (horizontal branch) He-core burning stars that have evolved [rom the tip of the RGB: the thick black line in Fie., This population in the HR diagram in the main are red clump (horizontal branch) He-core burning stars that have evolved from the tip of the RGB: the thick black line in Fig.
+ 2 shows the predicted locations of clump stars., 2 shows the predicted locations of clump stars.
+ Extension of the logL/L.c1.7 concentration to cooler temperatures may include stars at the bump Iuminositw of the first ascent RGB stars (see red portions of (he evolutionary tracks)., Extension of the $\log L/L_\odot \simeq 1.7$ concentration to cooler temperatures may include stars at the bump luminosity of the first ascent RGB stars (see red portions of the evolutionary tracks).
+ Dominance of the HA diagram by clump stars is expected because the lifetime of a clump star is (vpically one (o two orders of magnitude longer (han for a bump stu., Dominance of the HR diagram by clump stars is expected because the lifetime of a clump star is typically one to two orders of magnitude longer than for a bump star.
+ This argument is not strictly applicable to the Li-rvich stars because their frequency is also dependent on the probabilities that clump and bunmp stars produce lithium., This argument is not strictly applicable to the Li-rich stars because their frequency is also dependent on the probabilities that clump and bump stars produce lithium.
+ Except for five Iuminous examples. (he Li-rich stars are either at the Iuminosity of the clump or within (he range of expected values for the bump. a range spanning (he clump's luminosity.," Except for five luminous examples, the Li-rich stars are either at the luminosity of the clump or within the range of expected values for the bump, a range spanning the clump's luminosity."
+ Thefive luminous exceptions must owe their Li richness to different phenomena, Thefive luminous exceptions must owe their Li richness to different phenomena
+and 2000 FebruaryAlarch as part of a larger. study to measure the rotation periods of the late-twpe RECKX stars in the 5g Cha cluster.,and 2000 February–March as part of a larger study to measure the rotation periods of the late-type RECX stars in the $\eta$ Cha cluster.
+ Lawson ct al. (, Lawson et al. (
+2001) found that all of these stars were variable. with periods attributed. to rotational modulation of cool starspots.,"2001) found that all of these stars were variable, with periods attributed to rotational modulation of cool starspots."
+ Dillerential V -band observations of the 2 new cluster members were obtained with respect to nearby stars within their CCD frames., Differential $V$ -band observations of the 2 new cluster members were obtained with respect to nearby stars within their CCD frames.
+ The 1999 and 2000 data sets for cach star were independently analysecl using the Lomb-Scargle Fourier method for non-equally spaced cata. with the frequeney range of f=O2 d+ examined for periodicities.," The 1999 and 2000 data sets for each star were independently analysed using the Lomb-Scargle Fourier method for non-equally spaced data, with the frequency range of $f = 0-2$ $^{-1}$ examined for periodicities."
+ Both stars showed periodic variations in cach vear., Both stars showed periodic variations in each year.
+ Phased light curves are shown in Fig., Phased light curves are shown in Fig.
+ 5. and details of the periodicities are listed in Table 2.," 5, and details of the periodicities are listed in Table 2."
+ The S/N ratio for each periodicity was determined by measuring the residual noise level in the pre-whitened data sets., The S/N ratio for each periodicity was determined by measuring the residual noise level in the pre-whitened data sets.
+ For ECLA J0841.57853. a low-amplitude 173-4 periodicity was recovered in cach vear.," For ECHA J0841.5–7853, a low-amplitude 1.73-d periodicity was recovered in each year."
+ As with 1o [ate-type RECN stars. we associate the variations with the rotationa moculation of cool spots.," As with the late-type RECX stars, we associate the variations with the rotational modulation of cool spots."
+ Following Allain et al. (, Following Allain et al. (
+1996) we calculated the lower limit to the fractional spot coverage by assuming the spots were dark.,1996) we calculated the lower limit to the fractional spot coverage by assuming the spots were dark.
+ The spot fractions are listec as percentages in Table 2., The spot fractions are listed as percentages in Table 2.
+ Lawson ct al. (, Lawson et al. (
+2001) correctec the photometry of the RECN stars for the ellects of the starspots by estimating the V mag and colours of the unspotted. or minimume-spotted. star from the phase of the photometry and. the V-band. amplitude of the star.,"2001) corrected the photometry of the RECX stars for the effects of the starspots by estimating the $V$ mag and colours of the unspotted, or minimum-spotted, star from the phase of the photometry and the $V$ -band amplitude of the star."
+ The measurements listed in Table 1. were obtained near maximum light and thus represent the unspotted. or minimunmespotted. star.," The measurements listed in Table 1 were obtained near maximum light and thus represent the unspotted, or minimum-spotted, star."
+ For ECLA J0S843.87905. a 12-d. periodicitv ας measured in cach vear. but with dillerent structure in the light curve.," For ECHA J0843.3–7905, a 12-d periodicity was measured in each year, but with different structure in the light curve."
+ During 1999 February. a short-cduration (2Lx d) peak was observed that may have been an optical Hare.," During 1999 February, a short-duration $2-3$ d) peak was observed that may have been an optical flare."
+ Underlving the peak of AV=0.65 mag is a 0.2)9 mag amplitude quasi-sinusoidal variation that dominates the Fourier analysis., Underlying the peak of $\Delta V = 0.65$ mag is a $0.2-0.3$ mag amplitude quasi-sinusoidal variation that dominates the Fourier analysis.
+ Phe light curve during 2000 has a cilferent appearance. but with the same underlying period ancl V-band amplitude.," The light curve during 2000 has a different appearance, but with the same underlying period and $V$ -band amplitude."
+ Phe photometry reported. in Table l was obtained in 2000 near maximum light., The photometry reported in Table 1 was obtained in 2000 near maximum light.
+ Given the large photometric amplitude and the variable structure of the light curve probably driven by accretion hotspots. these data may not represent the unspotted photospheric values.," Given the large photometric amplitude and the variable structure of the light curve probably driven by accretion hotspots, these data may not represent the unspotted photospheric values."
+ Analysis of the AISSSO spectra showed 3 of these stars (GSC 9402.11003. GSC 9398.00090. and GSC 940300831). were Iate-Ix. or earlv-M. spectral twpe giants with Lain narrow absorption with zLA ," Analysis of the MSSSO spectra showed 3 of these stars (GSC 1003, GSC 0099 and GSC 0831) were late-K or early-M spectral type giants with $\alpha$in narrow absorption with $\approx 1$ "
+assured. effectively unweighting this coustraiut.,"assumed, effectively unweighting this constraint."
+ The MDM. proper motion is (+7L3.—19.5) mas t. leading to a proper motion probability density peaking at x—15 mas.," The MDM proper motion is $+74.3,-19.5$ ) mas $^{-1}$, leading to a proper motion probability density peaking at $\pi = 15$ mas."
+ The final distance estimate is. 13.3(41.6.—1.5) pe.," The final distance estimate is $43.3 (+1.6, -1.5)$ pc."
+ The literature contaius many other estimates for the distance of WZ See., The literature contains many other estimates for the distance of WZ Sge.
+ Smak(1993) estimated 48+10 ye from the flux of the white dwarl. for which he adopts ία=ϐϱ.1211... Sionetal.," \citet{smak93} estimated $48 \pm 10$ pc from the flux of the white dwarf, for which he adopts $M_{\rm wd} = 0.45 M_{\odot}$."
+(1995) 11odel aun HST ultraviolet spectrum with a logy=5 white dwarf. aud fiud 69 pc lor the clistauce.," \citet{sionwz} model an HST ultraviolet spectrum with a $\log g = 8$ white dwarf, and find 69 pc for the distance."
+ However. 5oruit&Rutten(CLOS) point out that the temperature aud gravity are highly degenerate in this kiud of fit: largely [rom a study of the stream cyuamics they adopt logg=9 for the white dwarl. :uxl argue that te Sionetal.(1995) distauce should be adjusted downward to LS pc.," However, \citet{spruit}
+ point out that the temperature and gravity are highly degenerate in this kind of fit; largely from a study of the stream dynamics they adopt $\log g = 9$ for the white dwarf, and argue that the \citet{sionwz} distance should be adjusted downward to 48 pc."
+ The short ¢istance determied here supports their interpretation. auc st1geestsMOD that the white dwar ‘in WZ See is relatively hig-gravity (hence massive).," The short distance determined here supports their interpretation, and suggests that the white dwarf in WZ Sge is relatively high-gravity (hence massive)."
+ WZ See is evidently the closest. known cataclysmic binary., WZ Sge is evidently the closest known cataclysmic binary.
+“ While this paper was in the final stages of p'eparation. two otler Ineasurements of the parallax ol WZ See came to my attention.," While this paper was in the final stages of preparation, two other measurements of the parallax of WZ Sge came to my attention."
+ First. C. Dahu (privale coniuntuication) kindly passed aloug a USNO ayy for WZ See closely agreeing with tle present determination.," First, C. Dahn (private communication) kindly passed along a USNO $\pi_{\rm rel}$ for WZ Sge closely agreeing with the present determination."
+" Secoucd. anuounced an even more precise Tap,=22.97E015 nas [‘ot the the HST Fine Cruidance System. which also agrees accuratey with this one."," Second, \citet{harrison03a} announced an even more precise $\pi_{\rm abs} = 22.97 \pm 0.15$ mas from the the HST Fine Guidance System, which also agrees accurately with this one."
+ With tli‘ee Ineependeut determinations giving οΠΠ the same value. we cau ye Very COuiclent about the clistarn'e to WZ See.," With three independent determinations giving essentially the same value, we can be very confident about the distance to WZ Sge."
+ The proper is sinall. (+2.-16 uas ον which compares with (+11.2.—25.2) lias Lin the Lick NPM.," The proper is small, $(+2,-16)$ mas $^{-1}$ , which compares with $(+11.2,-25.2)$ mas $^{-1}$ in the Lick NPM."
+ The photonetrle corSraint is based ou Wo=+12.0 or normal outbuIsis (Roseuzweigetal. 2000).. aud the Warner reatiou with an ill-constraiued binary inclination o. 60 degrees. based ou the moclest velocity amplittice CThorstensen.Wade.&Oke1986).," The photometric constraint is based on $V = +12.0$ for normal outbursts \citep{rosensu}, , and the Warner relation with an ill-constrained binary inclination of 60 degrees, based on the modest velocity amplitude \citep{thorwade86}."
+.. The parallax. Habs=ντ tas. ds 1ot well-deteriinec. auc he clistance estimate cleyencds critically ou the e‘LOL estimate adopted: the scatter about the bes fit indicates σς 1.L1jns. while the scatter of sis. comparison stars within 2 mag of Sl UMa eives 2.1 mas.," The parallax, $\pi_{\rm abs} = 7.4$ mas, is not well-determined, and the distance estimate depends critically on the error estimate adopted; the scatter about the best fit indicates $\sigma = 1.4$ mas, while the scatter of six comparison stars within 2 mag of SU UMa gives 2.4 mas."
+ AcoptinD>2 the smaller error vielc Soa Bayesian estimate ὁ[ 189 (279.—13) pe: i£ te larger error is adopte. the Lt1z-Ixeker correction increases the estimae to 103 (2-230.2162) pe.," Adopting the smaller error yields a Bayesian estimate of 189 $(+79,-43)$ pc; if the larger error is adopted, the Lutz-Kelker correction increases the estimate to 403 $(+230,-162)$ pc."
+ While he distance is disappointingly iudeteriiuale. 120 pe is a reasonable lower liit.," While the distance is disappointingly indeterminate, 120 pc is a reasonable lower limit."
+ This star resembles WZ See iu that it otbursts only rarely aux with la‘oe amplitude., This star resembles WZ Sge in that it outbursts only rarely and with large amplitude.
+ Because of its faiuuess. the parallax is somewhat uncertain at Ways=25.5+2.2 mas.," Because of its faintness, the parallax is somewhat uncertain at $\pi_{\rm abs} = 5.8 \pm 2.2$ mas."
+ The proper motion is moclest. (219.—12) mas Ll.," The proper motion is modest, $(+19,-12)$ mas $^{-1}$."
+ in good ag'eelueut. with (4-22.—8) in USNO 51.0.," in good agreement with $(+22, -8)$ in USNO B1.0."
+ The superoutburst ligi curve (Leibowitzetal.1991) reaches V11.5. iudicating that if normal outbursts occurred Gwhich they apparently dout). tvey would reach ~12.5.," The superoutburst light curve \citep{leib94} reaches $V = 11.5$, indicating that if normal outbursts occurred (which they apparently don't), they would reach $\sim 12.5$."
+ ΤΙe orbital period. [rom low-state plotometry. is 0.05799 d. anc the presence of a low-state m0¢uallon suggests a substantial orbita| incliuation. vieldiug au estimate My=5.943.," The orbital period, from low-state photometry, is 0.05799 d, and the presence of a low-state modulation suggests a substantial orbital inclination, yielding an estimated $M_V = 5.9 \pm 3$."
+ The τοπilας photometric constraint peaks 1ear 7=5 mas., The resulting photometric constraint peaks near $\pi = 5$ mas.
+ The sinall proper motion puts the star more cli:stant. with a peak heal ü3 nas.," The small proper motion puts the star more distant, with a peak near $\pi = 3$ mas."
+" Therather weak parallax determination leads to a substantial al effect, and a final Bayesian cisauce estimate of 100 (4530.—180) pc."," Therather weak parallax determination leads to a substantial $\pi^{-4}$ effect, and a final Bayesian distance estimate of 460 $+530, -180$) pc."
+ For comparison. Szkocdyetal.(20024) ," For comparison, \citet{szkodyhv} "
+Star formation in the presence of turbulence is expected to be a stochastic process (Larson. 2001). and therefore a statistical ensemble of simulations must be performed in order to compare meaningfully with observations.,"Star formation in the presence of turbulence is expected to be a stochastic process (Larson 2001), and therefore a statistical ensemble of simulations must be performed in order to compare meaningfully with observations."
+ In this paper we present an ensemble of 20 simulations of star formation in a realistic (1e. observationally motivated). dense core.," In this paper we present an ensemble of 20 simulations of star formation in a realistic (i.e. observationally motivated), dense core."
+ Here we outline the observational constraints on the properties of molecular cores that are on the verge of protostellar collapse and deseribe the initial conditions that we adopt., Here we outline the observational constraints on the properties of molecular cores that are on the verge of protostellar collapse and describe the initial conditions that we adopt.
+ Molecular cores can be divided into those that appear to have formed protostars in their centres — 1.9. those associated with IRAS sources — and those that do not (e.g. Beichman et al., Molecular cores can be divided into those that appear to have formed protostars in their centres – i.e. those associated with IRAS sources – and those that do not (e.g. Beichman et al.
+ 1986)., 1986).
+ The latter show no evidence for outflows. and are usually referred to as starless cores.," The latter show no evidence for outflows, and are usually referred to as starless cores."
+ Starless cores typically have masses of a few M... volume densities of 10? to 10!em? radii ~0.1pe and are approximately isothermal. with temperatures ~IOK (e.g. Jijina et al.," Starless cores typically have masses of a few $M_{\odot}$, volume densities of $10^3$ to $10^4\,{\rm cm}^{-3}$, radii $\sim\,0.1\,{\rm pc}$ and are approximately isothermal, with temperatures $\sim\,10\,{\rm K}$ (e.g. Jijina et al."
+" 1999),", 1999).
+ Those starless cores that are most centrally condensed. and hence presumably closest to protostellar collapse. are known as pre-stellar cores (originally pre-protostellar cores — Ward-Thompson et al.," Those starless cores that are most centrally condensed, and hence presumably closest to protostellar collapse, are known as pre-stellar cores (originally pre-protostellar cores – Ward-Thompson et al."
+ 1994)., 1994).
+ Their density profiles are approximately flat within a central region of a few thousand au. and then fall as /77 before steepening even further to 17+ or r? at their edges (e.g. Ward-Thompson et al.," Their density profiles are approximately flat within a central region of a few thousand au, and then fall as $r^{-2}$ before steepening even further to $r^{-4}$ or $r^{-5}$ at their edges (e.g. Ward-Thompson et al."
+ 1994: André et al., 1994; André et al.
+ 1996; Ward-Thompson et al., 1996; Ward-Thompson et al.
+ 1999)., 1999).
+ They finally become indistinguishable from the background molecular cloud at their edges (see André et al. (, They finally become indistinguishable from the background molecular cloud at their edges (see André et al. (
+2000) for a review of core properties).,2000) for a review of core properties).
+ Observations. of cores show that they often have a significant non-thermal contribution to their line widths. which can be attributed to turbulence.," Observations of cores show that they often have a significant non-thermal contribution to their line widths, which can be attributed to turbulence."
+ Larson (1981) has shown that on average the line-of-sight velocity dispersions. c. of molecular clouds and cores are related to their tangential linear sizes. L. by a relation of the form although later studies have suggested that the exponent should be higher than 0.38 (e.g. Myers 1983) and depends on the presence of embedded IRAS sources (Jijma et al.," Larson (1981) has shown that on average the line-of-sight velocity dispersions, $\sigma$, of molecular clouds and cores are related to their tangential linear sizes, $L$, by a relation of the form although later studies have suggested that the exponent should be higher than 0.38 (e.g. Myers 1983) and depends on the presence of embedded IRAS sources (Jijina et al."
+ 1999)., 1999).
+ Equation (1)) i5 normally interpreted as evidence for turbulence within molecular clouds (see Larson 2003 for a review)., Equation \ref{LARSON1}) ) is normally interpreted as evidence for turbulence within molecular clouds (see Larson 2003 for a review).
+ A key parameter characterising the level of turbulence in à core Is the ratio of turbulent to gravitational energy. cuu.," A key parameter characterising the level of turbulence in a core is the ratio of turbulent to gravitational energy, $\alpha_{\rm turb}$."
+ Figure 2 shows a plot of ay» against core mass for the starless cores in the Jijina et al. (, Figure \ref{fig:nonthermal} shows a plot of $\alpha_{\rm turb}$ against core mass for the starless cores in the Jijina et al. (
+1999) catalogue.,1999) catalogue.
+ We note that there is some difficulty in accurately determining the masses of cores. and hence their gravitational energies. but we believe that the points plotted in Fig.," We note that there is some difficulty in accurately determining the masses of cores, and hence their gravitational energies, but we believe that the points plotted in Fig."
+ 2 wre accurate to within a factor of two or better. and representative of dense cores which are destined to form stars.," \ref{fig:nonthermal} are accurate to within a factor of two or better, and representative of dense cores which are destined to form stars."
+ Figure 2 shows that the average μην is low: almost all of the observeddense cores have ey.<0.5. and most have O«Gun<0.3.," Figure \ref{fig:nonthermal} shows that the average $\alpha_{\rm turb}$ is low; almost all of the observeddense cores have $\alpha_{\rm turb}<0.5$, and most have $0<\alpha_{\rm turb}<0.3$."
+ The star symbol in the lower half of Fig., The star symbol in the lower half of Fig.
+ 2 represents the parameters (mass and turbulence) of the cores we have simulated here. chosen to be representative of dense cores with low levels of turbulence.," \ref{fig:nonthermal} represents the parameters (mass and turbulence) of the cores we have simulated here, chosen to be representative of dense cores with low levels of turbulence."
+ The open circle in the upper right of Fig., The open circle in the upper right of Fig.
+ 2. represents the parameters of the core simulated by Bate et al. (, \ref{fig:nonthermal} represents the parameters of the core simulated by Bate et al. (
+2002a.b. 2003).,"2002a,b, 2003)."
+ A Plummer-like profile of the form aCO good fitGR) to the observed density profiles. of. densegives cores (see Fig. D)., A Plummer-like profile of the form gives a good fit to the observed density profiles of dense cores (see Fig. \ref{fig:profile}) ).
+ pia is the central density (here 3x gem™).," $\rho_{\rm kernel}$ is the central density (here $3 \times 10^{-18}\,{\rm g}\,{\rm cm}^{-3}$ )."
+ Ryemet is the radius of the central region of the core — hereafter the kernel — in which the density is approximately uniform (here 5.000 au).," $R_{\rm kernel}$ is the radius of the central region of the core – hereafter the kernel – in which the density is approximately uniform (here $5,000\,{\rm au}$ )."
+ In the outer envelope of the core the density falls off as 77. and the outer boundary of the core is at 50.000au.," In the outer envelope of the core the density falls off as $r^{-4}$, and the outer boundary of the core is at $50,000\,{\rm au}\,$."
+ The total mass of the core is 5.4M.... of which ~2M. is in the kernel.," The total mass of the core is $5.4 M_{\odot}$, of which $\sim 2 M_\odot$ is in the kernel."
+ The core is initially isothermal. with temperature 7=IOK. and hence eq= 0.45.," The core is initially isothermal, with temperature $T = 10\,{\rm K}\,$, and hence $\alpha_{\rm therm} = 0.45\,$ ."
+ A. divergence-free Gaussian random velocity field. is superimposed on the core to simulate turbulence (cf., A divergence-free Gaussian random velocity field is superimposed on the core to simulate turbulence (cf.
+ Bate et al., Bate et al.
+ 2002a.b. 2003: Bonnell et al.," 2002a,b, 2003; Bonnell et al."
+ 2003)., 2003).
+ The power spectrum of the turbulence is set to PO)okot. which mimics the Larson sealing relation of Eqn. (1)) (," The power spectrum of the turbulence is set to $P(k) \propto k^{-4}$, which mimics the Larson scaling relation of Eqn. \ref{LARSON1}) ) ("
+Burkert Bodenheimer 2000).,Burkert Bodenheimer 2000).
+ The level of turbulent energy is set to Quin=0.05.," The level of turbulent energy is set to $\alpha_{\rm turb} = 0.05\,$."
+ Hence the core is initially globally virialised. in the sense that Ciern+Curb=0.5. although it is not in detailed hydrostatic balance.," Hence the core is initially globally virialised, in the sense that $\alpha_{\rm therm} + \alpha_{\rm turb} = 0.5$, although it is not in detailed hydrostatic balance."
+ In a subsequent paper we will examine the effect of varying the level of turbulent energy., In a subsequent paper we will examine the effect of varying the level of turbulent energy.
+ We present an ensemble of 20 simulations with these initial conditions., We present an ensemble of 20 simulations with these initial conditions.
+ The only difference between individual simulations is that the random number seed for the turbulent velocity field is changed - thereby changing the detailed structure of the velocity field. but hot its overall magnitude.," The only difference between individual simulations is that the random number seed for the turbulent velocity field is changed - thereby changing the detailed structure of the velocity field, but not its overall magnitude."
+ In addition we have run 10 simulations with no turbulence for the purpose of comparison: in each of these simulations the positional random number seed is changed so that the resulting Poisson noise in the initial particle positions is different., In addition we have run 10 simulations with no turbulence for the purpose of comparison; in each of these simulations the positional random number seed is changed so that the resulting Poisson noise in the initial particle positions is different.
+ The simulations are allowed to run for 0.3 Myr., The simulations are allowed to run for 0.3 Myr.
+ After 0.2 Myr most of the dynamical evolutionis finished. although aceretion is still on-going.," After 0.2 Myr most of the dynamical evolutionis finished, although accretion is still on-going."
+ At around 0.2 to 0.3 Myr. feedback," At around 0.2 to $0.3\,{\rm Myr}$ , feedback"
+The Mauritius Raclio Telescope (MICE) (Colapetal.1995a:UdavaShankaretal.2002) is à Fourier svnthesis. T-shaped non-coplanar array operating at 151.5 MlIA.,"The Mauritius Radio Telescope (MRT) \citep{ias:golap95a,apss:uday02} is a Fourier synthesis, T-shaped non-coplanar array operating at 151.5 MHz."
+ The telescope was built to fill the gap in the availability of deep sky surveys at low racio frequencies in the southern hemisphere., The telescope was built to fill the gap in the availability of deep sky surveys at low radio frequencies in the southern hemisphere.
+ The aim of the survey with MEINT is to contribute to the database of southern sky sources in the declination (9) range τί to 10 covering the entire right ascension (a). with a svnthesised beam of 4’4.Gsecze and an expected. point source sensitivity (1-7) of ~110 mJy +.," The aim of the survey with MRT is to contribute to the database of southern sky sources in the declination $\delta$ ) range $-70^\circ$ to $-10^\circ$, covering the entire right ascension $\alpha$ ), with a synthesised beam of $4\arcmin\times4\arcmin.6\sec za$ and an expected point source sensitivity $\sigma$ ) of $\sim110$ mJy $^{-1}$."
+" The (sa) is given by (9Ó). where. 6 (52O14"") is the latitucle of MICE."," The $(za)$ is given by $(\delta - \phi)$, where, $\phi$ $\left(\approx -20.14^\circ\right)$ is the latitude of MRT."
+ MICE. has been designed to be the southern-sky equivalent of the Cambridge 6€ survey at 151.5 MlIz (Baldwinetal.1985)., MRT has been designed to be the southern-sky equivalent of the Cambridge 6C survey at 151.5 MHz \citep{mnras:baldwin85}.
+. The next generation radio telescopes. like the LOw Frequeney Altirav (LOFAR) and the Murchison. Widelfield Array (AIWA). that are being built are low frequeney arrays: clearly indicating a renewed interest in metre-waveleneth astronomy.," The next generation radio telescopes, like the LOw Frequency ARray (LOFAR) and the Murchison Widefield Array (MWA), that are being built are low frequency arrays; clearly indicating a renewed interest in metre-wavelength astronomy."
+ “Phe Κον astrophysical science drivers. include acceleration. turbulence. and propagation in the galactic interstellar medium. exploring the high redshift universe anc transient phenomenon. as well as searching for the redshifte signature of neutral hvdrogen from the cosmologically important epoch of reionisation (I5olt).," The key astrophysical science drivers include acceleration, turbulence and propagation in the galactic interstellar medium, exploring the high redshift universe and transient phenomenon, as well as searching for the redshifted signature of neutral hydrogen from the cosmologically important epoch of reionisation (EoR)."
+ The surveys mace using such arrays will provide critical information abou oregrounds which will also provide a useful clatabase for roth extragalactic ancl galactic sources., The surveys made using such arrays will provide critical information about foregrounds which will also provide a useful database for both extragalactic and galactic sources.
+ MICE survey a 151.5 Mllz is a step in that direction and. in addition. wil wovicle the crucial sky mocel for calibration.," MRT survey at 151.5 MHz is a step in that direction and, in addition, will provide the crucial sky model for calibration."
+ Imaging at MICE is presently done only on the meridian o minimise the problems of non-coplanaritv., Imaging at MRT is presently done only on the meridian to minimise the problems of non-coplanarity.
+ X. two-dimensional (2D) image in a-sinze coordinates is formed w stacking one-dimensional (1-D) images on the meridian at clillerent sidereal times., A two-dimensional (2-D) image in $\alpha$ $\sin za$ coordinates is formed by stacking one-dimensional (1-D) images on the meridian at different sidereal times.
+ Images of ~ a steracian ish: of the southern sky. with an rns noise in MNimages of ⋅⋅300 my beam.71 (1-0). were produced," Images of $\sim$ a steradian $(18^{\mbox{\small h}}\leq\alpha\leq24^{\mbox{\small h}},-70^{\circ}\leq\delta\leq-10^{\circ})$ of the southern sky, with an rms noise in images of $\sim 300$ mJy $^{-1}$ $\sigma$ ), were produced"
+Current theoretical models predict. that galaxies. form and evolve in cold. dark matter (CDA) halos.,Current theoretical models predict that galaxies form and evolve in cold dark matter (CDM) halos.
+ Galaxies consequently tend to trace the distribution of mass. although he manner in which they do this may bebiased (Ixaiser1984:Dardeenοἱal.1986:Coles1993:Mo&White 1996).," Galaxies consequently tend to trace the distribution of mass, although the manner in which they do this may be \citep{Kaiser1984,BBKS,Coles1993,Mo1996}."
+. In principle. therefore. once the bias is allowed. for. it is xossible to use measurements of the clustering of galaxies o determine the clustering properties of the dark matter. especially if measurements can be made as a function of redshift.," In principle, therefore, once the bias is allowed for, it is possible to use measurements of the clustering of galaxies to determine the clustering properties of the dark matter, especially if measurements can be made as a function of redshift."
+ Decause clustering evolution is sensitive to the xwvameters underlying the background cosmological mocel. such observations can provide another (independent) test of he concordance cosmological model: see. e.g. Coles(2005).," Because clustering evolution is sensitive to the parameters underlying the background cosmological model, such observations can provide another (independent) test of the concordance cosmological model; see, e.g. \cite{Coles2005}."
+ In addition. clustering observations can be used to constrain ooperties of the galaxies themselves. such as the minimum alo mass within which they reside. which may vield clues about the processes of galaxy formation and evolution.," In addition, clustering observations can be used to constrain properties of the galaxies themselves, such as the minimum halo mass within which they reside, which may yield clues about the processes of galaxy formation and evolution."
+ A steadily increasing number of surveys of large-scale galaxy clustering are now becoming available., A steadily increasing number of surveys of large-scale galaxy clustering are now becoming available.
+ In the optical wavebands there are. projects such as the UIXIDSS Ultra Deep Survey (Llartleyetal.2010) and the SDSS Redshift Survey (Connollyetal.2010:Ross&Brunner2009) which are being used to extract information on clustering as a function of redshift.," In the optical wavebands there are projects such as the UKIDSS Ultra Deep Survey \citep{Hartley2010} and the SDSS Redshift Survey \citep{Connolly2010, Ross2009} which are being used to extract information on clustering as a function of redshift."
+ The Space Observatory was launched in 2009 and is the only space observatory to cover a spectral range from the far infrared. to. sub-millimetrc and therefore provides a new ancl unique window through which to study hieh-redshilt ewlaxy clustering., The Space Observatory was launched in 2009 and is the only space observatory to cover a spectral range from the far infrared to sub-millimetre and therefore provides a new and unique window through which to study high-redshift galaxy clustering.
+ Two surveys of particular interest to this article. HerMISS. (Oliverοἱal.2010). and L-ATLAS (Ealesetal.2010).. have already released angular correlation results (Coorayctal.2010:Aladdoxetal.2010) from their Science Demonstration Phase (SDP) which we will cliscuss further later. so it is timely to raise the issue of mocdelling sub-millimetre galaxy clustering.," Two surveys of particular interest to this article, HerMES \citep{Oliver2010} and H-ATLAS \citep{Eales2010}, have already released angular correlation results \citep{Cooray2010,Maddox2010} from their Science Demonstration Phase (SDP) which we will discuss further later, so it is timely to raise the issue of modelling sub-millimetre galaxy clustering."
+ Phe so-called Halo Mocel has been used. extensively over the past few vears in modelling galaxy clustering in a variety of contexts., The so-called Halo Model has been used extensively over the past few years in modelling galaxy clustering in a variety of contexts.
+ The. Lalo Model combines approximations of the dark matter profile within individual halos. the mass function. and bias models to estimate the correlation function for given cosmological parameters ancl characteristic halo masses.," The Halo Model combines approximations of the dark matter profile within individual halos, the mass function, and bias models to estimate the correlation function for given cosmological parameters and characteristic halo masses."
+ Lt has been shown to provide accurate ancl reliable predictions. of clustering measurements. at the price of a certain degree of complexiE and mocelling freedom.," It has been shown to provide accurate and reliable predictions of clustering measurements, at the price of a certain degree of complexity and modelling freedom."
+ The main focus of this paper is a comparison of the Llalo Alodel and a fitting function method which was initially. introduced. by Hamilton.etal.(1991)... ancl subsequently developed by Peacock&Dodds (19904).. Jain.Mo&White and especially Peacock&Doclels(19906).. to estimate the matter correlation Function.," The main focus of this paper is a comparison of the Halo Model and a fitting function method which was initially introduced by \cite{Hamilton1991}, and subsequently developed by \cite{Peacock1994}, , \cite{JMW95} and especially \cite{Peacock1996}, to estimate the matter correlation function."
+ Work by Matarrese(1997) followed by AMoscardinietal.(1998). and Colesetal. (1998)... showed how to incorporate this idea into a technique [or providing detailed: predictions of angular correlation in hieh-redshift galaxy surveys.," Work by \cite{Matarrese97} followed by \cite{Moscardini1998} and \cite{Coles1998}, showed how to incorporate this idea into a technique for providing detailed predictions of angular correlation in high-redshift galaxy surveys."
+ In this paperwe compare the, In this paperwe compare the
+In the case of the jet less dense than the ambient medium (hereafter “light jet”) which best reproduces observations. the integration domain extends over 300 AU in the radial direction and over 0000 AU in the + direction.,"In the case of the jet less dense than the ambient medium (hereafter “light jet”) which best reproduces observations, the integration domain extends over $300$ AU in the radial direction and over $6000$ AU in the $z$ direction."
+" In the case of the jet with the same mitial density as the ambient medium (hereafter ""equal-density jet”) which best reproduces observations. the domain Is (+ντ)m(600.6000) AU."," In the case of the jet with the same initial density as the ambient medium (hereafter “equal-density jet”) which best reproduces observations, the domain is $(r\times z) \approx
+(600\times6000)$ AU."
+ In this case. the radial axis is twice as large than in the light jet case because the cocoon surrounding the equal density jet has a radial extension greater than in the light jet.," In this case, the radial axis is twice as large than in the light jet case because the cocoon surrounding the equal density jet has a radial extension greater than in the light jet."
+ The dimension of the computational domain in the case of the jet denser than the ambient (hereafter “heavy jet”) which best reproduces observations 1s (τους27000) AU., The dimension of the computational domain in the case of the jet denser than the ambient (hereafter “heavy jet”) which best reproduces observations is $(r\times z) \approx (700 \times 27000)$ AU.
+" In all the cases. the initial jet velocity is along the + axis. coincident with the jet axis. and has a radial profile of the form where Vi ts the yeosh(r/r;)""on-axis velocity. 7 is the ambient to jet density ratio. η is the initial jet radius and αν=4 is the steepness parameter for the shear layer (as an example. see continuous line in Fig. |."," In all the cases, the initial jet velocity is along the $z$ axis, coincident with the jet axis, and has a radial profile of the form where $V_{0}$ is the on-axis velocity, $\nu$ is the ambient to jet density ratio, $r_{\rm j}$ is the initial jet radius and $w \bf= 4$ is the steepness parameter for the shear layer (as an example, see continuous line in Fig. \ref{fig1},"
+ for the light jet case discussed in Sect., for the light jet case discussed in Sect.
+ 3.2. with parameters in Tab. 3)).," \ref{Hydrodynamic evolution} with parameters in Tab. \ref{tab_mod}) ),"
+ to have a smooth transition of the kinetic energy at the interface between the jet and the ambient medium., to have a smooth transition of the kinetic energy at the interface between the jet and the ambient medium.
+ The density variation in the radial direction (dashed line in Fig. 1)), The density variation in the radial direction (dashed line in Fig. \ref{fig1}) )
+ ts where p is the jet densityrj (?))., is where $\rho_{\rm j}$ is the jet density \citealt{bmf94}) ).
+ Reflection boundary conditions are imposed along the jet axis. inflow boundary conditions at 2=0 and rzx+; and outflow boundary conditions elsewhere.," Reflection boundary conditions are imposed along the jet axis, inflow boundary conditions at $z = 0$ and $r \leq {r_{\rm j}}$ and outflow boundary conditions elsewhere."
+ The maximum spatial resolution achieved in the best light Jet case. (in both + and : directions). is z1.3 AU according to the PARAMESH methodology. using | refinement levels. corresponding to covering the jet radius with 25 points at the maximum resolution.," The maximum spatial resolution achieved in the best light jet case, (in both $r$ and $z$ directions), is $\approx 1.3$ AU according to the PARAMESH methodology, using $4$ refinement levels, corresponding to covering the jet radius with $25$ points at the maximum resolution."
+ The spatial resolution achieved in the equal-density case is half the one in the light jet case., The spatial resolution achieved in the equal-density case is half the one in the light jet case.
+ In the best heavy jet model. the spatial resolution achieved is 5 times lower than in the light jet case.," In the best heavy jet model, the spatial resolution achieved is $8$ times lower than in the light jet case."
+" Our choice of different spatial resolution for the three cases. aimed at reducing computational cost. was necessary because the thermal conduction is solved explicitly in FLASH and. therefore. a time-step limiter depending on density. p. temperature. 7. and spatial resolution, Ar. is required to avoid numerical instability (see. for instance. 2))."," Our choice of different spatial resolution for the three cases, aimed at reducing computational cost, was necessary because the thermal conduction is solved explicitly in FLASH and, therefore, a time-step limiter depending on density, $\rho$, temperature, $T$, and spatial resolution, $\Delta x$, is required to avoid numerical instability (see, for instance, \citealt{opr05}) )."
+" Stability is guaranteed for Aft<0.5Av?/D. where D is the diffusion coefficient. related to the conductivity. #. and to the specific heat at constant volume. ος, by D=#(1)/(pe,)."," Stability is guaranteed for $\Delta t <
+0.5~ \Delta x^2/D$, where $D$ is the diffusion coefficient, related to the conductivity, $\kappa$, and to the specific heat at constant volume, $c_v$, by $D = \kappa (T)/(\rho c_v)$."
+ In models characterized by high values of temperature (as. for instance. in the heavy jet case). therefore. a lower spatial resolution was required to avoid à very small time-step. Af.," In models characterized by high values of temperature (as, for instance, in the heavy jet case), therefore, a lower spatial resolution was required to avoid a very small time-step, $\Delta t$ ."
+ Condensations of plasma. due to radiative cooling effects. can become thermally unstable; however the presence of thermal conduction can prevent such instabilities.," Condensations of plasma, due to radiative cooling effects, can become thermally unstable; however the presence of thermal conduction can prevent such instabilities."
+" By comparing the radiative. 7,44. and thermal conduction. 7,4. characteristic times where / represents the characteristic length of temperature variations. we can infer which of the two competing processes dominates during the jet/ambient interaction."," By comparing the radiative, $\tau_{\rm rad}$, and thermal conduction, $\tau_{\rm cond}$ , characteristic times where $l$ represents the characteristic length of temperature variations, we can infer which of the two competing processes dominates during the jet/ambient interaction."
+ From the condition we can derive the cutoff length scale for instability. /p (2)). which indicates the maximum length over which thermal conduction dominates over radiative effects.in the classical conduction. regime.," From the condition we can derive the cutoff length scale for instability, $l_{\rm F}$ \citealt{fie65}) ), which indicates the maximum length over which thermal conduction dominates over radiative effects,in the classical conduction regime."
+ An analogous estimate in the saturation regime leads to which is one order of magnitude longer than the characteristic length 1n the classical regime., An analogous estimate in the saturation regime leads to which is one order of magnitude longer than the characteristic length in the classical regime.
+ As discussed later in Sect. 3.2..," As discussed later in Sect. \ref{Hydrodynamic
+evolution},"
+ the comparison between the classical Field length (the shortest characteristic length) and the size of the region behind the shock at the head of the jet will allow us to determine if this region is thermally stable or not., the comparison between the classical Field length (the shortest characteristic length) and the size of the region behind the shock at the head of the jet will allow us to determine if this region is thermally stable or not.
+ In order to verify our hypothesis of a fully tonized gas. we computed the ionization time scale of the most relevant elements in the X-ray spectrum of a shocked plasma at 7=Jobs10° Κ. assuming a post-shock density of about 10! ° (the light jet case).," In order to verify our hypothesis of a fully ionized gas, we computed the ionization time scale of the most relevant elements in the X-ray spectrum of a shocked plasma at $T = 3.4\times10^{6}$ K, assuming a post-shock density of about $10^{4}$ $^{-3}$ (the light jet case)."
+ As an example. we derived that the ionization time scale for C and O is | to 2 orders of magnitudes smaller than theradiative and thermal conduction time scales so that the plasma canbe considered in equilibrium.," As an example, we derived that the ionization time scale for C and O is 1 to 2 orders of magnitudes smaller than theradiative and thermal conduction time scales so that the plasma canbe considered in equilibrium."
+ ~1020% Ballet(2006))) ~109 (Dirurv(1983).. Jones&Ellison(1991).. Drury(2001))). (Amato&Blasi(2005))). (Auato&Blasi(2006). Viadimurov.Ellisπι«&DBvkov(2006))) ⋜⋯∖⋜↧↑∪≺∐↴∖↴↖↖⇁↕↑↕↓∪↴⋝↴∖↴↸∖↥⋅↖⇁⋜↧↑↕∪∐↴∖↴⋀∐⋃⋟∿−≽∩⋅↱⊐∩∙⊱⋯⊳∐↕⋜∐⋅∶↴∙⊾↸∖↴∖↴∐∪↸⊳↨↘↽∐↓∪≼∐∏↸⊳⋜↧⊓∪∐↴∖↴∐∪↖↖⇁↸∖↖↽↸∖↥⋅ ∫↿⋟⊤∿⊤↕∩∙↽∕∏∐∖ Πε (Warrenetal(2005))) O1 he backerouud plasma (Melkxeuzie&Volks(1982).. jereafter MISV82).," $\sim 10-20\%$ \cite{ballet}) $\sim 10^6$ \cite{druryrev}, \cite{je91}, \cite{maldrury}) \cite{amato1}) \cite{amato2} \cite{elli06}) $\Rt\sim M_0^{3/4} \sim 20-50$ $\Rt\sim 7-10$ $\Rt$ \cite{warren}) \cite{bv97} on the background plasma \cite{mck-v82}, hereafter MKV82)."
+ Iu fact this mechanism was orginally oxoposed in order to keep the AIF amplification iu the incar regime (e.g. ÓD/D< l1). but is now conuuonlv applied to cases 1u which 6B/Bx1.," In fact this mechanism was originally proposed in order to keep the MF amplification in the linear regime (e.g. $\delta B/B\ll 1$ ), but is now commonly applied to cases in which $\delta B/B\gg 1$."
+ Uufortunatelv. the icating process 1s quite model dependent aud even its applicability to situatious of interest for SNRs can aud should be seriously questioned.," Unfortunately, the heating process is quite model dependent and even its applicability to situations of interest for SNRs can and should be seriously questioned."
+ The effecIveness of the ienine process can easily be reduced to neglieible levels or artificially iuried to uuplivysical levels., The effectiveness of the heating process can easily be reduced to negligible levels or artificially amplified to unphysical levels.
+" As πιστοί above. a breakthrough iu the field jas recentlv beei provided bv XN-rav observations: the detection of Naw bright filaments in the outskirts of ποιο SNRs allow"" one to infer the streneth of the ME in these filamens. found to be B-—10050006."," As mentioned above, a breakthrough in the field has recently been provided by X-ray observations: the detection of X-ray bright filaments in the outskirts of some SNRs allows one to infer the strength of the MF in these filaments, found to be $B\sim 100-500\mu
+G$."
+ Such strong Πο]ls are generally attributed to the SI induced by CRs eficiently accelerated at the shock front. although. alternaive explanations have been proposed (Cdacalone&Jokipii (2007))).," Such strong fields are generally attributed to the SI induced by CRs efficiently accelerated at the shock front, although alternative explanations have been proposed \cite{gj07}) )."
+ Tn Tab., In Tab.
+ 1 we list some SNRs with estiaated MEs: αμ is the shock velocity. D» is the ME clownsrea ofthe shock as inferred from the N-ray bniehtuess profile aud Py2=BS(Szpou) is the downstream magnetic pressure normalized to the bulk one.," \ref{tab} we list some SNRs with estimated MFs: $u_0$ is the shock velocity, $B_2$ is the MF downstream of the shock as inferred from the X-ray brightness profile and $P_{w2}=B_2^2/(8\pi\rho_0 u_0^2)$ is the downstream magnetic pressure normalized to the bulk one."
+ The values of the parameters are from Parizotet(2006). aud from VBINOS (in pareuthesis iu Tab. 1))., The values of the parameters are from \cite{P+06} and from VBK05 (in parenthesis in Tab. \ref{tab}) ).
+ Weshow belowthat for he field streeths inferred for SNRs. the magnetic pressre is comparableor evenimexcess of the therlal pressire of the background plasma," Weshow belowthat for the field stregths inferred for SNRs, the magnetic pressure is comparableor eveninexcess of the thermal pressure of the background plasma"
+lmprovements in precision radial velocity (RV) measurements anc continued monitoring— are permitting the detection of lower amplitude planetary signatures.,Improvements in precision radial velocity (RV) measurements and continued monitoring are permitting the detection of lower amplitude planetary signatures.
+ One example of the fruits of this work is the detection of a super earth in the habitable zone surrounding Gliese 581 (Ucdryοἱal. 2007)., One example of the fruits of this work is the detection of a super earth in the habitable zone surrounding Gliese 581 \citep{Udry2007}.
+. Phis and other remarkable successes on the part of the observers is motivating a significant ellort to improve the statistical tools for analvzing radial velocity data (e.g. Loredo&Chernoll2003... Loreco 2004... Cumming. 2004.. Gregory 2005a b. Ford 2005 2006. Ford&Gregory 2006.. Cumming&Dragomir 2010)).," This and other remarkable successes on the part of the observers is motivating a significant effort to improve the statistical tools for analyzing radial velocity data (e.g., \citealt{LoredoChernoff2003}, \citealt{Loredo2004}, \citealt{Cumming2004}, Gregory 2005a b, Ford 2005 2006, \citealt{FordGregory2006}, \citealt{CummingDragomir2010}) )."
+ Much of the recent work has highlighted a Bayesian. AICAIC approach as a wav to better understand. parameter. uncertainties and degeneracies and to compute model probabilities., Much of the recent work has highlighted a Bayesian MCMC approach as a way to better understand parameter uncertainties and degeneracies and to compute model probabilities.
+ Gregory (2005a. b. c and 2007a. b. c) presented. a Bavesian MCAIC algorithm that makes use of parallel tempering (PT) to ellicientlv explore a large. moclel," Gregory (2005a, b, c and 2007a, b, c) presented a Bayesian MCMC algorithm that makes use of parallel tempering (PT) to efficiently explore a large model"
+"We now introduce the following canonical change of variables, of multiplier πο, n being the mean orbital motion of Mercury: In order to be consistent with the sign minus in the equations and before /., the wobble of the pseudo-core J. has to be replaced by π—Je.","We now introduce the following canonical change of variables, of multiplier $\frac{1}{nC}$ , $n$ being the mean orbital motion of Mercury: In order to be consistent with the sign minus in the equations and before $l_c$, the wobble of the pseudo-core $J_c$ has to be replaced by $\pi-J_c$."
+" In this way, we have Lc=G.cos(m—Jc)—G.cos(Jc)."," In this way, we have $L_c=G_c\cos(\pi-J_c)=-G_c\cos(J_c)$."
+" In this new setof variables, we have"," In this new setof variables, we have"
+FranceAfrica The hallmark of planetary microlensing events is a short deviation from an otherwise normal. point-source/point-lens (hereafter PSPL) event.," 	The hallmark of planetary microlensing events is a short deviation from an otherwise normal, point-source/point-lens (hereafter PSPL) event."
+ showed that extrasolar planets could be detected from such events. and gave an explicit prescription for how the planet/star mass ratio g (« 1) and the angular separation ¢ (1n units of the angular Einstein radius (c) could be reconstructed by decomposing the event light curve into its “normal” and “perturbed” components.," showed that extrasolar planets could be detected from such events, and gave an explicit prescription for how the planet/star mass ratio $q$ $\ll 1$ ) and the angular separation $d$ (in units of the angular Einstein radius $\theta_{\rm E}$ ) could be reconstructed by decomposing the event light curve into its “normal” and “perturbed” components."
+ Work during the ensuing decade has elucidated many additional subtleties of planetary light curves. but their fundamental characterization as briefly perturbed PSPL events has remained intact.," 	Work during the ensuing decade has elucidated many additional subtleties of planetary light curves, but their fundamental characterization as briefly perturbed PSPL events has remained intact."
+" Of particular importance in the present context. showed that events with small impact parameter (0j«1: where vy is the minimum separation between the source and the lens center of mass m units of (/r) probe the so-called ""central caustic” of the lens geometry. making them much more sensitive to the presence of planets than the larger impact-parameter events analyzed by(1992).. which probe the outer “planetary caustic”."," Of particular importance in the present context, showed that events with small impact parameter $u_0\ll 1$; where $u_0$ is the minimum separation between the source and the lens center of mass in units of $\theta_{\rm E}$ ) probe the so-called “central caustic” of the lens geometry, making them much more sensitive to the presence of planets than the larger impact-parameter events analyzed by, which probe the outer “planetary caustic”."
+ These central-caustic events are of exceptional importance. even though they are intrinsically rare.," These central-caustic events are of exceptional importance, even though they are intrinsically rare."
+ They are rare simply because the central-caustie is much smaller than the planetary caustic. so the great majority of planet-induced deviations (of fixed fractional amplitude) are due to planetary caustics.," They are rare simply because the central-caustic is much smaller than the planetary caustic, so the great majority of planet-induced deviations (of fixed fractional amplitude) are due to planetary caustics."
+ However. the probability of detecting a planet is much greater in small events. partly because the source is guaranteed to pass close to the central caustic. which almost by definition is near the center of the lens geometry (t= 0) and partly because even the sensitivity of planetary caustics is enhanced for wy« 1. (," However, the probability of detecting a planet is much greater in small impact-parameter events, partly because the source is guaranteed to pass close to the central caustic, which almost by definition is near the center of the lens geometry $\vect{u}=0$ ) and partly because even the sensitivity of planetary caustics is enhanced for $u_0\ll 1$ . ("
+Here u denotes the source position on the sky. normalized to p. with respect to the lens center of mass.),"Here $\vect{u}$ denotes the source position on the sky, normalized to $\theta_{\rm E}$, with respect to the lens center of mass.)"
+ By contrast. higher events miss the central caustic. and they are likely to miss the planetary caustic as well because it lies in a random position relative to the source trajectory.," By contrast, higher impact-parameter events miss the central caustic, and they are likely to miss the planetary caustic as well because it lies in a random position relative to the source trajectory."
+ Because of their higher sensitivity to planets. and because they can be recognized in real time. low impact-parameter events are monitored more intensively than typical events by microlensing follow-up networks. which in turn further enhancestheir sensitivity to planets.," Because of their higher sensitivity to planets, and because they can be recognized in real time, low impact-parameter events are monitored more intensively than typical events by microlensing follow-up networks, which in turn further enhancestheir sensitivity to planets."
+The Jorcan-Brans-Dicke (JBD) (Brans&Dicke1961) theory is one of the earliest and most well-motivated alternatives to the Einstein theory of eravitation.,The Jordan-Brans-Dicke (JBD) \citep{brans} theory is one of the earliest and most well-motivated alternatives to the Einstein theory of gravitation.
+ Phe value of the gravitational “constant” is set by the inverse ofa time-dependent classical scalar field witha coupling parameter w., The value of the gravitational “constant” is set by the inverse ofa time-dependent classical scalar field witha coupling parameter $\omega$.
+ General relativity is recovered in the limit of w7x., General relativity is recovered in the limit of $\omega \to \infty$.
+ Solar system observations impose lower bounds on w (Bertottiοἱal.2003:Will 2006).," Solar system observations impose lower bounds on $\omega$ \citep{obser1,obser}."
+. Generalized JBD models are obtained in the low energy limit of higher dimensional heories., Generalized JBD models are obtained in the low energy limit of higher dimensional theories.
+ String theoretic (Gasperini&Veneziano2007) and. Ixaluza-Ixlein (Appelquistetal.1987) models after conpacification of the extra. cimensions viele several variants of JBL mocels (or eencral scalar-tensor models) in which the the scalar field coupling w may become dynamical. and also models. with potential terms for the JBD scalar field.," String theoretic \citep{string} and Kaluza-Klein \citep{kk} models after compactification of the extra dimensions yield several variants of JBD models (or general scalar-tensor models) in which the the scalar field coupling $\omega$ may become dynamical, and also models with potential terms for the JBD scalar field."
+ JBD models have been used for tackling several cosmological problems pertaining to dilferent eras of evolution of the Fricclman-Robertson-Walker (PFRW) universe., JBD models have been used for tackling several cosmological problems pertaining to different eras of evolution of the Friedman-Robertson-Walker (FRW) universe.
+ The scenario of extended inllation (Mathiazhagan&Johri1984:LaSteinhardt.1989) was ooposed within the context. of JBD cosmology.," The scenario of extended inflation \citep{extended1,extended2} was proposed within the context of JBD cosmology."
+ Generalized 191). nocels resulting from the compactification of higher dimensional actions (Alajumdar&Sethi1992:Majumcdarctal.1993:Alajumelar1997) enable eflicient transition to the post-inflationary radiation clominatec phase (or reheating).," Generalized JBD models resulting from the compactification of higher dimensional actions \citep{asm11,asm12,asm13} enable efficient transition to the post-inflationary radiation dominated phase (or reheating)."
+ The JBD scalar field has also been incorporated in quintessence scenarios for obtaining the present acceleration. of the universe (Bertolami&Martins2000:BanerjeePavon2001:Sen&Seshadri 2003).," The JBD scalar field has also been incorporated in quintessence scenarios for obtaining the present acceleration of the universe \citep{quint1,quint2,quint3}."
+.. A more general class of JBD theory where the coupling parameter w is an arbitrary function. of the scalar ield (Bergmann1968:Wagoncr1970:Bergmann1970) has interesting consequences on cosmic evolution (Sahoo&Singh2002.2003) and has also been applied. in the conext of obtaining singularity [ree cosmology (IxalvanaRamaLOOT) and relic gravitational waves (Barrowctal.1993:Singh&Sahoo2004).," A more general class of JBD theory where the coupling parameter $\omega$ is an arbitrary function of the scalar field \citep{bergmann1,bergmann2,bergmann3} has interesting consequences on cosmic evolution \citep{sahoo1,sahoo2} and has also been applied in the context of obtaining singularity free cosmology \citep{kalyana} and relic gravitational waves \citep{barrow,sahoo3}."
+. Primorcial black roles (PBIIs) are potentially [ascinating cosmo-archeological tools (Carr 2003).., Primordial black holes (PBHs) are potentially fascinating cosmo-archeological tools \citep{pbh}. .
+ Lt is possible for PBs to impact through their evaporation products diverse cosmological processes such as barvogenesis ancl nucleosvnthesis. (Care1976:Zeldovichelal.1977:1991:Majumcdaretal.1995:LiddleandGreen 1998).. the cosmic microwave background. radiation. (Chapline1975:AlacGibbon&Carr 1991).. and the growth of perturbations as well CXfshordiοἱal.," It is possible for PBHs to impact through their evaporation products diverse cosmological processes such as baryogenesis and nucleosynthesis \citep{baryo1,baryo2,baryo3,baryo7,baryo4,baryo5,baryo6}, the cosmic microwave background radiation \citep{cmbr1,cmbr2}, and the growth of perturbations as well \citep{afshordi}."
+2003).. PBLIs could act as seeds for structure formation (Macketal.2006) and could also form a significant component of darκ. matter through ellicient early accretion in. braneworlel scenarios (Alajuniclar2003:Alukherjee 2005)..," PBHs could act as seeds for structure formation \citep{ostriker} and could also form a significant component of dark matter through efficient early accretion in braneworld scenarios \citep{brane1,brane2,brane3,brane4}. ."
+" Mechanisms or growth of supermassive black holes by PBIIs accreting dark. cnerey have been proposed (Bean&Magueijo20 )2).. though recent. results bv ανασα&Carr(2005a) and Llaraclaetal.(2006) indicate the lack of self-similar erowth of black holesby accreting quintessence,"," Mechanisms for growth of supermassive black holes by PBHs accreting dark energy have been proposed \citep{bean}, , though recent results by \citet{harada11} and \citet{harada12} indicate the lack of self-similar growth of black holesby accreting quintessence."
+Long duration Gammarav bursts (CBs) are thought to originate [rom the collapse of massive stars.,Long duration Gamma–ray bursts (GRBs) are thought to originate from the collapse of massive stars.
+ Several lines of evidence points toward. this conclusion. ranging [rom the association to type Ib/c supernovae (Woosley Bloom 2006 and references therein). to the occurrence of GiBs in the most luminous part of their host galaxies (Svensson et al.," Several lines of evidence points toward this conclusion, ranging from the association to type Ib/c supernovae (Woosley Bloom 2006 and references therein), to the occurrence of GRBs in the most luminous part of their host galaxies (Svensson et al."
+ 2010)., 2010).
+ The ambient. medium in which GBs explode is expected. to be denser than the interstellar medium and typical of star forming regions., The ambient medium in which GRBs explode is expected to be denser than the interstellar medium and typical of star forming regions.
+ The values of the absorbing column densities as measured in Xrays are hieh (Calama Wijers 2001: Stratta et al., The values of the absorbing column densities as measured in X–rays are high (Galama Wijers 2001; Stratta et al.
+ 2004: Campana ct al., 2004; Campana et al.
+ 2006: Watson et al., 2006; Watson et al.
+ 2007)., 2007).
+ An analysis of the intrinsic column densities of allοι GRBs observed within 1.000 s and with a known redshift has been carried out by Campana ct al. (," An analysis of the intrinsic column densities of all GRBs observed within 1,000 s and with a known redshift has been carried out by Campana et al. ("
+2010).,2010).
+ The selected sample consisting of 93 GRBs was biased., The selected sample consisting of 93 GRBs was biased.
+ The cistribution of the intrinsic N.rav absorption column clensity is consistent with a lognormal distribution with mean logeNy(s)=219+0.1 (90% confidence level)., The distribution of the intrinsic X–ray absorption column density is consistent with a lognormal distribution with mean $\log N_H(z)=21.9\pm0.1$ $90\%$ confidence level).
+ This distribution is in agreement with the expected distribution for GRBs occurring randomly in giant molecular clouds similar to those within the Milkv Way (Campana et al., This distribution is in agreement with the expected distribution for GRBs occurring randomly in giant molecular clouds similar to those within the Milky Way (Campana et al.
+ 2006: Reichart Price 2002)., 2006; Reichart Price 2002).
+ Looking at he cistribution ofXrav column densities vs. redshift. there is à lack of non-absorbed CRBs at high redshift and a lack of heavily absorbed CRBs at low redshift.," Looking at the distribution of X–ray column densities vs. redshift, there is a lack of non-absorbed GRBs at high redshift and a lack of heavily absorbed GRBs at low redshift."
+ This might. be he outcome of biases present in the sample., This might be the outcome of biases present in the sample.
+ Looking a he distribution of Xray column clensities versus recdshill a lack of non-absorbed GRBs at high redshift ancl a lack of heavily absorbed ας at low redshift were found. in oevious studies (Campana ct al., Looking at the distribution of X–ray column densities versus redshift a lack of non-absorbed GRBs at high redshift and a lack of heavily absorbed GRBs at low redshift were found in previous studies (Campana et al.
+ 2010)., 2010).
+ The former migh »' explained in terms of more compact and dense star ormation regions in the voung Universe (or to a sizable contribution [from intervening svstems)., The former might be explained in terms of more compact and dense star formation regions in the young Universe (or to a sizable contribution from intervening systems).
+ The latter migh » interpreted. as due to a change in the dust. properties with redshift. with GRBs at redshift z223 having a uigher dust to gas ratio for the same X.ray column densiE (c.g. dillerent grain size or composition).," The latter might be interpreted as due to a change in the dust properties with redshift, with GRBs at redshift $z\lsim 2-3$ having a higher dust to gas ratio for the same X–ray column density (e.g. different grain size or composition)."
+ This will naturally wovicle a lack of heavily (X.rav) absorbed. GRBs at small redshifts., This will naturally provide a lack of heavily (X–ray) absorbed GRBs at small redshifts.
+ In the optical the presence. of à large. amount. of absorbing material is much less clear. since a large number of CRB afterglows are not alfected by absorption (Ixann. Klose Zeh 2006: Sehady et al.," In the optical the presence of a large amount of absorbing material is much less clear, since a large number of GRB afterglows are not affected by absorption (Kann, Klose Zeh 2006; Schady et al."
+ 2010: Zafar et al., 2010; Zafar et al.
+ 2011: but see Greiner et al., 2011; but see Greiner et al.
+ 2011 and Covino et al., 2011 and Covino et al.
+ 2011. in preparation).," 2011, in preparation)."
+Dynamical studies in low mass X-rav binaries (LAINBs hereafter) offer a promissing route to test the equation of state of nuclear matter.,Dynamical studies in low mass X-ray binaries (LMXBs hereafter) offer a promissing route to test the equation of state of nuclear matter.
+ Soft equations of state are not stable above 1.6 M. (e.g.Brown&Bethe1994). and hence finding a neutron star more massive than this limit would be a major advance in our knowledge of nuclear matter physics., Soft equations of state are not stable above 1.6 $_{\odot}$ \citep[e.g.][]{brown94} and hence finding a neutron star more massive than this limit would be a major advance in our knowledge of nuclear matter physics.
+ LAINBs and. in particular. accreting millisecond pulsars (AALS) are expected to harbour the most massive neutron stars oecause of the sustained. acerction of matter curing their ong lifes (vandenHeuvel&Bitzaraki1995).," LMXBs and, in particular, accreting millisecond pulsars (AMPs) are expected to harbour the most massive neutron stars because of the sustained accretion of matter during their long lifes \citep{vandenheuvel95}."
+. Unfortunately. dynamical studies are hampered by (1) the overwhelming accretion Luminosity in persistent LAINBs (which swanips he conor star's spectrum) and. (2) the extreme faintness of the companion star in transient AALPs during quicseenee (D'Avanzoetal.2009).," Unfortunately, dynamical studies are hampered by (1) the overwhelming accretion luminosity in persistent LMXBs (which swamps the donor star's spectrum) and (2) the extreme faintness of the companion star in transient AMPs during quiescence \citep{davanzo09}."
+. A promissing route to overcome these limitations was oposed by Steeghs&Casares(2002)., A promissing route to overcome these limitations was proposed by \cite{steeghs02}.
+.. X-ray irracliated donor stars can be betraved. by the detection of high excitation lluorescence NILE and CLL lines within the Bowen end at 224634-50A., X-ray irradiated donor stars can be betrayed by the detection of high excitation fluorescence NIII and CIII lines within the Bowen blend at $\simeq$ 4634-50.
+ Phe Doppler shift of these narrow lines races the orbit of the donor star and. therefore. can vield dynamical constraints in persistent LAINBs ancl transient AAIPs during outburst.," The Doppler shift of these narrow lines traces the orbit of the donor star and, therefore, can yield dynamical constraints in persistent LMXBs and transient AMPs during outburst."
+ As a result of this strateey some &ood cancidates for massive neutron stars have been found. namely N1822-371 (Munoz-dariasetal.2005). and λα N-1 (Cornelisseetal.2007).," As a result of this strategy some good candidates for massive neutron stars have been found, namely X1822-371 \citep{munoz05} and Aql X-1 \citep{cornelisse07}."
+. In a few long period. (P?z 1 day) systems. the donor is an evolved star which can be spectroscopically detected over the ipraciated accretion disc.," In a few long period $P>$ 1 day) systems, the donor is an evolved star which can be spectroscopically detected over the irradiated accretion disc."
+ Cvg N-2 represents one of this exceptional cases., Cyg X-2 represents one of this exceptional cases.
+ Is first orbital solution was presented bv Cowleyetal.(1979). who show that the donor star is an AD-P2 HE orbiting the neutron star in 9.843 days with a projected velocity Avs=ST43 kms +., Its first orbital solution was presented by \cite{cowley79} who show that the donor star is an A5-F2 III orbiting the neutron star in 9.843 days with a projected velocity $K_2=87\pm3$ km $^{-1}$.
+ Almost 20 vears later. this was refined by Casaresetal.(1998). (COS) who find an A9 donor with orbital parameters 214=9.8444+0.0003 cd and vo= SS.0£14kms +.," Almost 20 years later, this was refined by \cite{casares98} (C98) who find an A9 donor with orbital parameters $P_{\rm orb}=9.8444 \pm 0.0003$ d and $K_2=88.0\pm1.4$ km $^{-1}$."
+ This work also reports the first determination of the rotational broadening of the conor’s absorption features (Vsin;=342+2.5 kms 1) which. in turn. implies a binary mass ratio g=AleΑι0.34d0.04.," This work also reports the first determination of the rotational broadening of the donor's absorption features $V \sin i = 34.2 
+\pm 2.5$ km $^{-1}$ ) which, in turn, implies a binary mass ratio $q=M_{2}/M_{1}=0.34 \pm 0.04$."
+ This value is remarkable as it implies a peculiar donor star. very undermassive for its spectral type ancl the probable outcome of an intermediate mass binary (lxing&Ritter1999:Pocdsiadlowski&Rappaport2000:Ixolbetal. 2000).," This value is remarkable as it implies a peculiar donor star, very undermassive for its spectral type and the probable outcome of an intermediate mass binary \citep{king99, podsiadlowski00, kolb00}."
+. jut the implications for the compact objects mass are also remarkable., But the implications for the compact object's mass are also remarkable.
+ Ehe orbital solution. combined with inclination constraints /=62.5d4 derived through ellipsoidal fits to UDV light curves gives à neutron star mass of 1.78 + 0.23 M. (Orosz&Ixuulkers1990). and. hence. a good candidate for a massive neutron star.," The orbital solution, combined with inclination constraints $i=62.5\pm4^{\circ}$ derived through ellipsoidal fits to UBV light curves gives a neutron star mass of 1.78 $\pm$ 0.23 $_{\odot}$ \citep{orosz99} and, hence, a good candidate for a massive neutron star."
+ This result. has been recently challenged. by Elebertctal.(2009) (E209). who measure ο=7943 kms Land hence Aljp=l.5 £ 0.3 M., This result has been recently challenged by \cite{elebert09} (E09) who measure $K_2=79\pm3$ km $^{-1}$ and hence $M_1$ =1.5 $\pm$ 0.3 $_{\odot}$.
+ llere we present new high resolution spectroscopic observations of (νο N-2 with the main aim of refining the rotational broadening of the donor star and update the orbital parameters., Here we present new high resolution spectroscopic observations of Cyg X-2 with the main aim of refining the rotational broadening of the donor star and update the orbital parameters.
+ In. contrast with 1209. we find a good agreement with previous results reported in C98 which give support to the presence of a massive neutron star in (νο X-2.," In contrast with E09, we find a good agreement with previous results reported in C98 which give support to the presence of a massive neutron star in Cyg X-2."
+"[a1732"" 0x1160 —24 V Lewinetal.(1971). itis 1976).. (Dotvetal.1981:Rickettsal.1983).","$\alpha$$^h$$^m$ $\delta$ $-$ $V$ \citet{letal71} \citet{d76}
+ \citep{betal76}. \citep{detal81, retal82,
+metal88}."
+. Bo 33x 10PG. (Cui1997).. (Predehletal.1995), $_\circ$ $\sim$ $\times$ $^{13}$ \citep{c97}.
+ (Jablonskietal.1997).. Glass&," \citep{pfs95}
+ \citep{jetal97}."
+"Feast(1973) A, ~ (Davidsenοἱal.1977:Shahbazet1996: (Chakrabarty&"," \citet{gf73} $A_v$ $\sim$ \citep{detal77, setal96, cr97}. \citep{cr97}."
+Roche1997).. (IKkholopovetal.1981). 73560R... (—G6x10Davidsenοἱal.(1977) ," \citet{detal77} \citep{ketal81} $\sim$$_\odot$ $\sim$$\times$$^{8}$\citet{detal77}
+ "
+agreement with the measured ones: equivalently one could raise the value of the critical density. ὃς required for collapse above that of the top-hat model.,"agreement with the measured ones; equivalently one could raise the value of the critical density, $\delta_c$, required for collapse above that of the top-hat model."
+ Alore serious is variance of x. approximately 0.2. which means that two halos with the same assigned density can have quite disparate true overdensities.," More serious is variance of $\chi$, approximately 0.2, which means that two halos with the same assigned density can have quite disparate true overdensities."
+ The model supposes that they collapse zt. the. same time. whereas the ful non-linear evolution would presumably show otherwise.," The model supposes that they collapse at the same time, whereas the full non-linear evolution would presumably show otherwise."
+ We occasionally find some high-mass halos (mass greater than 8) with big x. which contribute significantly to the enlargemen of the error bars at. those scales.," We occasionally find some high-mass halos (mass greater than 8) with big $\chi$, which contribute significantly to the enlargement of the error bars at those scales."
+ Phese are cllectively leftovers of the merging process anc should not be treater as collapsed halos in subsequent applications of the methoc (that is. in a realistic galaxy formation modelling they should be considered as sources of material to be accrete at a later stage of the hierarchy).," These are effectively leftovers of the merging process and should not be treated as collapsed halos in subsequent applications of the method (that is, in a realistic galaxy formation modelling they should be considered as sources of material to be accreted at a later stage of the hierarchy)."
+ The varianee in X is unmweleome but is only one contribution to the dispersion. between overdensity anc collapse time., The variance in $\chi$ is unwelcome but is only one contribution to the dispersion between overdensity and collapse time.
+ We note that N-bocly simulations show [ου cach particle a poor correspondence between the expecte mass ofits parent halo (predicted from the initial conditions) and the true value measured from a numerical simulation. evolved from the same initial conditions (White 1995. Bond 11992).," We note that N-body simulations show for each particle a poor correspondence between the expected mass of its parent halo (predicted from the initial conditions) and the true value measured from a numerical simulation, evolved from the same initial conditions (White 1995, Bond 1992)."
+ Moreover. gravitational collapse is clearly not. as simple as the spherical model assumes., Moreover gravitational collapse is clearly not as simple as the spherical model assumes.
+ Phere is. for example. no guarantee that underdense regions never collapse or that high-dense regions will do so (Dertschinger Jain 1904).," There is, for example, no guarantee that underdense regions never collapse or that high-dense regions will do so (Bertschinger Jain 1994)."
+ However. if a simple semi-analytical mocdel of the gravitational clustering is desired. then the simple relation between collapse redshift and. initial overdensity given by the spherical model seems the most obvious choice.," However, if a simple semi-analytical model of the gravitational clustering is desired, then the simple relation between collapse redshift and initial overdensity given by the spherical model seems the most obvious choice."
+ Our method passes the test for self-similaritv. vet for n=0 it predicts far more high-mass halos than Press-Schechter or other methods based. on similar ideas. such as theMocdel.," Our method passes the test for self-similarity, yet for $n=0$ it predicts far more high-mass halos than Press-Schechter or other methods based on similar ideas, such as the."
+. This is an expected. outcome of our method. and points more to a deficiency in the PS model than anything else. as we attempt to explain below.," This is an expected outcome of our method and points more to a deficiency in the PS model than anything else, as we attempt to explain below."
+ Press-Schechter does not count the number of halos of a given mass., Press-Schechter does not count the number of halos of a given mass.
+ Rather. it counts the fraction of the Universe where. if one were to put down a top-hat filter of the appropriate mass. the overdensity. would exceed a certain eritical threshold.," Rather, it counts the fraction of the Universe where, if one were to put down a top-hat filter of the appropriate mass, the overdensity would exceed a certain critical threshold."
+ Regions which just poke above this hreshold. for a single position of their centres. contribute nothing to the mass-function.," Regions which just poke above this threshold for a single position of their centres, contribute nothing to the mass-function."
+ It is easy to see that this xcomes increasingly likely as one moves to rarer and rarer objects (of higher ancl higher mass)., It is easy to see that this becomes increasingly likely as one moves to rarer and rarer objects (of higher and higher mass).
+ For these it is much ter simply to count the number of peaks which exceed he threshold density after filtering on the appropriate scale (on the other hand Peaks Theory. predicts far too many low-mass objects as it does not distinguish between overlapping jos)., For these it is much better simply to count the number of peaks which exceed the threshold density after filtering on the appropriate scale (on the other hand Peaks Theory predicts far too many low-mass objects as it does not distinguish between overlapping halos).
+ Phe necessary theory has been exhaustively analvzed » Bardeen (1986) who showed that uncorrected. PS (without the extra factor of two) underestimates the number of high-mass halos by a factor i30n7/77. where 9.=ra(m) (this result is lor a gaussian filter but similar results will wold for all filters with just a small cillerence in scaling).," The necessary theory has been exhaustively analyzed by Bardeen (1986) who showed that uncorrected PS (without the extra factor of two) underestimates the number of high-mass halos by a factor $\alpha^{3/2}\nu^3$, where $\delta_c=\nu\sigma(m)$ (this result is for a gaussian filter but similar results will hold for all filters with just a small difference in scaling)."
+ One way to visualize this result is to think of cach peak as jwing an overdensity profile where 2 is the radius of the top-hat filter., One way to visualize this result is to think of each peak as having an overdensity profile where $R$ is the radius of the top-hat filter.
+ Ht is then easy o estimate the contribution to the PS mass fraction and o integrate over all values of £j greater than the threshold. δισ. τοσο the total number of halos.," It is then easy to estimate the contribution to the PS mass fraction and to integrate over all values of $\nu_0$ greater than the threshold, $\delta_c/\sigma$, to get the total number of halos."
+ This method suggests hat PS should. predict V2/Orv? times as many halos as oaks VPheory. in rough agreement with the above [ος p=1 Oo Q0.," This method suggests that PS should predict $\sqrt{2/9\pi}\,\nu^3$ times as many halos as Peaks Theory, in rough agreement with the above for $n\approx -1$ to 0."
+ A more direct. demonstration of the above cillerence tween. Press-Schechter and the actual number of high-mass peaks in the density field is shown by the numbers in ‘Table 1.., A more direct demonstration of the above difference between Press-Schechter and the actual number of high-mass peaks in the density field is shown by the numbers in Table \ref{tab:npeaks}.
+ Columns 29 show the measured number of blocks which exceed the density threshold given in the first column in cach of the eight sub-grids of mass 512., Columns 2–9 show the measured number of blocks which exceed the density threshold given in the first column in each of the eight sub-grids of mass 512.
+ These agree with the Press-Schechter prediction. as," These agree with the Press-Schechter prediction, as"
+at the candidate positions.,at the candidate positions.
+" Last, we have now approached a regime where there are enough known classes of sources (from SNe to variable star types) that reliable cross-validated classification can be employed to run machine-learned classifications instead of the manually tuned classification algorithm (§3))."," Last, we have now approached a regime where there are enough known classes of sources (from SNe to variable star types) that reliable cross-validated classification can be employed to run machine-learned classifications instead of the manually tuned classification algorithm \ref{sec:class}) )."
+" It is clear from §3.4 that ML-based classifications are reasonably predictive, withTRANSIENT/VARSTAR Classification errors at the level."," It is clear from \ref{sec:ml} that ML-based classifications are reasonably predictive, with classification errors at the level."
+ It is clearly early days for large-scale discovery and classification frameworks for synoptic astronomical surveys., It is clearly early days for large-scale discovery and classification frameworks for synoptic astronomical surveys.
+" As we look to future implementations, there are several avenues and questions to explore:"," As we look to future implementations, there are several avenues and questions to explore:"
+field.. and y?=&2/(67 K2).,", and $\gamma^2=k_z^2/(k_y^2+k_z^2)$."
+" Channel modes have &,=0 so that y=I.", Channel modes have $k_y=0$ so that $\gamma=1$.
+ This equation is used later on to compute the channel mode growth rates in the conditions of our numerical simulations., This equation is used later on to compute the channel mode growth rates in the conditions of our numerical simulations.
+ It is most conveniently solved in dimensionless form. with S as unit of time and V4/S as unit of length.," It is most conveniently solved in dimensionless form, with $S$ as unit of time and $V_A/S$ as unit of length."
+" We introduce σ΄. C/S. Kk= R/S. ko VAk/S; defining the viscous (E) and resistive (E,) Elsasser-like numbers by one has Siv=p[Εν anddq y=1/£,."," We introduce $\sigma^*=\sigma/S$ , $\kappa^*=\kappa/S$ , $k^*=V_A k/S$; defining the viscous $E_\nu$ ) and resistive $E_\eta$ ) Elsasser-like numbers by one has $\nu^*=1/E_\nu$ and $\eta^*=1/E_\eta$."
+ Note that with a power-law velocity profile in à disk (Q 77). Q*=1l/q and αἲΞ[202-aq]7/q.," Note that with a power-law velocity profile in a disk $\Omega\propto
+r^{-q}$ ), $\Omega^*=1/q$ and $\kappa^*=[2(2-q)]^{1/2}/q$."
+" Our definitions of the Elsasser numbers (and other dimensionless quantities) differ from the usual ones (A, and A,. see e.g. ? and ?)) by a factor $/O."," Our definitions of the Elsasser numbers (and other dimensionless quantities) differ from the usual ones $\Lambda_\nu$ and $\Lambda_\eta$, see e.g. \citealt{PC08} and \citealt{PG09}) ) by a factor $S/\Omega$."
+ This choice has several motivations: S rather than © is the relevant time-scale of linear growth rates: this definition is more consistent with our previous choice of units: and it makes the connection between the Elsasser and Reynolds numbers simpler., This choice has several motivations: $S$ rather than $\Omega$ is the relevant time-scale of linear growth rates; this definition is more consistent with our previous choice of units; and it makes the connection between the Elsasser and Reynolds numbers simpler.
+ In the dissipation-free limit. the dispersion relation can be solved exactly.," In the dissipation-free limit, the dispersion relation can be solved exactly."
+ The magnetic tension stabilizes the instability fork>Ki80Ο)=(2/9) 7., The magnetic tension stabilizes the instability for $k^* > k_c^*\equiv (2\Omega^*)^{1/2}=(2/q)^{1/2}$ .
+ In a box of finite vertical extent L-. the vertical modes wavelengths are multiple of Kyi.= 2n/L..," In a box of finite vertical extent $L_z$ , the vertical modes wavelengths are multiple of $k_{min}=2\pi/L_z$ ."
+ Axisymmetric MRI modes are therefore stabilized when ko< Kus. Which translates intoB>B.=27$/Qz30.," Axisymmetric MRI modes are therefore stabilized when $k_c < k_{min}$ which translates into $\beta \ge \beta_c\equiv 2\pi^2 S/\Omega\simeq
+30$."
+" The maximum growth rate in the dissipation-free regime. $/2. 1s achieved for a wavenumber ky=k,./2."," The maximum growth rate in the dissipation-free regime, $\sigma_M=S/2$ , is achieved for a wavenumber $k_M\simeq k_c/2$."
+ Resistivity and viscosity modify the marginal stability limit in the (Re. Riv. B) space.," Resistivity and viscosity modify the marginal stability limit in the $Re$, $Rm$, $\beta$ ) space."
+ umerical investigation indicates that only one of the four roots is unstable for g<2. so that the marginal stability limits obtains when the o-independent term of the dispersion relation is equal to zero.," Numerical investigation indicates that only one of the four roots is unstable for $q < 2$, so that the marginal stability limits obtains when the $\sigma$ -independent term of the dispersion relation is equal to zero."
+ Also. as shown by ?.. for B2100 and Re3400 (two conditions that are satisfied in the present investigation). the marginal. stability limit is nearly independent of the Reynolds number.," Also, as shown by \cite{LL07}, for $\beta \ge 100$ and $Re\gtrsim 400$ (two conditions that are satisfied in the present investigation), the marginal stability limit is nearly independent of the Reynolds number."
+ This leads to the following expression of the marginal stability limit where the second equality applies for keplerian flows., This leads to the following expression of the marginal stability limit where the second equality applies for keplerian flows.
+ The last approximation holds only when f is large enough. and was given in ?:: the first expression is new and sensibly more precise.," The last approximation holds only when $\beta$ is large enough, and was given in \cite{LL07}; the first expression is new and sensibly more precise."
+" The maximum growth rate achieved by the instability is modified when viscosity and resistivity become important. Le. when either E, or E, €1."," The maximum growth rate achieved by the instability is modified when viscosity and resistivity become important, i.e. when either $E_\nu$ or $E_\eta$ $\lesssim 1$."
+ ? give an asymptotic form of the growth rate in the dissipation regime for small or of order unity Pm (see their Fig., \cite{PC08} give an asymptotic form of the growth rate in the dissipation regime for small or of order unity $Pm$ (see their Fig.
+ 5 and their Eq., 5 and their Eq.
+" 91): it reads ση=209E,/3 where c=S/2 1s the dissipation-free growth rate.", 91); it reads $\sigma_M=2\sigma_0 E_\eta/3$ where $\sigma_0=S/2$ is the dissipation-free growth rate.
+" The range of E, and Pri examined in this work spans the transition from a dissipative to à non dissipative regime for the linear instability.", The range of $E_\eta$ and $Pm$ examined in this work spans the transition from a dissipative to a non dissipative regime for the linear instability.
+" The corresponding variation of maximum growth rate with both E, and E, (through the Prandtl] number) is shown on Fig. 2.."," The corresponding variation of maximum growth rate with both $E_\eta$ and $E_\nu$ (through the Prandtl number) is shown on Fig. \ref{growthrate},"
+ along with the asymptotic expressions just recalled., along with the asymptotic expressions just recalled.
+" In order to explore a significant range of parameters. most of our runs are performed at a standard resolution (N,. N4. N-) = (256. 128. 64) in real space."," In order to explore a significant range of parameters, most of our runs are performed at a standard resolution $N_r$, $N_\phi$, $N_z$ ) = (256, 128, 64) in real space."
+ All simulations have an aspect ratio R:é:z24:4:]1:as argued earlier. this aspect ratio allows us to capture the fastest growing channel and parasitic modes.," All simulations have an aspect ratio $R:\phi:z=4:4:1$; as argued earlier, this aspect ratio allows us to capture the fastest growing channel and parasitic modes."
+ Three simulations have been performed at a resolution twice as large (512. 256. 128) i order to reach a higher Reynolds number (Re=20000). and lower Prandtl numbers (down to Pm=1/16).," Three simulations have been performed at a resolution twice as large (512, 256, 128) in order to reach a higher Reynolds number $Re=20000)$, and lower Prandtl numbers (down to $Pm=1/16$ )."
+ Note that spectral codes are intrinsically more resolved than finite difference codes with the same number of points: as a consequence. one needs to adopt resolutions larger by a factor =2 in all directions in a finite difference code such as ZEUS. Athena or RAMSES to obtain the same results. a point to bear in mind when comparing our conclusions with those published in the literature with finite difference codes e.g.?.foranexplicit comparison).," Note that spectral codes are intrinsically more resolved than finite difference codes with the same number of points; as a consequence, one needs to adopt resolutions larger by a factor $\simeq 2$ in all directions in a finite difference code such as ZEUS, Athena or RAMSES to obtain the same results, a point to bear in mind when comparing our conclusions with those published in the literature with finite difference codes \citep[see e.g.][for
+an explicit comparison]{FPLH07}."
+ The simulations performed in this work are in the large Reynolds number limit (400<Re 20000)., The simulations performed in this work are in the large Reynolds number limit $400 < Re < 20000$ ).
+ The relevant marginal stability limit depends only on Aii and Is precisely given by Eq. (23))., The relevant marginal stability limit depends only on $Rm$ and is precisely given by Eq. \ref{etalim}) ).
+ When the magnetic Reynolds number is too close to this limit. nonaxisymmetric perturbations damp out. and the flow becomes nearly identical to an axisymmetric one.," When the magnetic Reynolds number is too close to this limit, nonaxisymmetric perturbations damp out, and the flow becomes nearly identical to an axisymmetric one."
+ We have checked that none of the simulations discussedbelow is affected by this issue., We have checked that none of the simulations discussedbelow is affected by this issue.
+ With our adopted standard resolution. we can confidently resolve the dissipative scales up to Reynolds numbers ~10°.," With our adopted standard resolution, we can confidently resolve the dissipative scales up to Reynolds numbers $\simeq 10^4$."
+ Consequently. we have explored Reynolds numbers in the range 400 — 6400 and magnetic Reynolds numbers in the range 300 — 12800.," Consequently, we have explored Reynolds numbers in the range $400$ – $6400$ and magnetic Reynolds numbers in the range $400$ – $12800$."
+ The magnetic Prandtl number of the simulations is either 1/3. 1: or 4.," The magnetic Prandtl number of the simulations is either $1/4$, $1$ or $4$."
+ A smaller value (Par= 1/16) has also been reached with one of our higher resolution runs with Re= 20000., A smaller value $Pm=1/16$ ) has also been reached with one of our higher resolution runs with $Re=20000$ .
+ We have performed simulations for three different vertical field strengths: 8=107. 10°. and 107.," We have performed simulations for three different vertical field strengths: $\beta=10^2$, $10^3$, and $10^4$."
+ This range is constrained by two considerations., This range is constrained by two considerations.
+ On the low end (large field). one must stay away from the marginal stability limit due to the magnetic tension.," On the low end (large field), one must stay away from the marginal stability limit due to the magnetic tension."
+ Indeed. when the incompressible MRI works too close to this stability threshold. the flow is strongly biased by the presence of channel modes (an effect that incompressible simulations are more prone to capture. as exemplified in 2)).," Indeed, when the incompressible MRI works too close to this stability threshold, the flow is strongly biased by the presence of channel modes (an effect that incompressible simulations are more prone to capture, as exemplified in \citealt{LL07}) )."
+ On the large end (weak field). one is limited by thefact that at some point. the flow Is going to be dominated by thezero flux transport process.whose dependence on the viscous and resistive dissipation might be sensibly different.," On the large end (weak field), one is limited by thefact that at some point, the flow is going to be dominated by thezero flux transport process,whose dependence on the viscous and resistive dissipation might be sensibly different."
+ If one simply assumes that one shifts from one regime to the other when the efficiency of the transport processes are equal. one finds that the transition between the mean field and zero net field regimes should be in the rangeB~10? — 10: this limit applies," If one simply assumes that one shifts from one regime to the other when the efficiency of the transport processes are equal, one finds that the transition between the mean field and zero net field regimes should be in the range $\beta\sim 10^{5}$ –$10^{4}$ ; this limit applies"
+The above does not explain why our sources should have deeper silicate absorption than local radio-loud populations.,The above does not explain why our sources should have deeper silicate absorption than local radio-loud populations.
+ The answer hinges on (he unknown origin of the mid-In obscuration — s.a., The answer hinges on the unknown origin of the mid-IR obscuration – s.a.
+ cdustiy torus or host galaxy., dusty torus or host galaxy.
+ The deep silicate features of our sources imply oplically Chick obseuring media with high filling factors (i.e. likely pulfed-up geometry. but not 4x given the radio jets) as thev have to effectively hide the hot inner regions of the AGN (Pier&Ixrolik1992:Levensonetal.2007).," The deep silicate features of our sources imply optically thick obscuring media with high filling factors (i.e. likely puffed-up geometry, but not $\pi$ given the radio jets) as they have to effectively `hide' the hot inner regions of the AGN \citep{pk92,levenson07}."
+. AGN tori are believed to be ebumpsy. which precludes high filling factors (Elitzur&Shlosman2006).. consistent with their observed weak silicate features (Shietal.2006).," AGN tori are believed to be clumpy, which precludes high filling factors \citep{elitzur06}, consistent with their observed weak silicate features \citep{shi06}."
+. Iligher τος can be due to a cold dust screen 2006).. such as expected from a dusty host galaxy.," Higher $\tau_{9.7}$ can be due to a cold dust screen \citep{ogle06}, such as expected from a dusty host galaxy."
+ Hieh-z radio galaxies/radio-loud quasars have been detected in the sub-mm indicating cold dust (Willott 2003)..," $z$ radio galaxies/radio-loud quasars have been detected in the sub-mm indicating cold dust \citep{willott02,dusty_radio}."
+ We already. know that the number density of deeply obscured sources is about. 10 limes higher at z — 22 than at z 00 (Sajinaetal.2007)., We already know that the number density of deeply obscured sources is about 10 times higher at $z$ $\sim$ 2 than at $z$ $\sim$ 0 \citep{sajina07}.
+. Lastly. MIPSI6122 is consistent with a jet axis close to the line-o[-sight (<< 445°).," Lastly, MIPS16122 is consistent with a jet axis close to the line-of-sight $<<$ $^{\circ}$ )."
+ The dusty torus of an AGN is believed to be coupled with the accretion disk (e.g.Elitzur&Shlosman2006).. implying a nearly torus lor MIPSI6122. which is unlikely given its deep silicate absorption feature.," The dusty torus of an AGN is believed to be coupled with the accretion disk \citep[e.g.][]{elitzur06}, implying a nearly face-on torus for MIPS16122, which is unlikely given its deep silicate absorption feature."
+ On the other hand. a mis-mateh between the orientation of the host galaxy anc AGN is frequentle observed (Schmittοἱal.2002).," On the other hand, a mis-match between the orientation of the host galaxy and AGN is frequently observed \citep{schmitt02}."
+. Thus. although the origin of the mic-IR. obscuration remains an open question. host galaxy obscuration is a good candidate.," Thus, although the origin of the mid-IR obscuration remains an open question, host galaxy obscuration is a good candidate."
+ Many thanks to Patrick Ogle for a critical reading of the manuscript prior to submission., Many thanks to Patrick Ogle for a critical reading of the manuscript prior to submission.
+ We are also grateful to N. Sevinour. Ix. Dasvra. D. Partridge. D. Lutz. and D. Scott for useful discussions.," We are also grateful to N. Seymour, K. Dasyra, B. Partridge, D. Lutz, and D. Scott for useful discussions."
+ We are grateful to (he anonvinous referee for (heir helpfull comments., We are grateful to the anonymous referee for their helpfull comments.
+" This work is based on observations made with theSpitzer Space Telescope. which is operated bv the Jet Propulsion Laboratory, California Institute of Technology under contract with NASA."," This work is based on observations made with the Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under contract with NASA."
+ Support for this work was provided by NASA through an awarel issued by JPL/Caltech., Support for this work was provided by NASA through an award issued by JPL/Caltech.
+ This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated bv the Jet Propulsion Laboratory. California Institute ol Technology. under contract with the National Aeronautics ancl Space Administration.," This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."
+Although bright irregulu galaxies have rotating eas disks. it is unccar whether the faintest dwart ireeular galaxies are rotaionallv supported or not.,"	Although bright irregular galaxies have rotating gas disks, it is unclear whether the faintest dwarf irregular galaxies are rotationally supported or not."
+ Ctte et al. (, Côtté et al. (
+2000). based ou a study of the liuemiatics of elelt dxvarf regular ealaxies (with maenitides varving frou Alp=16.7 to Mp=113 ) sugecst tlat jormial rotalon is seen only iu dwarts brighter than Bvτι.,"2000), based on a study of the kinematics of eight dwarf irregular galaxies (with magnitudes varying from $_{\rm B} = -16.7$ to $_{\rm B} = -11.3$ ) suggest that normal rotation is seen only in dwarfs brighter than $_B \sim -14$."
+ This Is consistent with the earlier findings of Loctal.(1993).. who. based on au iuterferoimnetrie studs MEC faiit ¢wait eaaxies (with Mp~9 to Mp~ 15) oud. hat ouly two of their sample Q [nine galaxies show Orelered velociv fields.," This is consistent with the earlier findings of \cite{lo93}, who, based on an interferometric study of faint dwarf galaxies (with $_{B\rm} 
+\sim -9$ to $_{B\rm} \sim -15$ ) found that only two of their sample of nine galaxies showed ordered velocity fields."
+ Towever. this conclusion las cel. questioned |w Skillman(1996) who pointed out tliat the interferomeiic observations O: Loetal.(1993) 1.icked sensitivity to low extended ITI distribution and coud thus have heen luscnsitive to the argo scale velocity field.," However, this conclusion has been questioned by \cite{skillman96} who pointed out that the interferometric observations of \cite{lo93} lacked sensitivity to low extended HI distribution and could thus have been insensitive to the large scale velocity field."
+ Further. Beeunetal.(2003) showed that the dwarf ireeular galaxy Cauclopardalis D. despite being extremely faint (Mp~ 10.9) nonetheless jas a regular velocity field. cosistent with tlat expeced roni a rotating dis£o," Further, \cite{begum03} showed that the dwarf irregular galaxy Camelopardalis B, despite being extremely faint $M_B \sim -10.9$ ) nonetheless has a regular velocity field, consistent with that expected from a rotating disk."
+ So. sone faint σαaxies at least. rave rotating III disks.," So, some faint galaxies at least, have rotating HI disks."
+ Wha about the vest?, What about the rest?
+ Tf eas in aut dwarf galaxies is not supported by rotation. what oxovides the energy hat keeps it frou collapse?," If gas in faint dwarf galaxies is not supported by rotation, what provides the energy that keeps it from collapse?"
+ For very dut dwarf ireeular ealaxies. the dineling energy of the eas Is HO much larecr than he energy out of a few supernovac.," For very faint dwarf irregular galaxies, the binding energy of the gas is not much larger than the energy output of a few supernovae."
+ Star fornation iu such σαlaxies could heuce rave a profound effect ou t1ο kinematics of the ISAL, Star formation in such galaxies could hence have a profound effect on the kinematics of the ISM.
+ ludeed. the faiutest dwarf galaxies are expected to k)8C a substantial part of heir eas due to the enerev deposied in the ISAL by superunovae frou the first burst of star formation (c.g. De|] Silk 198).," Indeed, the faintest dwarf galaxies are expected to lose a substantial part of their gas due to the energy deposited in the ISM by supernovae from the first burst of star formation (e.g. Dekel Silk 1986)."
+ 8ome observational supxort for such models is provider bvthe large expanding II supershells seen in the ISAL of some dwarf regular ealaxies with active star formatiol1 (e.g. Ott et al., Some observational support for such models is provided by the large expanding HI supershells seen in the ISM of some dwarf irregular galaxies with active star formation (e.g. Ott et al.
+ 2001)., 2001).
+ Iu this paper we disciiss the issue of the làueiaties of fait dwarf galaxies. aud its possible counectious with energy input from stellar processes. in the specific context of the faint (Mp~12.1) dwarf nreeular galaxy CR.," In this paper we discuss the issue of the kinematics of faint dwarf galaxies, and its possible connections with energy input from stellar processes, in the specific context of the faint $M_B \sim -12.1$ ) dwarf irregular galaxy GR8."
+ CRs was first discovered by Reaves(1956) in the course of a survey for dwarf galaxies in the direction of the Vireo Cluster., 	GR8 was first discovered by \cite{reaves56} in the course of a survey for dwarf galaxies in the direction of the Virgo Cluster.
+ It has also been cataloged as DDO 155 by vandenBer:el(1959)..., It has also been cataloged as DDO 155 by \cite{vdberg59}. .
+ The original distance cstimates for GRs WCTC in the rage 1.0 - L.t Alpe (Hodee(1967).DeVaucoileurs&Moss(1983).Iloessel&Danielson (1983))). whicji would malse GRs à probable member of the local eroup.," The original distance estimates for GR8 were in the range 1.0 - 1.4 Mpc \cite{hodge67,devac83,
+hoessel83}) ), which would make GR8 a probable member of the local group."
+ Ilowever. recent estimates eive somewhat hueer clistances.," However, recent estimates give somewhat larger distances."
+ Tolstoyetal.(1995) estimated a distauce of 2.2 Ape (based on observations of the oulv detected Cepheid variable)., \cite{tolstoy95} estimated a distance of 2.2 Mpc (based on observations of the only detected Cepheid variable).
+ Tus estimate is 1n ονοςlent agreementC» with that of Doluw-Palueretal.(1998) which Is based ou the brightuess of the tip of the red giaut branch., This estimate is in excellent agreement with that of \cite{dohm98} which is based on the brightness of the tip of the red giant branch.
+ From the ocation for the local group barvecuter eiven bv Courteau&van¢chBerel(1999) One call compute that this distance aces CRs well outside the local group zero velocity sirface., From the location for the local group barycenter given by \cite{courteau99} one can compute that this distance places GR8 well outside the local group zero velocity surface.
+ Cosistent with this. vali(euBereh(2000) does not classify GRA as a member of the local eroup.," Consistent with this, \cite{vdberg00} does not classify GR8 as a member of the local group."
+ GR& has a patchy appearance niofical images. wihi the emission απο dominated by brigit blue knots.," 	GR8 has a patchy appearance in optical images, with the emission being dominated by bright blue knots."
+ Πα inaeiue (Ilodee 1967) shows that these knots are sites of active star formation., $\alpha$ imaging (Hodge 1967) shows that these knots are sites of active star formation.
+ Towever. in adition to the brigit hhe shots. COURS also possesses faint extended: emissiDn (IIxlee(1967).DeViuicouleurs&Moss( 1983))). nmicative of earlier episodes of star formation.," However, in addition to the bright blue knots, GR8 also possesses faint extended emission \cite{hodge67, devac83}) ), indicative of earlier episodes of star formation."
+ Tudeed. CAL diagranis (based ou IST inaeglueg. Doluu-Pahueretal. (1998))). show that althoeh the bright star fornüug krots in CRS have stars which are VOInecr than ~10 Avr. the ealaxy also contaiis stars which are older than a fow Cr.," Indeed, CM diagrams (based on HST imaging, \cite{dohm98}) ), show that although the bright star forming knots in GR8 have stars which are younger than $\sim 10$ Myr, the galaxy also contains stars which are older than a few Gyr."
+ Despite this long history of star formation. the metalicity. of the star ornmiug knots in GR& is extremely QW. solar (Skilnua et al.," Despite this long history of star formation, the metallicity of the star forming knots in GR8 is extremely low, $\sim 3\%$ solar (Skillman et al."
+ 1988b)., 1988b).
+ This makes it ouk| of the lowest iueallicitv ealaxies known (Ikuuth Osxlin 2000)., This makes it one of the lowest metallicity galaxies known (Kunth Össtlin 2000).
+ In keepi1ο with this low metallicity. despite the| fact that the star ornmiug regions are expected to be ass¢dated with moleclay Oogas. CO has not been detected in tιο galaxy (Verter Tocdec1995).," In keeping with this low metallicity, despite the fact that the star forming regions are expected to be associated with molecular gas, CO has not been detected in the galaxy (Verter Hodge1995)."
+used in Fieures 5b. aud 10..,used in Figures \ref{fg5} and \ref{fg10}.
+. sources in the sample of Titovetal.(2011). have proper motions with values iu the range within the 3o of the mean in cach component., sources in the sample of \cite{ti11} have proper motions with values in the range within the $3\sigma$ of the mean in each component.
+ Figure Ll indicates that the mean of proper motions with high accuracy for quasars is close to zero., Figure \ref{fg14} indicates that the mean of proper motions with high accuracy for quasars is close to zero.
+ Based o- a sanuple of 62 radio sources. Moóretal.(2001). poiute«a out that the proper mofious of the radio sources are caused bv the chauges in their brightuess distribution structure.," Based on a sample of 62 radio sources, \cite{mo11} pointed out that the proper motions of the radio sources are caused by the changes in their brightness distribution structure."
+" The systematic motion of quasars can be detected at present is the secular aberration dift with a value of ο,£41.5 µας l(Titov etal.2011).", The systematic motion of quasars can be detected at present is the secular aberration drift with a value of $6.4\pm 1.5$ $\mu$ as $^{-1}$ \citep{ti11}.
+. In eoncral. the mean of proper motions of quasars can be considered as zero with preseut observation accuracy.," In general, the mean of proper motions of quasars can be considered as zero with present observation accuracy."
+ The systematic motion of 6.1 pas 1 of quasars cannot be detected by the present ground-based observation iu the optical baud., The systematic motion of $6.4$ $\mu$ as $^{-1}$ of quasars cannot be detected by the present ground-based observation in the optical band.
+ Based ou the assunption that the mean of proper motions for a sample of quasars should be zero. iu deviation from zero of the mean of the proper motions of the sample can be considered as the systematic error of the proper motions. aud the dispersion of the proper motions from the mean cau be considered as the random error of the proper motions of the sample.," Based on the assumption that the mean of proper motions for a sample of quasars should be zero, any deviation from zero of the mean of the proper motions of the sample can be considered as the systematic error of the proper motions, and the dispersion of the proper motions from the mean can be considered as the random error of the proper motions of the sample."
+ A sample of 117053 quasars ideutified iu the PPMXL catalog were used to derive the svstematic aud raudoni errors in the proper motions of the PPAINE catalog., A sample of 117053 quasars identified in the PPMXL catalog were used to derive the systematic and random errors in the proper motions of the PPMXL catalog.
+" The systematic error of the proper motious for the whole quasar sample ix derived as 2.0 and 2.1 mas | du pr,cosd and fis. respectively."," The systematic error of the proper motions for the whole quasar sample is derived as $-2.0$ and $-2.1$ mas $^{-1}$ in $\mu_{\alpha}\cos\delta$ and $\mu_{\delta}$ , respectively."
+ The random error of the proper motious for the whole quasar sample is 5.3 and L8 mas vr tin Hacos9 and fr. respectively.," The random error of the proper motions for the whole quasar sample is $5.3$ and $4.8$ mas $^{-1}$ in $\mu_{\alpha}\cos\delta$ and $\mu_{\delta}$, respectively."
+ No obvious magnitude depeudeuce of the systematic errors of the proper motions iu both componcuts are ound., No obvious magnitude dependence of the systematic errors of the proper motions in both components are found.
+ The nus of the systematic error of the proper notions in cach coniponeut for quasars with different uagnitudes is ~O.5 anas d., The rms of the systematic error of the proper motions in each component for quasars with different magnitudes is $\sim 0.5$ mas $^{-1}$.
+ There ds obvious naguitude depeudenuce of the random errors of the proper notions iu both components., There is obvious magnitude dependence of the random errors of the proper motions in both components.
+ The random eror of the xoper motions in cach component can change frou ~LU to ~RU mas + alone with the iuerease of the naguitude of quasars., The random error of the proper motions in each component can change from $\sim4.0$ to $\sim8.0$ mas $^{-1}$ along with the increase of the magnitude of quasars.
+ No obvious color depeucdence of he systematic and random errors of the proper motions in both components are found., No obvious color dependence of the systematic and random errors of the proper motions in both components are found.
+ The xius of the svsteatic errors of the proper motions iu each conrponeutfor quasars with different colors is ~0.5 mas vrft., The rms of the systematic errors of the proper motions in each componentfor quasars with different colors is $\sim 0.5$ mas $^{-1}$.
+ The rni of the random error of the proper motions iu cach component for quasars with different colors is 0.5 mas band ~0.5 mas | for photographie BoR and SDSS gysor. respectively.," The rms of the random error of the proper motions in each component for quasars with different colors is $\sim 0.8$ mas $^{-1}$ and $\sim 0.5$ mas $^{-1}$ for photographic $B-R$ and SDSS $g-r$, respectively."
+ The systematic errors of the proper motions iu both colponents Fare correlated with à of quasars: 8x sina., The systematic errors of the proper motions in both components $\overline{\mu}$ are correlated with $\alpha$ of quasars: $\overline{\mu}\propto \sin\alpha$ .
+ No obvious à depeudence of the random crrors of the proper motions iu both compoucuts are found., No obvious $\alpha$ dependence of the random errors of the proper motions in both components are found.
+ The rni of the random error of the proper motions iu cach compoucut for quasars with different à is ~I0 mias |., The rms of the random error of the proper motions in each component for quasars with different $\alpha$ is $\sim 1.0$ mas $^{-1}$.
+ For quasars with à« 0. their raudom errors of the proper motions in both conponeuts are bigeer thin those for quasars with à> 0: and their svstematic errors of the proper motious in both coniponents are very different from those for quasars with à>0.," For quasars with $\delta<0$ , their random errors of the proper motions in both components are bigger than those for quasars with $\delta >0$ ; and their systematic errors of the proper motions in both components are very different from those for quasars with $\delta>0$."
+ Especially. quasars iu the neighborhoodof 6=30° have the ιαπ of the systematic error of the proper motious iu both colmpoucuts ~5 mias L," Especially, quasars in the neighborhoodof $\delta=-30\degr$ have the maximum of the systematic error of the proper motions in both components $\sim 5$ mas $^{-1}$."
+ For quasars with à>0. there is no obvious 6 dependence of the systematic aud random errors of νιcos.," For quasars with $\delta>0$, there is no obvious $\delta$ dependence of the systematic and random errors of $\mu_{\alpha}\cos\delta$."
+" The rus of the systematic and random errors of ji,co80. for quasars with à>(0 is ~0.5 aud ~O.7Fimasvr ο, respectively."," The rms of the systematic and random errors of $\mu_{\alpha}\cos\delta$ for quasars with $\delta>0$ is $\sim 0.5$ and $\sim 0.7$ mas $^{-1}$, respectively."
+ But for quasars with à> the systematic errors of p5 increase with the à aud can be best described as fixX0.076.," But for quasars with $\delta>0$, the systematic errors of $\mu_{\delta}$ increase with the $\delta$ and can be best described as $\overline{\mu_{\delta}}\propto 0.07 \delta$."
+ There is no obyious à epeudence of the random errors of jj; for quasars with à>0., There is no obvious $\delta$ dependence of the random errors of $\mu_{\delta}$ for quasars with $\delta>0$.
+ The ruis of the random errors of p for quasars with àc0ds ~02 mas vro/., The rms of the random errors of $\mu_{\delta}$ for quasars with $\delta>0$ is $\sim 0.3$ mas $^{-1}$.
+ The 2MASS observation data ire important in the erivation of the proper motions in the PPMXLE catalog., The 2MASS observation data are important in the derivation of the proper motions in the PPMXL catalog.
+" The systematic error of jf,cos0 for quasars with 241ÀASS ata is about half of that for quasars without 2\TASS ata.", The systematic error of $\mu_{\alpha}\cos\delta$ for quasars with 2MASS data is about half of that for quasars without 2MASS data.
+ For comparison. the svstematic and raundoni errors of the USNO-SDSS proper motions for a quasar sample identified in the PPAINL aud SDSS DR7 catalogs are erived.," For comparison, the systematic and random errors of the USNO-SDSS proper motions for a quasar sample identified in the PPMXL and SDSS DR7 catalogs are derived."
+" The systematic error of the USNO-SDSS proper motions is 0.9 and 0.5 mas | in yr,cosd and pry. respectively,"," The systematic error of the USNO-SDSS proper motions is $0.9$ and $0.5$ mas $^{-1}$ in $\mu_{\alpha}\cos\delta$ and $\mu_{\delta}$, respectively."
+" The random error of the USNO-SDSS proper motions is 3.3 and 3.3 amas tin 45,cosó aud fis. respectively."," The random error of the USNO-SDSS proper motions is $3.3$ and $3.3$ mas $^{-1}$ in $\mu_{\alpha}\cos\delta$ and $\mu_{\delta}$, respectively."
+ No obvious magnitude dependence of the &vstematic errors of the USNO-SDSS proper motions in both conmonuenuts are found., No obvious magnitude dependence of the systematic errors of the USNO-SDSS proper motions in both components are found.
+" The ruis of the systematic errors of the USNO-SDSS proper motious iu each component for quasars with different maenitudes is ~0. [anas vr+, ", The rms of the systematic errors of the USNO-SDSS proper motions in each component for quasars with different magnitudes is $\sim 0.4$ mas $^{-1}$ .
+The random error of the USNO-SDSS proper motions iu cach coniponent increases frou ~2.5 mas | to ~6.0 mas | along with the imcerease of the SDSS r inaenitude of quasars., The random error of the USNO-SDSS proper motions in each component increases from $\sim2.5$ mas $^{-1}$ to $\sim6.0$ mas $^{-1}$ along with the increase of the SDSS $r$ magnitude of quasars.
+ No obvious color depeudenuce of the systematic errors of the USNO-SDSS. proper motions iun both components is found., No obvious color dependence of the systematic errors of the USNO-SDSS proper motions in both components is found.
+ The rus of the systematic errors of tle USNO-SDSS proper motions iu cach component for quasars with different gor color is ~O.3 mas P., The rms of the systematic errors of the USNO-SDSS proper motions in each component for quasars with different $g-r$ color is $\sim 0.3$ mas $^{-1}$.
+ But. the random errors of the USNO-SDSS proper motious in each component cau chanec from ~2.5 mas to ~6.0 mas | alone with the SDSS g rcolor of [|between 0.5 and 1.5.," But, the random errors of the USNO-SDSS proper motions in each component can change from $\sim2.5$ mas $^{-1}$ to $\sim6.0$ mas $^{-1}$ along with the SDSS $g-r$ color of between $-0.5$ and $1.5$."
+ There is no a dependence of the systematic ane random errors of the USNO-SDSS proper motions in both two componcuts., There is no $\alpha$ dependence of the systematic and random errors of the USNO-SDSS proper motions in both two components.
+ The rms of the systematic anc random errors of the USNO-SDSS proper motions in cach componcut for quasars with differeut à is O.1 are ~O.Linasyr ft. respectively.," The rms of the systematic and random errors of the USNO-SDSS proper motions in each component for quasars with different $\alpha$ is $\sim 0.4$ and $\sim 0.4$ mas $^{-1}$, respectively."
+ There is uo 6 dependence of the random errors of the USNO-SDSS proper motions iu both two components., There is no $\delta$ dependence of the random errors of the USNO-SDSS proper motions in both two components.
+" The rus of the random errors of the USNO-SDSS proper motions in cach component. for quasars with different 8 is ~0.1 aqnas vr|,", The rms of the random errors of the USNO-SDSS proper motions in each component for quasars with different $\delta$ is $\sim 0.4$ mas $^{-1}$.
+" No obvious à dependence of the svstematic errors of the USNO- proper motions in Jn,cos0 is found.", No obvious $\delta$ dependence of the systematic errors of the USNO-SDSS proper motions in $\mu_{\alpha}\cos\delta$ is found.
+ The rus of the systematic errors of tle USNO-SDSS proper motions iu Hacosfor quasars with different à is ~0.2 mas vr il , The rms of the systematic errors of the USNO-SDSS proper motions in $\mu_{\alpha}\cos\delta$for quasars with different $\delta$ is $\sim 0.2$ mas $^{-1}$ .
+"The svstematic errors of the USNO-SDSS proper motions in ply increasewith the àand can bebest described as FOX 0.036, ", The systematic errors of the USNO-SDSS proper motions in $\mu_{\delta}$ increasewith the $\delta$and can bebest described as $\overline{\mu_{\delta}}\propto 0.03\delta$ .
+Tn general. comparing with the proper notions in he PPAINL catalog. using the SDSS photometric and astrometric data can siguificautly reduce the systematic and random crrors of the derived proper notions.," In general, comparing with the proper motions in the PPMXL catalog, using the SDSS photometric and astrometric data can significantly reduce the systematic and random errors of the derived proper motions."
+ Thus. combining the photographic dataand he present observation data obtained from the erouud- telescopes with CCD. the proper motions with systematic error less than Lamas vr? aud raucous error ess than 5 mas  for objects with r<20 mag can be derived.," Thus, combining the photographic dataand the present observation data obtained from the ground-based telescopes with CCD, the proper motions with systematic error less than 1 mas $^{-1}$ and random error less than 5 mas $^{-1}$ for objects with $r < 20$ mag can be derived."
+halfinass relaxation time may therefore imply rapid mass segregation for the massive stars or support agereeation (?) as à mechauisui for high mass star formation.,half-mass relaxation time may therefore imply rapid mass segregation for the massive stars or support aggregation \citep{bonnell02} as a mechanism for high mass star formation.
+" We present uew. high resolution IKeck/NIRSPEC echelle spectrometry and new IIST/ACS imagine of the superstar cluster MIS2-F. The main results of this study are: We would like to thauk the staff of the Necks Observatory. aud observing assistant Cary Puniwai specifically,"," We present new, high resolution Keck/NIRSPEC echelle spectrometry and new HST/ACS imaging of the super--star cluster M82-F. The main results of this study are: We would like to thank the staff of the Keck Observatory, and observing assistant Gary Puniwai specifically."
+ We also thank L. J. Snuüth for providing us with optical spectroscopy of MS2-F.. NAL thanks I. R. Wine for helpful discussions ou elliptical clusters. A. Cottou-Clay. for mathematical assistance. aud J. Terrell for generous computing support.," We also thank L. J. Smith for providing us with optical spectroscopy of M82-F. NM thanks I. R. King for helpful discussions on elliptical clusters, A. Cotton-Clay for mathematical assistance, and J. Terrell for generous computing support."
+ The authors wish to recognize and acknowledge the very significant cultural role aud reverence that the sununüt of Mauna EKea has always had within the indigenous UWawaiian community., The authors wish to recognize and acknowledge the very significant cultural role and reverence that the summit of Mauna Kea has always had within the indigenous Hawaiian community.
+ We are most fortunate to have the opportunity to conduct observations from this mountain., We are most fortunate to have the opportunity to conduct observations from this mountain.
+ This work has been supported by NSF Craut AST0205999., This work has been supported by NSF Grant AST–0205999.
+1976: Frail et al. 1991)).,"\nocite{am76}; Frail et al. \nocite{fchw91}) ),"
+ which validates our experimental technique., which validates our experimental technique.
+ We thus consider. the presented: absorption spectrum of PSR. 6545 to be free of any appreciable contamination from features in the emission spectrum., We thus consider the presented absorption spectrum of PSR $-$ 6545 to be free of any appreciable contamination from features in the emission spectrum.
+ The distance ascribed. to this pulsar [from its dispersion measure is 3.2 kpc., The distance ascribed to this pulsar from its dispersion measure is 3.2 kpc.
+ The limit obtained here serves as an independent distance measure anc allows a number of consequences to be explored., The limit obtained here serves as an independent distance measure and allows a number of consequences to be explored.
+ As the DM of the pulsar is known. we can use the neutral hyerogen distance to give an upper limit on the electron clensity in this direction.," As the DM of the pulsar is known, we can use the neutral hydrogen distance to give an upper limit on the electron density in this direction."
+ “Phe DAL of PSR 6545 is approximately 116.2 pe em7., The DM of PSR $-$ 6545 is approximately 116.2 pc $^{-3}$.
+ A distance limit of greater than 3.7 kpe provides an upper limit to the mean electron density along this line of sight of 0.03 cm*., A distance limit of greater than 3.7 kpc provides an upper limit to the mean electron density along this line of sight of 0.03 $^{-3}$.
+ The secular decrease of binary orbital period. £5. has been successfully observed. in PSI. D1913]16 ClEavlor& and PSR D15341|12. (Stairsetal.1998).," The secular decrease of binary orbital period, $\dot{P_b}$, has been successfully observed in PSR B1913+16 \cite{tw89,dt91,tay94} and PSR B1534+12 \cite{sac+98}."
+.. While the measured value is. dominated by the emission. of eravitational radiation. unfortunately there are many possible contributions to an observed orbital period derivative (Damour&Vavlor1991).," While the measured value is dominated by the emission of gravitational radiation, unfortunately there are many possible contributions to an observed orbital period derivative \cite{dt91}."
+. They include: the component predicted by Git: the acceleration of the centre of mass of the binary system. with respect to the Sun clue to the gravitational field. of the Calaxy. and the apparent acceleration induced by the proper motion of the system.," They include; the component predicted by GR; the acceleration of the centre of mass of the binary system with respect to the Sun due to the gravitational field of the Galaxy, and the apparent acceleration induced by the proper motion of the system."
+ ‘Tauris and Sennels have predicted that the space velocity of the binary should be greater than 150 kms ! but ]xaspi ct al., Tauris and Sennels \nocite{ts00} have predicted that the space velocity of the binary should be greater than 150 km $^{-1}$ but Kaspi et al.
+ point out that the age and Galactic atitude of the svstem suggest. its. velocity perpendicular o plane must be —150(47/3.2kpe) Km I., \nocite{klm+00a} point out that the age and Galactic latitude of the system suggest its velocity perpendicular to plane must be $\sim 150(d/3.2kpc)$ km $^{-1}$.
+ Lt is possible o examine the magnitude of the kinematical contribution o anv measured orbital period. derivative., It is possible to examine the magnitude of the kinematical contribution to any measured orbital period derivative.
+" For transverse velocities in the range 100 - 200 km the level of the kinematical contribution to 2, is just a few percent. but without a measurement of the proper motion. we will be unable to confirm GR."," For transverse velocities in the range 100 - 200 km $^{-1}$ the level of the kinematical contribution to $\dot{P_b}$ is just a few percent, but without a measurement of the proper motion, we will be unable to confirm GR."
+ Phe prospects of making a quick Oper motion measurement are not clear., The prospects of making a quick proper motion measurement are not clear.
+ Young pulsars ike PSR. 6545 have intrinsic timing noise which make proper motion measurements difficult from timing on short. baselines., Young pulsars like PSR $-$ 6545 have intrinsic timing noise which make proper motion measurements difficult from timing on short baselines.
+ Interferometric observations may be more useful., Interferometric observations may be more useful.
+ PSR 6545 has a very narrow pulse and a reasonable [lux density at 20 cm wavelength., PSR $-$ 6545 has a very narrow pulse and a reasonable flux density at 20 cm wavelength.
+ Loa suitable reference could be found. an accurate proper motion should be obtainable via VLBI within a few vears.," If a suitable reference could be found, an accurate proper motion should be obtainable via VLBI within a few years."
+ An upper limit to the distance of PSR 6545 will be more cillicult. to. obtain., An upper limit to the distance of PSR $-$ 6545 will be more difficult to obtain.
+ The distance method emploved. by Stairs. et al., The distance method employed by Stairs et al.
+ for PSR DI1534| 12 will be less accurate in the case of PSR. 6545. because the kinematic terms are a much smaller fraction of the relativistic contribution to the observed. orbital period derivative.," \nocite{sac+98} for PSR $+$ 12 will be less accurate in the case of PSR $-$ 6545, because the kinematic terms are a much smaller fraction of the relativistic contribution to the observed orbital period derivative."
+ The measurement of the absorption spectrum of. this rnulsar has allowed a lower limit to be placed upon its distance of 3.7 kpc. the tangent point distance preclictec w the Galactic rotation curve.," The measurement of the absorption spectrum of this pulsar has allowed a lower limit to be placed upon its distance of 3.7 kpc, the tangent point distance predicted by the Galactic rotation curve."
+ The mean electron density in the direction. of this pulsar is therefore αἲ mos L03 7., The mean electron density in the direction of this pulsar is therefore at most 0.03 $^{-3}$.
+" We have also examined the Calactic anc kinematic contribution to the observed. D, and. found tha it should. only be at the few percent level for reasonable oulsar transverse velocities making PSR. 6545 an excellent. gravitational laboratory.", We have also examined the Galactic and kinematic contribution to the observed $\dot{P_b}$ and found that it should only be at the few percent level for reasonable pulsar transverse velocities making PSR $-$ 6545 an excellent gravitational laboratory.
+ Future measurements of he system's proper motion will help to further define the evel at which this svstem can be used to test Cit., Future measurements of the system's proper motion will help to further define the level at which this system can be used to test GR.
+The Orszag-Tang vortex (?) is a standard test problem that is used to validate many numerical MHD schemes.,The Orszag-Tang vortex \citep{1979JFM....90..129O} is a standard test problem that is used to validate many numerical MHD schemes.
+" The setup involves periodic domain of size [0,1]x with an adiabatic equation of state with y=5/3."," The setup involves periodic domain of size $[0,1]\times[0,1]$ with an adiabatic equation of state with $\gamma = 5/3$."
+ The initial density and pressure are set in all computational domain to 25/(367) and 5/(127) respectively., The initial density and pressure are set in all computational domain to $25/(36\pi)$ and $5/(12\pi)$ respectively.
+"The velocity v=(—sin(27y),+sin(272),0) and magnetic field B=(—Bosin(27:),Bosin(472),0.0), where Bo=1/νΑπ.","The velocity ${\bf v} = (-\sin(2\pi y), +\sin(2\pi x), 0)$ and magnetic field ${\bf B} = (-B_0\sin(2\pi y), B_0\sin(4\pi x), 0.0)$, where $B_0 =
+1/\sqrt{4\pi}$."
+" The second simulation involves the same initial conditions, with the exception that the problem is solved a boosted frame, with the initial velocity v=Vrest+Vboost, where Vboost=(10,10,10)."," The second simulation involves the same initial conditions, with the exception that the problem is solved a boosted frame, with the initial velocity ${\bf v} =
+{\bf v}_{\rm rest} + {\bf v}_{\rm boost}$, where ${\bf v}_{\rm boost} = (10, 10, 10)$."
+" The simulation is solved with 10° particles, where first 5-10* where randomly sampled within the square [0,1]x and the second 5-10* in the square of half size, [0.25,0.75]x0.75]."," The simulation is solved with $10^5$ particles, where first $5\cdot 10^4$ where randomly sampled within the square $[0,1]\times[0,1]$ and the second $5\cdot 10^4$ in the square of half size, $[0.25, 0.75]\times[0.25,0.75]$."
+ This permitted to achieve high resolution in the central region of the problem., This permitted to achieve high resolution in the central region of the problem.
+" The initial conditions were set, as soon as distribution was relaxed."," The initial conditions were set, as soon as distribution was relaxed."
+" In the we show density at t=0.5 for both rest-frame and boost-frame initial conditions, and in we show particle distribution."," In the we show density at $t = 0.5$ for both rest-frame and boost-frame initial conditions, and in we show particle distribution."
+" As expected, there are no discernible differences, and both simulations resolve discontinuities well."," As expected, there are no discernible differences, and both simulations resolve discontinuities well."
+" Furthermore, due to Lagrangian nature of the method, the particle number density correlates with the mass density."," Furthermore, due to Lagrangian nature of the method, the particle number density correlates with the mass density."
+" It is worth pointing out that in this particular simulation, the particle distribution appears to remain regular even in the vicinity of the shocks."," It is worth pointing out that in this particular simulation, the particle distribution appears to remain regular even in the vicinity of the shocks."
+ The show 1D pressure profile at y=0.3125 (top) and y=0.427 (bottom) for both boosted and rest-frame simulations.," The show 1D pressure profile at $y =
+0.3125$ (top) and $y = 0.427$ (bottom) for both boosted and rest-frame simulations."
+" The agreement with finite-difference scheme is excellent, ac can be compared to published results (e.g. ?,, ?,, ?))."," The agreement with finite-difference scheme is excellent, ac can be compared to published results (e.g. \citealp{2000JCoPh.161..605T}, , \citealp{2007MNRAS.379..915R}, \citealp{2008ApJS..178..137S}) )."
+" In contrast to shock-tube problems, no pressure blips are visible here."," In contrast to shock-tube problems, no pressure blips are visible here."
+" 'The rotor problem, introduced by ? to test propagation of strong torsional Alfvénn waves, is also considered as one of the standard candles to validate numericalMHD schemes."," The rotor problem, introduced by \cite{1999JCoPh.149..270B} to test propagation of strong torsional Alfvénn waves, is also considered as one of the standard candles to validate numericalMHD schemes."
+" Here, the computational domain is a unit square, [0,1]x 1]. The initial pressure and magnetic field are uniform with values p=1 and B=(6/ν4π, 0,0)."," Here, the computational domain is a unit square, $[0,1]\times[0,1]$ The initial pressure and magnetic field are uniform with values $p =
+1$ and ${\bf B} = (5/\sqrt{4\pi}, 0, 0)$ ."
+ Inside R<Ro=0.1 there is a dense, Inside $R < R_0 = 0.1$ there is a dense
+"p because da,zz1/(S2s) sometime before the jet becomes non-relativistic and à»~0 at this time. therebygiving στp (see eq. [12]]","$p$ because $\delta\alpha_1\approx 1/(8-2s)$ sometime before the jet becomes non-relativistic and $\alpha_2\approx 0$ at this time, therebygiving $\beta\approx p$ (see eq. \ref{fnua}] ]"
+ and Figures | 2): .) can exceed p. as can be seen in Figure 2. however the decrease in à» during the mildly relativistic phase prevents ο) from getting much larger than p.," and Figures 1 2); $\beta$ can exceed $p$, as can be seen in Figure 2, however the decrease in $\alpha_2$ during the mildly relativistic phase prevents $\beta$ from getting much larger than $p$ ."
+" This result can be extended to any observing frequency v>v,, after an appropriate modification of equation (12).", This result can be extended to any observing frequency $\nu > \nu_m$ after an appropriate modification of equation \ref{fnua}) ).
+" For instance. to consider the case of 5.>n»7rr, the right side of the equation should be multiplied by a factor of (uS).38)4), which has little effeet on the evolution of .7."," For instance, to consider the case of $\nu_c > \nu > \nu_m$ the right side of the equation should be multiplied by a factor of $(t_{obs}\Gamma^2)^{(1-3s/4)}$, which has little effect on the evolution of $\beta$ ."
+ The time scale for the increase in ./ due to sideways expansion is of order 107 (10°) for sz0 (2) (see Figure 2)., The time scale for the increase in $\beta$ due to sideways expansion is of order $10^2$ $10^3$ ) for s=0 (2) (see Figure 2).
+ Therefore this effect is smaller than that resulting from seeing the jet edge. and it extends over a much longer time.," Therefore this effect is smaller than that resulting from seeing the jet edge, and it extends over a much longer time."
+ To conclude. we wish to emphasize that for most jets propagating in a uniform ISM we are likely to see an increase to 3 of only 0.6-0.9: the remainder of the increase takes place on a long time scale. and thus is hard to detect.," To conclude, we wish to emphasize that for most jets propagating in a uniform ISM we are likely to see an increase to $\beta$ of only 0.6–0.9; the remainder of the increase takes place on a long time scale, and thus is hard to detect."
+ For jets in àmedium. s=2. 3 changes by less than about 0.5 and the transition time Πε—10°.," For jets in a, $s=2$, $\beta$ changes by less than about $0.5$ and the transition time $R_{t_e}\sim 10^3$."
+ Such a gradual increase to the afterglow light-curve power-law index is extremely difficult to detect (see Figure 2)., Such a gradual increase to the afterglow light-curve power-law index is extremely difficult to detect (see Figure 2).
+" For instance. if the edge of the Jetbecomes visible at £,5,~1 day. the difference in the optical flux at the end of 10 days with and without Jet is 0.25 mag. which can be easily missed."," For instance, if the edge of the jet becomes visible at $t_{obs}\sim 1$ day, the difference in the optical flux at the end of 10 days with and without jet is $\sim 0.25$ mag, which can be easily missed."
+" Thus. the GRBs studied by Chevalier and Li (1999), which show evidence for themodel. could in fact have had a collimated ejection of material."," Thus, the GRBs studied by Chevalier and Li (1999), which show evidence for the, could in fact have had a collimated ejection of material."
+ The optical emission of the afterglow of GRB 990510 was measured in the V. R and I bands between 0.15 and 7 days after the burst and showed the power-law index of the light-curve. ./. to have increased from (0.52+0.02 to 2.18+0.05 (Harrison et al.," The optical emission of the afterglow of GRB 990510 was measured in the V, R and I bands between 0.15 and 7 days after the burst and showed the power-law index of the light-curve, $\beta$, to have increased from $0.82 \pm 0.02$ to $2.18 \pm 0.05$ (Harrison et al."
+ 1999) or from 0.763:0.01 to 2.[0370.02 (Stanek et al., 1999) or from $0.76 \pm 0.01$ to $2.40 \pm 0.02$ (Stanek et al.
+ 1999) during a dimensionless time Ryz30 which. as deseribed previously. is not possible to obtain through the effects of the jet sideways expansion alone.," 1999) during a dimensionless time $R_{t_e}\approx 30$ which, as described previously, is not possible to obtain through the effects of the jet sideways expansion alone."
+" Therefore there must be some contribution to the light-curve steepening due to the passage of one (or both) of the spectral breaks: the synchrotron peak 7,,, and the cooling frequency 7...", Therefore there must be some contribution to the light-curve steepening due to the passage of one (or both) of the spectral breaks: the synchrotron peak $\nu_m$ and the cooling frequency $\nu_c$.
+ In Figure 3 we show a comparison between the light-curves of GRB 990510 in the V. A. 7 bands and the 8.7 GHz radio data. with a model where the cooling frequency 7. crosses the optical band at fj; day.," In Figure 3 we show a comparison between the light-curves of GRB 990510 in the $V$, $R$, $I$ bands and the 8.7 GHz radio data, with a model where the cooling frequency $\nu_c$ crosses the optical band at $t_{obs} \sim 1$ day."
+ The steepening of the light- has little dependence on the observing band because the ratio of the largest to the smallest optical wavelength is  1.5., The steepening of the light-curve has little dependence on the observing band because the ratio of the largest to the smallest optical wavelength is $\sim 1.5$ .
+ Moreover. the integration over angle spreads in time the steepening of .}. making it nearly achromatic.," Moreover, the integration over angle spreads in time the steepening of $\beta$, making it nearly achromatic."
+ An increase of Jj by ~0.8 is caused by the jet edge and the sideways expansion. and an increase of 0.25 results from the passage of r. through the observing band.," An increase of $\beta$ by $\sim0.8$ is caused by the jet edge and the sideways expansion, and an increase of 0.25 results from the passage of $\nu_c$ through the observing band."
+" A further increase of .} of ~0.15 ts caused by the passage of v,,, through the observing band at νι~0.03 day (see lower panel of Figure 3): the transition time for.) to increase by ~(3p)=1)/1 due to the 7/4, crossing is about a decade in the observer frame as a result of integration over equal arrival time surface. hence one should be careful in deducing p from ./ at early times."," A further increase of $\beta$ of $\sim 0.15$ is caused by the passage of $\nu_m$ through the observing band at $t_{obs} \sim 0.03$ day (see lower panel of Figure 3); the transition time for $\beta$ to increase by $\sim(3p-1)/4$ due to the $\nu_m$ crossing is about a decade in the observer frame as a result of integration over equal arrival time surface, hence one should be careful in deducing $p$ from $\beta$ at early times."
+ All these together give rise to a light-curve that 1s consistent with the data., All these together give rise to a light-curve that is consistent with the data.
+ The model is also consistent with the HST V- observation carried out at about à month after the burst (Fruchter et al 1999)., The model is also consistent with the HST $V$ observation carried out at about a month after the burst (Fruchter et al 1999).
+ The parameters for the fit are given in the caption for fig., The parameters for the fit are given in the caption for fig.
+ 3 which yields the energy in the burst to be 10? ere., 3 which yields the energy in the burst to be $^{49}$ erg.
+" Correcting for the radiative losses the energy in the burst increases by a factor of a few to <10°"" erg.", Correcting for the radiative losses the energy in the burst increases by a factor of a few to $\lta 10^{50}$ erg.
+ We estimate the uncertainty in model parameters by varying them m such a way that the numerically calculated light-curve Hes within 3-0 of the observed data points., We estimate the uncertainty in model parameters by varying them in such a way that the numerically calculated light-curve lies within $\sigma$ of the observed data points.
+" We find the uncertainty in the jet opening angle and the burst energy to be a factor of two. and €,. n and ej; are found to be uncertain by factors of about 4. 40 and 7 respectively: we note that the radio observations are very important in constraining the model parameters."," We find the uncertainty in the jet opening angle and the burst energy to be a factor of two, and $\epsilon_e$, $n$ and $\epsilon_B$ are found to be uncertain by factors of about 4, 40 and 7 respectively; we note that the radio observations are very important in constraining the model parameters."
+ The electron index p is constrained by the observed .} before and after the | day break: the error in p is ~54., The electron index $p$ is constrained by the observed $\beta$ before and after the $\sim 1$ day break; the error in $p$ is $\sim 5\%$.
+ The optical emission of the afterelow of GRB 990510 can also be explained by a model where the synchrotron peak frequency crosses the observed band at ~0.1 day., The optical emission of the afterglow of GRB 990510 can also be explained by a model where the synchrotron peak frequency crosses the observed band at $\sim 0.1$ day.
+ Its effect on Jj persists for up to ~1 day and yields an increase of ὐ of ~0.5 during the early observations., Its effect on $\beta$ persists for up to $\sim 1$ day and yields an increase of $\beta$ of $\sim 0.5$ during the early observations.
+ The parameters for the second model differ from the one described above (see Figure 3) somewhat., The parameters for the second model differ from the one described above (see Figure 3) somewhat.
+ In particular. ει. is larger by a factor of two. the energy per solid angle is smaller by a factor of two. and Jy is ~20% larger.," In particular, $\epsilon_e$ is larger by a factor of two, the energy per solid angle is smaller by a factor of two, and $\theta_0$ is $\sim 20\%$ larger."
+ One of the main results of this work is to show that afterglows from well collimated Gamma-Ray Burst remnants going off in a medium with density decreasing as r? show little evidence for light-curve steepening due to jet edge and sideways expansion., One of the main results of this work is to show that afterglows from well collimated Gamma-Ray Burst remnants going off in a medium with density decreasing as $r^{-2}$ show little evidence for light-curve steepening due to jet edge and sideways expansion.
+ This could explain the lack of breaks in the afterglows of GRB 980326 and GRB 980519. which Chevalier Li (1999) found to offer support for themodel.," This could explain the lack of breaks in the afterglows of GRB 980326 and GRB 980519, which Chevalier Li (1999) found to offer support for the."
+ Jets can perhaps be detected by the measurement of time dependent polarization., Jets can perhaps be detected by the measurement of time dependent polarization.
+ In a collimated outflow the sharpest break in the light-curve is produced in a uniform density eircum-stellar medium. and is associated with the edge of the jet coming within the relativistic beaming cone (the effect).," In a collimated outflow the sharpest break in the light-curve is produced in a uniform density circum-stellar medium, and is associated with the edge of the jet coming within the relativistic beaming cone (the )."
+ The magnitude of this break is ~0.7 (0.4) for a uniform ISM (wind model) and occurs over about | decade (2 decades) in time., The magnitude of this break is $\sim0.7$ (0.4) for a uniform ISM (wind model) and occurs over about 1 decade (2 decades) in time.
+ Further steepening of the light-curve. associated with the sideways expansion of the jet. oceurs on a much longer time seale of Ry ~10? (104). ie. weeks to months.," Further steepening of the light-curve, associated with the sideways expansion of the jet, occurs on a much longer time scale of $R_{t_e}\sim$ $^2$ $^4$ ), i.e. weeks to months."
+ The power-law index for the light-curve of GRB 990510 increased between days 0.8 and 3 by about 1.35., The power-law index for the light-curve of GRB 990510 increased between days 0.8 and 3 by about 1.35.
+ This is too large and too fast to result from jet edge sideways expansion effects., This is too large and too fast to result from jet edge sideways expansion effects.
+ However. the observations can be explained if either the cooling or the synchrotron peak frequency passedthrough the observing band at about| or 0.1 day. respectively.," However, the observations can be explained if either the cooling or the synchrotron peak frequency passedthrough the observing band at about1 or 0.1 day, respectively."
+" Models that are consistent both the optical and radio data of this afterglow have an opening angle of ~5"" andenergy in the explosion is X10?"" ere (see Figure 3).", Models that are consistent both the optical and radio data of this afterglow have an opening angle of $\sim 5^o$ andenergy in the explosion is $\lta 10^{50}$ erg (see Figure 3).
+" For the afterglow of GRB 990123 the power-law index of the light-curve increased by 0.55 between days 1.5 and 3. which can be explained by the effect alone Rees 1999),"," For the afterglow of GRB 990123 the power-law index of the light-curve increased by 0.55 between days 1.5 and 3, which can be explained by the alone Rees 1999)."
+The Milky Way GCS is no exception to this general treud.,The Milky Way GCS is no exception to this general trend.
+ Because it is a completely uncontamiuatecd sample of clusters with well established metallicities. it provides what is still perliaps tle best classic example of a =.bimodal NDF (see Figure 1).," Because it is a completely uncontaminated sample of clusters with well established metallicities, it provides what is still perhaps the best classic example of a bimodal MDF (see Figure 1)."
+ The shape of the MDF aud its close connection with the cluster kinematics was first clearly established by Zinn(1985) and has held up through many subsequeut analyses (see.e.g.Harris2001.forareview)..," The shape of the MDF and its close connection with the cluster kinematics was first clearly established by \citet{zin85} and has held up through many subsequent analyses \citep[see, e.g.][for a review]{h01}."
+ The MDF. in its usual plotted form as number per unit [Fe/H] (left panel of Figure 1). is often approximated by a pair of Caussian [uuctious (οἱ.," The MDF, in its usual plotted form as number per unit [Fe/H] (left panel of Figure 1), is often approximated by a pair of Gaussian functions (cf."
+ the references cited above)., the references cited above).
+ These numerical fits have uo pliysical significance (ara have not usually been claimed to). aud eeive little insighte into the evolutionary history of the system except lor the strong lint that a distiuct two-pliase formation history is required.," These numerical fits have no physical significance (and have not usually been claimed to), and give little insight into the evolutionary history of the system except for the strong hint that a distinct two-phase formation history is required."
+ Iu this brief paper. we take oue step further iuto interpreting these bimocal NDFs using some simple and. pliysically plausible assumptions.," In this brief paper, we take one step further into interpreting these bimodal MDFs using some simple and physically plausible assumptions."
+" Standard first-order chemical evolution models. such as the one-zoue closed-box or “simple model"" (Pagel&Patchett1975: work reasonably well at inatehiug the broad. uniuodal MDFs of the halo field-star populatious in the Milky Way (e.g.Pagel&Patchett1972:Hartwick1976:ναιNorris1991:Prautzos2003) and iu M31 (Durrell.Harris.&Priteliet2001) with appropriate choices of the elective iucleosvntlieticH vieldH veg."," Standard first-order chemical evolution models, such as the one-zone closed-box or “simple model” \citep{pag75,bin98} work reasonably well at matching the broad, unimodal MDFs of the halo field-star populations in the Milky Way \citep[e.g.][]{pag75,har76,rya91,pra03} and in M31 \citep{dur01} with appropriate choices of the effective nucleosynthetic yield $y_{\rm eff}$."
+" ExteusiousH to the simpleH model to incorporateH some smoothly varyingH eas lnflow (the ""accretiug-box model) have recently beeu used to model the unimodal. metal-rich DF of the halo stars in NGC 5128 (Harris&Harris2002)."," Extensions to the simple model to incorporate some smoothly varying gas inflow (the “accreting-box” model) have recently been used to model the unimodal, metal-rich MDF of the halo stars in NGC 5128 \citep{har02}."
+. An early. brief phase of primordial gas inflow has even been applied to explain the relatively sinall uumber of theverv lowest-metallicity stars ([Fe/H] «—3) in the Milky Way halo (Prantzos2003).," An early, brief phase of primordial gas inflow has even been applied to explain the relatively small number of the lowest-metallicity stars ([Fe/H] $< -3$ ) in the Milky Way halo \citep{pra03}."
+. Accreting-box models can be thought “as approximate versions of more exteusive hierarchical-uiergiug models of galaxy. formation. iu “=‘hieh a large population of sinall pregalactic gas clouds merges successively iuto bigger aud biggerMDver ποmleces while star formation within the clouds goes on simultaneously (see.foranapplicationoftheCALFORMcodetocaseNGC 5128)...," Accreting-box models can be thought of as approximate versions of more extensive hierarchical-merging models of galaxy formation, in which a large population of small pregalactic gas clouds merges successively into bigger and bigger pieces while star formation within the clouds goes on simultaneously \citep[see, e.g.][for an application of the GALFORM code
+to the case of NGC 5128]{bea03}."
+ By contrast. the distiuctly bimodal MDEFs of globular cluster systems do uot easily fit into the continuous sequence of starforming events in normal hierarchical-imereing or accreting-box moclels.," By contrast, the distinctly bimodal MDFs of globular cluster systems do not easily fit into the continuous sequence of starforming events in normal hierarchical-merging or accreting-box models."
+ Some new feature of the history needs to be invoked. either to boost the formation of metal-poor clusters at early times. or tubibit the slightly later formation of intermectiate-metallicity clusters.," Some new feature of the history needs to be invoked, either to boost the formation of metal-poor clusters at early times, or inhibit the slightly later formation of intermediate-metallicity clusters."
+values that fully agrees with the latest estimates from Tycner et al.(2011).,values that fully agrees with the latest estimates from Tycner et al.(2011).
+ This is reasonable taking into account that their measurements dominates the sample both in term of number (96 out of 150) and uncertainties., This is reasonable taking into account that their measurements dominates the sample both in term of number (96 out of 150) and uncertainties.
+" Finally, we can conclude that the next periastron should take place around July 5, 2011 (x4 days)."," Finally, we can conclude that the next periastron should take place around July 5, 2011 $\pm$ 4 days)."
+ We can constrain the total mass of the ó Sco binary system using Kepler's third law and our estimation of the binary period and semi major-axis., We can constrain the total mass of the $\delta$ Sco binary system using Kepler's third law and our estimation of the binary period and semi major-axis.
+" Since this last parameter is given in angular units (i.e. mas), the binary system distance is also needed for the mass calculation."," Since this last parameter is given in angular units (i.e. mas), the binary system distance is also needed for the mass calculation."
+ Using the distance inferred from the Hipparcos parallax (Perryman et al., Using the distance inferred from the Hipparcos parallax (Perryman et al.
+" 1997), i.e. d123 ppc, we derived a total mass of Ma;g=15.2 +5Mo."," 1997), i.e. $d = 123$ pc, we derived a total mass of $M_{\rm A+B} = 15.2 \pm 5$ $M_\odot$."
+" On the other hand, using the distance deduced from van Leeuween (2007) revised Hipparcos parallax, i.e. d = 150pc, we obtained a total mass MA,p =27.7 + 10 Mo."," On the other hand, using the distance deduced from van Leeuween (2007) revised Hipparcos parallax, i.e. d = 150pc, we obtained a total mass $_{\rm A+B}$ =27.7 $\pm$ 10 $M_\odot$."
+ The second estimation is in agreement with the total mass determined previously by Tango et al. (, The second estimation is in agreement with the total mass determined previously by Tango et al. (
+"2009), i.e. €10 Mo.","2009), i.e. $\pm 10$ $M_\odot$."
+" From physical modeling of the spectral energy distribution, Carciofi et al. ("," From physical modeling of the spectral energy distribution, Carciofi et al. ("
+2006) estimated the mass of the primary to be on the order of ΜΑ=14 Mg.,2006) estimated the mass of the primary to be on the order of $M_{\rm A} = 14$ $M_\odot$.
+" Assuming that both stars have the same age and are currently on the main sequence, and considering the flux ratio measured in the K band, i.e. Fp/F4=0.07/0.93=0.075+ 0.012, the secondary should also be a quite massive hot star."," Assuming that both stars have the same age and are currently on the main sequence, and considering the flux ratio measured in the K band, i.e. $F_{B}/F_{A} = 0.07/0.93 = 0.075 \pm 0.012$ , the secondary should also be a quite massive hot star."
+" Thus, it seems that Ma,p=15.25 Mao is quite unrealistic."," Thus, it seems that $M_{\rm A+B} = 15.2 \pm 5$ $M_\odot$ is quite unrealistic."
+ On the other hand Maip=27.7€ 1OMo would implies that both components have approximatively the same mass., On the other hand $M_{\rm A+B} = 27.7 \pm 10$ $M_\odot$ would implies that both components have approximatively the same mass.
+" Finally, Tango et al. ("," Finally, Tango et al. ("
+2009) give an estimation of the secondary mass of Mg=8+3.6 Me.,2009) give an estimation of the secondary mass of $M_{\rm B} = 8 \pm 3.6$ $M_\odot$.
+" From the mass ratio they calculated, i.e. MA/Mg=1.957+0.011 we can deduce the Luminosity ratio using Mass-Luminosity relation from Griffiths et al. ("," From the mass ratio they calculated, i.e. $M_{\rm A}/M_{\rm B} = 1.957 \pm 0.011$ we can deduce the Luminosity ratio using Mass-Luminosity relation from Griffiths et al. ("
+"1988): Thus, (ane)we can deduce that Lg/L4=0.095+0.01.","1988): Thus, we can deduce that $L_{\rm B}/L_{\rm A} = 0.095 \pm 0.01$."
+" This is compatible with our K band flux ratio measurements, i.e. Fpg/F4=0.075+0.012."," This is compatible with our K band flux ratio measurements, i.e. $F_{\rm B}/F_{\rm A} = 0.075 \pm 0.012$."
+ This gives an additional indication that the two components have similar spectral classes., This gives an additional indication that the two components have similar spectral classes.
+" Considering the measured flux ratio, this would imply that the secondary spectral class ranges between B2V and Β4Ν. During our observing campaign, 6 Sco was observed three times with AMBER in medium resolution (MR): in 2007, 2009, and 2010."," Considering the measured flux ratio, this would imply that the secondary spectral class ranges between B2V and B4V. During our observing campaign, $\delta$ Sco was observed three times with AMBER in medium resolution (MR): in 2007, 2009, and 2010."
+ All the observations were centered on the fist half of the K band um) in order to study the circumstellar environment geometry and kinematics in the Bry emission line., All the observations were centered on the fist half of the K band $\mu$ m) in order to study the circumstellar environment geometry and kinematics in the $\gamma$ emission line.
+" The quality of the data taken in 2007 was too poor to detect any signal in the line, whereas, as seen in Fig}, visibility and phase variations were successfully detectedin 2009."," The quality of the data taken in 2007 was too poor to detect any signal in the line, whereas, as seen in Fig \ref{visMR}, visibility and phase variations were successfully detectedin 2009."
+" We also discovered that two other K-band emission lines were strong enough to support the kinematics study: the line located at jum, and the Bró line around um. However, the latter is located too close to the edge of the so that the data SNR is significantly lower than for the two other lines."," We also discovered that two other K-band emission lines were strong enough to support the kinematics study: the line located at $\mu$ m, and the $\delta$ line around $\mu$ m. However, the latter is located too close to the edge of the K-band so that the data SNR is significantly lower than for the two other lines."
+" We should have decided to study the circumstellar environment geometry and kinematics in MR mode, as it was already done for two other classical Be stars: α Ara (Meilland et al."," We should have decided to study the circumstellar environment geometry and kinematics in MR mode, as it was already done for two other classical Be stars: $\alpha$ Ara (Meilland et al."
+ 2007a) and x CMa (Meilland et al., 2007a) and $\kappa$ CMa (Meilland et al.
+ 2007b)., 2007b).
+" However, since the star was bright enough, we decided to take advantage of the eight times higher spectral resolution offered by the HR mode."," However, since the star was bright enough, we decided to take advantage of the eight times higher spectral resolution offered by the HR mode."
+" With this resolution of R=12000, kinematics details of about ss“! are achievable."," With this resolution of $R = 12000$, kinematics details of about $^{-1}$ are achievable."
+" Consequently, in 2010, 6 Sco was observed twice with the VLTI/AMBER in HR mode."," Consequently, in 2010, $\delta$ Sco was observed twice with the VLTI/AMBER in HR mode."
+ The first observations were centered on the Bry line and the second on the line., The first observations were centered on the $\gamma$ line and the second on the line.
+" Unfortunately, we did not manage to calibrate the observation properly."," Unfortunately, we did not manage to calibrate the observation properly."
+" However, this does not affect either the differential visibilities and phases or the closure phase."," However, this does not affect either the differential visibilities and phases or the closure phase."
+" Finally, quasi-simultaneous medium-spectral resolution observations centered on Ha were carried out with the CHARA/VEGA instrument."," Finally, quasi-simultaneous medium-spectral resolution observations centered on $\alpha$ were carried out with the CHARA/VEGA instrument."
+" The differential visibilities and phases, the closure phases (only for VLTI/AMBERdata), and line profiles for these observations are presented in Fig D]."," The differential visibilities and phases, the closure phases (only for VLTI/AMBERdata), and line profiles for these observations are presented in Fig \ref{lines}. ."
+ Their morphology is similar to those shown by Meilland et al. (, Their morphology is similar to those shown by Meilland et al. (
+2007a) for the Be,2007a) for the Be
+30 s. These MBPs were found to have an average velocity of approximately 1 +. and a mean lifetime of 91 s. The highest velocity recorded in this study was 7 Loa value comparable to the photospheric sound speed. while the longest MDP lifetime observed was approximately 20 minutes.,"30 s. These MBPs were found to have an average velocity of approximately 1 $^{-1}$, and a mean lifetime of 91 s. The highest velocity recorded in this study was 7 $^{-1}$, a value comparable to the photospheric sound speed, while the longest MBP lifetime observed was approximately 20 minutes."
+ Similar values were obtained when (he same methodology was applied to the numerical simulations., Similar values were obtained when the same methodology was applied to the numerical simulations.
+" A total of 721 AIBPs were detected in the 16""x 200 G simulation. with an average velocity of 0.6 toa maximum of 4.1 +. an average lifetime of 103.2 s. ancl a maximum of just over 11 minutes."," A total of 721 MBPs were detected in the $16''\times16''$ 200 G simulation, with an average velocity of 0.6 $^{-1}$, a maximum of 4.1 $^{-1}$, an average lifetime of 108.2 s, and a maximum of just over 11 minutes."
+ In the 400 G simulation. we detected 1024 MBPs. with an average velocity of 0.65 | anda maximum of 6.2 kms... an average lifetime of 168.4 s and a maximum of nearly 17 minutes.," In the 400 G simulation, we detected 1024 MBPs, with an average velocity of 0.65 $^{-1}$ and a maximum of 6.2 $^{-1}$, an average lifetime of 168.4 s and a maximum of nearly 17 minutes."
+ Histograms indicating the distributions of velocities for all three data sets are shown in Figure ία D). with the results sumnmarised in Table 1.," Histograms indicating the distributions of velocities for all three data sets are shown in Figure \ref{Fig2}( (d–f), with the results summarised in Table 1."
+ A dependence of the velocity distribution on the average magnetic field strength. (and thus on the solar magnetic activity level) suggests an alternative. and novel. way of estimating the net magnetic flux in (he solar photosphere.," A dependence of the velocity distribution on the average magnetic field strength (and thus on the solar magnetic activity level) suggests an alternative, and novel, way of estimating the net magnetic flux in the solar photosphere."
+ Hence. as the best agreement with our observations is lound for the 400 G simulation. we can infer the average magnetic [lux of our observations as being close to 400 G. Small discrepancies. between our observed and simulated datasets may be due to the large MDP sample present in our observations (6236 versus 1024).," Hence, as the best agreement with our observations is found for the 400 G simulation, we can infer the average magnetic flux of our observations as being close to 400 G. Small discrepancies between our observed and simulated datasets may be due to the large MBP sample present in our observations (6236 versus 1024)."
+ The distribution of velocities. for both the observed and simulated datasets. are shown as histograms in Figure 4((d[).," The distribution of velocities, for both the observed and simulated datasets, are shown as histograms in Figure \ref{Fig2}( (d–f)."
+ These demonstrate that the majority of AIBPs have velocities between 0 and 1 !. while some MDPs (i.e. in the observations. in the 200 G simulated data and in the 400 G) have velocities in excess of 2 |.," These demonstrate that the majority of MBPs have velocities between 0 and 1 $^{-1}$, while some MBPs (i.e. in the observations, in the 200 G simulated data and in the 400 G) have velocities in excess of 2 $^{-1}$ ."
+ The considerable fraction of MBPs with velocities greater than 2 J| is important. as rapid bursts in their motion may induce magnetosonic kink waves (Chouclhurietal.1993).," The considerable fraction of MBPs with velocities greater than 2 $^{-1}$ is important, as rapid bursts in their motion may induce magnetosonic kink waves \citep{Choud93}."
+.. As shown previously. these kink waves may act as a conduit for imparting energy into (he upper solar atmosphere.," As shown previously, these kink waves may act as a conduit for imparting energy into the upper solar atmosphere."
+ Although there is considerable agreement between the results obtained in the observed and simulated data. there are some slight discrepancies that need to be addressed for a more complete picture.," Although there is considerable agreement between the results obtained in the observed and simulated data, there are some slight discrepancies that need to be addressed for a more complete picture."
+ For example. the maximum velocity. (he maximum lifetime and the proportion of AIBPs with velocities greater than 2 1 are all smaller in the simulated datasets.," For example, the maximum velocity, the maximum lifetime and the proportion of MBPs with velocities greater than 2 $^{-1}$ are all smaller in the simulated datasets."
+ A possible reason is that the net (or average) magnetic field applied within the computational domain remains constant throughout the 200 G or 400 G time series. whereas the [lux present in the observed data may constantly change wilh time.," A possible reason is that the net (or average) magnetic field applied within the computational domain remains constant throughout the 200 G or 400 G time series, whereas the flux present in the observed data may constantly change with time."
+ This dynamic nature ol flux generation and annihilation could directly lead to the higher velocity distributions detected for the observed. MDPs., This dynamic nature of flux generation and annihilation could directly lead to the higher velocity distributions detected for the observed MBPs.
+ An alternative algorithm based on local correlation tracking (LCT:Welschetal.Matlochetal.2010) wasapplied to the observations and simulated time series. ancl vielded sinilar velocity characteristics to within 0.31 +.," An alternative algorithm based on local correlation tracking \citep[LCT;][]{Wel04, Mat10} wasapplied to the observations and simulated time series, and yielded similar velocity characteristics to within 0.3$^{-1}$ ."
+ A significant number of MBPs. 221% ," A significant number of MBPs, $\approx$ "
+Our analvsis shows that nearly all azimuthal wavenumbers m are unstable under reasonable conditions.,Our analysis shows that nearly all azimuthal wavenumbers $m$ are unstable under reasonable conditions.
+ Although modes with higher values of n grow more rapidly. we expect that these modes will saturate at relatively small aanplitudes.," Although modes with higher values of $m$ grow more rapidly, we expect that these modes will saturate at relatively small amplitudes."
+ The ΙΟ modes. on the other hand. are likely to grow to large aanplitude and are therefore of most interest for understanding QPOs.," The $m$ modes, on the other hand, are likely to grow to large amplitude and are therefore of most interest for understanding QPOs."
+" For q in the range 3/2 (o 2 (Ixepleriau to constant angular momentum). aud reasonable assumptions about the density profile. we have shown that the observed mode Irequency tends to be of order 1(2,,."," For $q$ in the range 3/2 to 2 (Keplerian to constant angular momentum), and reasonable assumptions about the density profile, we have shown that the observed mode frequency tends to be of order $m \Omega_m$."
+ Thus. in the simplest version of (he model. we expect the observed QPO frequencies to be in the ratio 1:2:3. ete.," Thus, in the simplest version of the model, we expect the observed QPO frequencies to be in the ratio 1:2:3, etc."
+ However. as we showed in34 (see Fig.," However, as we showed in (see Fig."
+ 4). il may often be the case that the m=1 mode is stable and that only modes with m>2 are unstable.," 4), it may often be the case that the $m=1$ mode is stable and that only modes with $m\geq2$ are unstable."
+ This should happen whenever the effective gravity in the racial direction is weak. ie. When (he gas pressure in (he disk is low.," This should happen whenever the effective gravity in the radial direction is weak, i.e., when the gas pressure in the disk is low."
+ In this case. the QDPO frequencies should be roughly in the ratio 2:3:4. ete.," In this case, the QPO frequencies should be roughly in the ratio 2:3:4, etc."
+ It is interesting that a frequency ratio 2:3 is lvequently seen in both black hole and neutron star svstems20022)., It is interesting that a frequency ratio 2:3 is frequently seen in both black hole and neutron star systems.
+ We have emphasized that. in our opinion. QPOs should be produced by the growing modes in the disk.," We have emphasized that, in our opinion, QPOs should be produced by the growing modes in the disk."
+ Our belief is based on the fact that the amplitude of intensity fhictuations observed in QPOs is often quite large., Our belief is based on the fact that the amplitude of intensity fluctuations observed in QPOs is often quite large.
+ Observations also show Chat QPOs have a finite frequency width. which indicates that the unstable modes in (he disk have a [finite lile time (hence the term quasi-periodic oscillations).," Observations also show that QPOs have a finite frequency width, which indicates that the unstable modes in the disk have a finite life time (hence the term quasi-periodic oscillations)."
+" Apart [rom the rotation period. there are two natural time scales in the disk: the time scale associated with the effective gravity. which is ev 1/Qey. ancl the viscous (ime scale. which is ~r/e,. where ve, is the mean racial velocity of the disk fluid."," Apart from the rotation period, there are two natural time scales in the disk: the time scale associated with the effective gravity, which is $\sim 1/\Omega_{\rm eff}$ and the viscous time scale, which is $\sim
+r/v_r$, where $v_r$ is the mean radial velocity of the disk fluid."
+" For a standard disk. Quy~ ή, where c; is the sound speed."," For a standard disk, $\Omega_{\rm eff} \sim c_s/r$ , where $c_s$ is the sound speed."
+" Since C.>C,. the lime scale 1/4 is generally much shorter than ο."," Since $c_s \gg v_r$, the time scale $1/\Omega_{\rm eff}$ is generally much shorter than $r/v_r$."
+ Therefore. the likely lifetime of blobs created by the instability is 1/OQeg.," Therefore, the likely lifetime of blobs created by the instability is $1/\Omega_{\rm eff}$."
+ Once a mass blob is formed. it will diit toward the central object under Cae action of gravity. causing a displacement in the center frequency and a width to the QPO feature in the power spectrim.," Once a mass blob is formed, it will drift toward the central object under the action of gravity, causing a displacement in the center frequency and a width to the QPO feature in the power spectrum."
+" The width is likely to be Af~Oop/2ac(h/r)(Q,,/21). where fr is the vertical thickness of the disk. and the displacement speed is approximately. df/dl~(Qe/28)f."," The width is likely to be $\Delta f \sim \Omega_{\rm eff}/2\pi \sim
+(h/r)(\Omega_m/2\pi)$, where $h$ is the vertical thickness of the disk, and the displacement speed is approximately $df/dt \sim 
+(\Omega_{\rm eff}/2\pi) f$."
+ Another issue concerns how the presence of à nonaxisvmmetrie mode translates to a time modulation of (he observed flux., Another issue concerns how the presence of a nonaxisymmetric mode translates to a time modulation of the observed flux.
+" Two possibilities are likely,", Two possibilities are likely.
+ One is that the svstem is viewed in a nearly edge-on configuration so that the accreting star eclipses the far side of the disk., One is that the system is viewed in a nearly edge-on configuration so that the accreting star eclipses the far side of the disk.
+ Then. as bright ancl faint seements of the disk are successively eclipsed the signal at the observer will be modulated.," Then, as bright and faint segments of the disk are successively eclipsed the signal at the observer will be modulated."
+ The other possibility. which also requires fairly high inclination. is that the motion of the gas is relativistic and (he observed. signal is dominated by the blue-shifted segment of the disk.," The other possibility, which also requires fairly high inclination, is that the motion of the gas is relativistic and the observed signal is dominated by the blue-shifted segment of the disk."
+ Once again. as the nonaxisvimnietric pattern rotates. (he signal will oscillate.," Once again, as the nonaxisymmetric pattern rotates, the signal will oscillate."
+ Both mechanisms require that the bright and faint patches on the disk should have large areas.since otherwise the fractional modulationof the observed. X-ray παν," Both mechanisms require that the bright and faint patches on the disk should have large areas,since otherwise the fractional modulationof the observed X-ray flux"
+"We expand each of (he wave-field variables into an incident component (subscript ""inc) and a seattered component (subscript ""sc): Dv inserting the above expansions into equations (11))-(15)) and retaining the terms of order ε we obtain The terms on the right-hand side of the above equations act as sources for the scattered waves: (his is (he Dorn approximation.",We expand each of the wave-field variables into an incident component (subscript “inc”) and a scattered component (subscript “sc”): By inserting the above expansions into equations \ref{cont}) \ref{temp}) ) and retaining the terms of order $\epsilon$ we obtain The terms on the right-hand side of the above equations act as sources for the scattered waves: this is the Born approximation.
+ Writing all wave variables in (he form of equation (16)) and using the fact that the magnetic field is solenoidal. (he above equations reduce to a forced Ποο] equation for the (&..0) Fourier component of the scattered pressure lield. p: where Ap is the two-dimensional Laplacian with respect to r and the source function S(r) is given by The first term in S is due to the density jump at the tube boundary. and the other two terms are due to the direct effect of the Lorentz lorce on the wave.," Writing all wave variables in the form of equation \ref{eq.expand}) ) and using the fact that the magnetic field is solenoidal, the above equations reduce to a forced Helmholtz equation for the $(k_z, \omega)$ Fourier component of the scattered pressure field, $\tilde{p}_{\rm sc}$: where $\Delta_\br$ is the two-dimensional Laplacian with respect to $\br$ and the source function $\tilde{S}(\br)$ is given by The first term in $\tilde{S}$ is due to the density jump at the tube boundary and the other two terms are due to the direct effect of the Lorentz force on the wave."
+ For the incoming wave. Pine. We lake the same plane wave as in the exact solution (Eq. |17]," For the incoming wave, $\tilde{p}_{\rm inc}$, we take the same plane wave as in the exact solution (Eq. \ref{eq.definc}] ])."
+"In this section, we will test the analytical predictions of the previous section by a coagulation/fragmentation code (?,, see also ?)).","In this section, we will test the analytical predictions of the previous section by a coagulation/fragmentation code \citealp{Birnstiel:2010p9709}, see also \citealp{Brauer:2008p215}) )."
+ The code solves for the time evolution of the grain size distribution using an implicit integration scheme., The code solves for the time evolution of the grain size distribution using an implicit integration scheme.
+ This enables us to find the steady-state distribution by using large time steps., This enables us to find the steady-state distribution by using large time steps.
+" In this way, the time evolution is not resolved, but the state distribution is reliably and very quickly derived."," In this way, the time evolution is not resolved, but the steady-state distribution is reliably and very quickly derived."
+ We start out with the simplest case of a constant kernel and then — step by step — approach a more realistic scenario (in the context of a protoplanetary disk)., We start out with the simplest case of a constant kernel and then – step by step – approach a more realistic scenario (in the context of a protoplanetary disk).
+" In Sect. ??,,"," In Sect. \ref{sec:distri:diskdistris},"
+ we will then consider a kernel taking into account relative velocities of Brownian motion and turbulent velocities and also a fragmentation probability which depends on the masses and the relative velocities of the colliding particles., we will then consider a kernel taking into account relative velocities of Brownian motion and turbulent velocities and also a fragmentation probability which depends on the masses and the relative velocities of the colliding particles.
+" In the following section, we consider the case of a constant kernel and will include fragmentation above particle sizes of 1 mm because this represents an instructive test case."," In the following section, we consider the case of a constant kernel and will include fragmentation above particle sizes of 1 mm because this represents an instructive test case."
+ We iteratively solve for the steady-state size distribution between coagulation and fragmentation., We iteratively solve for the steady-state size distribution between coagulation and fragmentation.
+" The outcome of these simulations for a constant kernel (i.e., ν= 0) are power-law distributions where the slope of the distribution depends on the fragmentation law (the slope £, see Eq. 10))."," The outcome of these simulations for a constant kernel (i.e., $\nu=0$ ) are power-law distributions where the slope of the distribution depends on the fragmentation law (the slope $\xi$, see Eq. \ref{eq:distri:frag_powerlaw}) )."
+ Figure 2 shows the corresponding size distributions for some of the different fragment distributions: the steepest distribution corresponds to the case of €=0.5., Figure \ref{fig:distri:spectra} shows the corresponding size distributions for some of the different fragment distributions: the steepest distribution corresponds to the case of $\xi=0.5$.
+" For larger £-values, the slope of the mass distribution flattens."," For larger $\xi$ -values, the slope of the mass distribution flattens."
+" In all cases, a bump develops towards the upper end of the distributions."," In all cases, a bump develops towards the upper end of the distributions."
+ The reason for this is the following: grains typically grow mostly through collisions with similar-sized or larger particles., The reason for this is the following: grains typically grow mostly through collisions with similar-sized or larger particles.
+" Since the distribution is truncated at the upper end (defined asdmax,, see also Eq. 43)),"," Since the distribution is truncated at the upper end (defined as, see also Eq. \ref{eq:distri:a_max}) ),"
+" particles close to the upper end lack larger collision partners, the growth rate at these sizes would decrease if the distribution keeps its power-law nature."," particles close to the upper end lack larger collision partners, the growth rate at these sizes would decrease if the distribution keeps its power-law nature."
+" This, in turn means that the flux could not be constant belowdmax."," This, in turn means that the flux could not be constant below."
+". To keep a steady state, the number of particles at that point has to increase in order to replace the missing collision partners at larger sizes."," To keep a steady state, the number of particles at that point has to increase in order to replace the missing collision partners at larger sizes."
+" Figure 3shows how the slope of the resulting distribution depends on the fragmentation slope ὅ, where the three previously discussed regimes can be identified:"," Figure \ref{fig:distri:exponent} shows how the slope of the resulting distribution depends on the fragmentation slope $\xi$ , where the three previously discussed regimes can be identified:"
+central stellar velocity dispersion in the dwarf.,central stellar velocity dispersion in the dwarf.
+ When the left hand side of the equation becomes larger than the right hand side gas is removed from the galaxy., When the left hand side of the equation becomes larger than the right hand side gas is removed from the galaxy.
+ It is assumed that the stripping is instantaneous and that the intergalactic gas density encountered by the moving galaxy remains constant., It is assumed that the stripping is instantaneous and that the intergalactic gas density encountered by the moving galaxy remains constant.
+ The use of this criterion as a relatively good approximation is supported by numerical models of gas stripping in clusters (Mori&Burkert 2000).., The use of this criterion as a relatively good approximation is supported by numerical models of gas stripping in clusters \citep{2000ApJ...538..559M}.
+" Using the IGM densities measured here we can explore whether ram-pressure stripping in groups is strong enough to remove significant reservoirs of neutral interstellar gas from dwarf galaxies, a key step in producing dSphs."," Using the IGM densities measured here we can explore whether ram-pressure stripping in groups is strong enough to remove significant reservoirs of neutral interstellar gas from dwarf galaxies, a key step in producing dSphs."
+" To model dwarf galaxies we choose the range of total mass, mass, o,, and interstellar neutral gas density shown in Table 2.."," To model dwarf galaxies we choose the range of total mass, mass, $\sigma_v$, and interstellar neutral gas density shown in Table \ref{tab:dwarfs}."
+ We compute the combination of galaxy velocity and IGM density necessary to instantaneously strip the neutral gas from a given dwarf galaxy using Equation 3 and show these curves in Figure 6 for the three dwarfs differentiating them by mass., We compute the combination of galaxy velocity and IGM density necessary to instantaneously strip the neutral gas from a given dwarf galaxy using Equation 3 and show these curves in Figure \ref{fig:dwarfs} for the three dwarfs differentiating them by mass.
+ The density and velocity characteristics of four groups from this paper are plotted to show which will have the necessary conditions to strip from the dwarfs., The density and velocity characteristics of four groups from this paper are plotted to show which will have the necessary conditions to strip from the dwarfs.
+" Comparing these stripping curves for generic dwarf galaxies to the IGM densities and average velocities of groups in this paper, we see that the dwarfs with Maur=>105Μο Mios=>10° Mo) and Mg;=106Mo (Mi=107 Mo) are likely to be stripped in all four groups shown."," Comparing these stripping curves for generic dwarf galaxies to the IGM densities and average velocities of groups in this paper, we see that the dwarfs with $M_{HI}=10^5\ \msol$ $M_{tot}=10^6\ \msol$ ) and $M_{HI}=10^6\ \msol$ $M_{tot}=10^7\ \msol$ ) are likely to be stripped in all four groups shown."
+" There are many additional factors that could increase the effectiveness of ram-pressure stripping including feedback from star-formation, the cosmic ionizing background, the orbital path of the galaxy, and a fluctuating density distribution."," There are many additional factors that could increase the effectiveness of ram-pressure stripping including feedback from star-formation, the cosmic ionizing background, the orbital path of the galaxy, and a fluctuating density distribution."
+ The densities we derive here suggest that ram-pressure may be a critically important process affecting the gas content of dwarf galaxies., The densities we derive here suggest that ram-pressure may be a critically important process affecting the gas content of dwarf galaxies.
+ The question of whether these dwarfs would be stripped in the Local or nearby groups with confirmed morphology-density relations is slightly more complicated., The question of whether these dwarfs would be stripped in the Local or nearby groups with confirmed morphology-density relations is slightly more complicated.
+ The Local Group is an order of magnitude less massive (~1—2x10!?Mo; Karachentsev (2005)) and thus not directly comparable with the systems that we consider here., The Local Group is an order of magnitude less massive $\sim 1-2 \times 10^{12}\ \msol$; \citet{2005AJ....129..178K}) ) and thus not directly comparable with the systems that we consider here.
+" We can consider Leo T, a dwarf galaxy in the Local Group which still has a significant amount of (~4x10? Mo) and is at a Galactocentric distance of 420 kpc."," We can consider Leo T, a dwarf galaxy in the Local Group which still has a significant amount of $\sim 4 \times 10^5\ \msol$ ) and is at a Galactocentric distance of 420 kpc."
+ Leo T is likely to be similar to the progenitors of the newly discoveredMilky Way satellites (Willman all of which are at radii «250 kpc and have no 2010)detected., Leo T is likely to be similar to the progenitors of the newly discoveredMilky Way satellites \citep{2010AdAst2010E..21W} all of which are at radii $< 250$ kpc and have no detected.
+. Grcevich&Putman(2009) estimate that Leo T would have to experience a halo gas density of 0.6—2.3x107?cm? to entirely remove its current content., \citet{2009ApJ...696..385G} estimate that Leo T would have to experience a halo gas density of $0.6-2.3 \times 10^{-3} \cmc$ to entirely remove its current content.
+" We consider ourselves to be probing intergalactic gas using the radio sources in this paper, however, for systems like the Local Group the dividing line between halo gas in individual galaxies and IGM gas is unclear."," We consider ourselves to be probing intergalactic gas using the radio sources in this paper, however, for systems like the Local Group the dividing line between halo gas in individual galaxies and IGM gas is unclear."
+ That dwarf galaxies in the Local Group could experience densities in this range is not impossible considering the IGM densities that we report here and the crossing time (~2 Gyr) for the Milky Way system., That dwarf galaxies in the Local Group could experience densities in this range is not impossible considering the IGM densities that we report here and the crossing time $\sim 2$ Gyr) for the Milky Way system.
+ In order to produce dSphs the neutral gas must be removed and the distribution of stars must also change., In order to produce dSphs the neutral gas must be removed and the distribution of stars must also change.
+" Tidal stirring caused by repeated shocking at the pericenter of the dwarf's orbit can dissipate angular momentum, transforming the rotationally supported stellar distribution into one that is pressure supported (Mayeretal.2006;Mayer2011).."," Tidal stirring caused by repeated shocking at the pericenter of the dwarf's orbit can dissipate angular momentum, transforming the rotationally supported stellar distribution into one that is pressure supported \citep{2006MNRAS.369.1021M,2011arXiv1104.4278M}."
+" In these simulations, tidal heating increases the effectiveness of ram-pressure stripping by expanding the dwarf galaxy mass distribution."," In these simulations, tidal heating increases the effectiveness of ram-pressure stripping by expanding the dwarf galaxy mass distribution."
+ Ram pressure stripping is likely not strong enough to remove large quantities of neutral gas from galaxies with total masses of ~10°—1019Mo which is where the for groups of galaxies begins to flatten out., Ram pressure stripping is likely not strong enough to remove large quantities of neutral gas from galaxies with total masses of $\sim 10^9 - 10^{10}\ \msol$ which is where the for groups of galaxies begins to flatten out.
+" In this regime, it may be more likely that these larger galaxies experience strong tidal stripping through interactions with neighboring galaxies similar to the interaction between the Large and Small Magellanic Clouds and the Milky Way."," In this regime, it may be more likely that these larger galaxies experience strong tidal stripping through interactions with neighboring galaxies similar to the interaction between the Large and Small Magellanic Clouds and the Milky Way."
+ In the group environment the space velocities are similar to the internal rotational velocities for large galaxies which makes tidal interactions especially damaging., In the group environment the space velocities are similar to the internal rotational velocities for large galaxies which makes tidal interactions especially damaging.
+ observations of galaxy groups show numerous examples of intergalactic that appears to be tidally stripped or large galaxies which are deficient for their morphological type (Yunetal.1994;Montenegroetal.2001;Freeland 2009)..," observations of galaxy groups show numerous examples of intergalactic that appears to be tidally stripped or large galaxies which are deficient for their morphological type \citep{1994Natur.372..530Y,2001A&A...377..812V,2009AJ....138..295F}."
+ 'The stripping of hot gas from the halos of galaxies preventing this gas from providing a fuel source for further star-formation is referred to as strangulation or starvation., The stripping of hot gas from the halos of galaxies preventing this gas from providing a fuel source for further star-formation is referred to as strangulation or starvation.
+ There are a handful of observations of X-ray tails from galaxies in a group environment (Jeltemaetal.2008;Machacek2007;Ras- 2004)..," There are a handful of observations of X-ray tails from galaxies in a group environment \citep{2008ApJ...679.1162J,2007ApJ...664..804M,2006MNRAS.370..453R,2005ApJ...630..280M,2004ApJ...617..262S}. ."
+ Most simulations focus on ram-pressure, Most simulations focus on ram-pressure
+The data analysis was carried out with the Science Analysis System (SAS) version 7.0 (?)..,The data analysis was carried out with the Science Analysis System (SAS) version 7.0 \citep{sas}.
+ Standard selection criteria were applied to the data and periods with enhanced background due to proton flares were discarded., Standard selection criteria were applied to the data and periods with enhanced background due to proton flares were discarded.
+ To increase the signal-to-noise-ratio of the RGS data we further restricted the time intervals with acceptable background rates (CCD9 rate «00.5 ccts/s) for fainter sources like 8 CCom and HD 81809 and merged the individual observations for each targets using the tool ’rgscombine’., To increase the signal-to-noise-ratio of the RGS data we further restricted the time intervals with acceptable background rates (CCD9 rate $<$ cts/s) for fainter sources like $\beta$ Com and HD 81809 and merged the individual observations for each targets using the tool '.
+ The data was analyzed with standard tools of CIAO 3.4 (?) we applied standard data processing and filtering and we used positive and negative first order spectra to measure line fluxes., The data was analyzed with standard tools of CIAO 3.4 \citep{ciao} we applied standard data processing and filtering and we used positive and negative first order spectra to measure line fluxes.
+ The individual LETGS exposures provide sufficient data quality and are analyzed separately., The individual LETGS exposures provide sufficient data quality and are analyzed separately.
+" To investigate global stellar X-ray properties, we use a combination of RGS/MOS data from the observations, for LETGS we use spectra in the range."," To investigate global stellar X-ray properties, we use a combination of RGS/MOS data from the observations, for LETGS we use spectra in the range."
+ For the determination of the Ne/O ratios only high resolution X-ray spectra with good SNR are taken into account., For the determination of the Ne/O ratios only high resolution X-ray spectra with good SNR are taken into account.
+ Global spectral analysis was performed with XSPEC V11.3 (2) using multi-temperature models (three temperature components are sufficient to describe our spectra) with variable abundances as calculated with the APEC code (?).., Global spectral analysis was performed with XSPEC V11.3 \citep{xspec} using multi-temperature models (three temperature components are sufficient to describe our spectra) with variable abundances as calculated with the APEC code \citep{apec}.
+" In our models, temperatures, emission measure and abundances of Ne, O and Fe are treated as free parameters."," In our models, temperatures, emission measure and abundances of Ne, O and Fe are treated as free parameters."
+" Depending on spectral quality, the abundances of other elements were also free parameters or set to solar values."," Depending on spectral quality, the abundances of other elements were also free parameters or set to solar values."
+" We also checked the results by fitting only those spectral ranges dominated by Ne and O lines, i.e.,,, and additionally for the LETGS."," We also checked the results by fitting only those spectral ranges dominated by Ne and O lines, i.e., and additionally for the LETGS."
+" We find that while the absolute abundances of Ne and O differ somewhat between the applied models, mainly due to the interdependence with the emission measure, the derived Ne/O ratio turns out to be a quite stable result."," We find that while the absolute abundances of Ne and O differ somewhat between the applied models, mainly due to the interdependence with the emission measure, the derived Ne/O ratio turns out to be a quite stable result."
+" For line fitting purposes we use the CORA program (?),, using identical line width and assuming Lorentzian line shapes for all lines of a respective detector."," For line fitting purposes we use the CORA program \citep{cora}, using identical line width and assuming Lorentzian line shapes for all lines of a respective detector."
+" This analysis uses total spectra, i.e. we determine one ((+continuum) level in the respective region around each line."," This analysis uses total spectra, i.e. we determine one (+continuum) level in the respective region around each line."
+ To derive the Ne/O ratios from linear combinations of emission lines we especially measure the Lya line (18.97 A)) and the resonance line (21.6 A)) as well as the Lya line (12.13 À)) and the resonance line (13.45 A))., To derive the Ne/O ratios from linear combinations of emission lines we especially measure the ${\alpha}$ line (18.97 ) and the resonance line (21.6 ) as well as the ${\alpha}$ line (12.13 ) and the resonance line (13.45 ).
+ Note that both Lya lines are actually unresolved doublets and that the strength of the resonance lines of and were derived froma fit to the respective triplet with all three lines included in the fit., Note that both ${\alpha}$ lines are actually unresolved doublets and that the strength of the resonance lines of and were derived from a fit to the respective triplet with all three lines included in the fit.
+" As an example, we show in refspec the X-ray spectrum of 61CCyg from the RGS detectors with the spectral lines of interest labeled."," As an example, we show in \\ref{spec} the X-ray spectrum of Cyg from the RGS detectors with the spectral lines of interest labeled."
+" The spectrum consists of 11 co-added observations, background subtraction and flux conversion were done with the tool."," The spectrum consists of 11 co-added observations, background subtraction and flux conversion were done with the tool."
+" In the line ratio approach one uses linear combinations of these line fluxes to determine the Ne/O abundance ratio, however, the specific linear combination can be differently chosen."," In the line ratio approach one uses linear combinations of these line fluxes to determine the Ne/O abundance ratio, however, the specific linear combination can be differently chosen."
+" In refchianti we show the emissivity curves, i.e. G(T), of the respective linear combinations for Ne and O as used by ? ((D&TT) and ? ((L&SS) and calculated with the CHIANTI 55.2 code (??).."," In \\ref{chianti} we show the emissivity curves, i.e. G(T), of the respective linear combinations for Ne and O as used by \cite{dra05} T) and \cite{lie06} S) and calculated with the CHIANTI 5.2 code \citep{chi,anti}."
+" Unfortunately the original weighting is based on energy flux in D&TT(Ovm vs. Nerx+0.15x Nex) and on photon flux in L&SS (0.67xOvi—0.17 vs. NeIX+0.02x X), we therefore normalized the distributions for comparison."," Unfortunately the original weighting is based on energy flux in T vs. $\ion{Ne}{ix}+0.15\times\ion{Ne}{x}$ ) and on photon flux in S $0.67\times\ion{O}{viii}-0.17\times\ion{O}{vii}$ vs. $\ion{Ne}{ix}+0.02\times\ion{Ne}{x}$ ), we therefore normalized the distributions for comparison."
+" Note that the L&SS emissivity for O is unphysically negativeat temperatures below MMK, corresponding to an LLya/O vu(r) energy flux ratio of 0.3 or lower."," Note that the S emissivity for O is unphysically negativeat temperatures below MK, corresponding to an $\alpha$ (r) energy flux ratio of $\sim 0.3$ or lower."
+ As is apparent from, As is apparent from
+using the Systemic Console (Aleschiarietal.2009)FAwww.,using the Systemic Console \citep{Meschiari09}.
+oklo.ore..OUS The errors on each parameter are estimated. using the ))otstrap technique with 5000) scrambled realizations ofn the RV+ datasets., The errors on each parameter are estimated using the bootstrap technique with 5000 scrambled realizations of the RV datasets.
+" For cach planet. we list. best-fit1 iod (P). ecceutricitv (6). semu-amplitude (AU. time of periastron passage (νο). longitude of periceuter (a). nininmna mass (VEsin /) and semianuajor axis (7),"," For each planet, we list best-fit period $P$ ), eccentricity $e$ ), semi-amplitude $K$ ), time of periastron passage $T_{peri}$ ), longitude of pericenter $\varpi$ ), minimum mass $\mass\ \sin i$ ) and semi-major axis $a$ )."
+ Additionally. we report approximate estimates of the ransit probability calculated as part of the Moute-Carlo uodeliug. assunine a putative radius R=Rrp.," Additionally, we report approximate estimates of the transit probability calculated as part of the Monte-Carlo modeling, assuming a putative radius $\radius = \radius_{JUP}$."
+ WD 31253 is a V ole7.133 magnitude star of spectral class Fa., HD 31253 is a V = 7.133 magnitude star of spectral class F8.
+ Relative to Sun. UD 31253 is modestly unctal-rich (|Fe/TI]= 0.16).," Relative to the Sun, HD 31253 is modestly metal-rich ([Fe/H] = 0.16)."
+ Table 3. shows the complete set of 39 relative radial velocity observations for IID 31253., 		 Table \ref{tab:rvdata_HD31253} shows the complete set of 39 relative radial velocity observations for HD 31253.
+ The radial velocity coverage spans approximately 13 vears of RV monitoriug., The radial velocity coverage spans approximately 13 years of RV monitoring.
+ The mecian‘ internalNN unecrtainty: for our observations fitis Emsand the peak-to-peak velocity variatiouNEN is DOT ur," The median internal uncertainty for our observations is 1.59, and the peak-to-peak velocity variation is 36.37."
+slovThe :velocity scatterDOT around the average RV iiin our measurements iSwsnrenients S92ds 8.92u, The velocity scatter around the average RV in our measurements is 8.92.
+s The top paucl of Figure 1.1 shows the individual RV observations for IID 31253, The top panel of Figure \ref{fig:data_HD31253} shows the individual RV observations for HD 31253.
+ SeareleThe middle pancl shows the error-weighted Lomb. (LS) periodogram of the full RV dataset. (Cilliland&Balinuas1987), The middle panel shows the error-weighted Lomb-Scargle (LS) periodogram of the full RV dataset \citep{Gilliland87}.
+ The three horizoutal− lines iu this figure aud other conrparable plots represent. from top to bottom. the and analytic False Alarii Probability(FAP) levels. respectively.," The three horizontal lines in this figure and other comparable plots represent, from top to bottom, the, and analytic False Alarm Probability (FAP) levels, respectively."
+⋅ The↴ analytic: FAPs) are computed using⋅↜ a straightforwardsh≽⋅⊓ ard:approach.ach. where we estimate:ὲ the number frequencies by analyzing⋅⋅↜ a setWe of 1.000 . do⋅⋅ with⋅ the⋅ same↸⊳∪∐↸⊳↕∏↴∖↴↕↖↽↸∖↕↖↽↥⋅∏↕↸∖∪∏↑↴∖↴↑↸∖∐⋜∐⋅↥⋅∪↑⋜↧⊓∪∐↴∖↴↕∶↴∙∐⋜↧⊓∐⋅↸∖↴∖↴⋜↧↴∖↴⋜↧↸⊳⋜⋯↴∖↴↸∖ timestamps as theeaussian . ≻⇁⋅⋅⋅ (Press1002 et," The analytic FAPs are computed using a straightforward approach, where we estimate the number of independent frequencies by analyzing a set of 1,000 gaussian deviates with the same timestamps as the original dataset \citep{Press}."
+ al.of1992)..1the. highest peak. the: quoted FAP is estimated lsuiofJXBNaus32 n* ⋅ ↻↸∖↥⋅↕∪≼↧∪↕⋜⊔∪∏↑−⋅≻≼⇂⋜∏↽↴∖↴∙↕⊔⋯⊳∐↴∖↴∐∪↥⋅↑↸∖↥⋅↑∐⋜⋯↑↕∐∖⊔≱∢⊢≺↧ . ∏↴∖↴∐↕∶↴∙⋜↕⋯∪↥⋅↸∖↥⋅∪↴⋝∏↴∖↴↑⋀∖↕∪∐↑↸∖≼⋜∐⋅↕∪⋜∏∏∐⋅∪⋜↧↸⊳∐∙↖↖⇁↕∐↸⊳∐ : ∙∙↕↘↸∖↻↕↸∖∐⋜⋯↻↸∖∐∪≼↧∙↽∕∏∐∖↴∖↴↸∖⋯↕≓⋜⋯∏≻∐↑⋯∐∖∪↕↑∐↸∖∪↴⋝↴∖↴↸∖↥⋅↖⇁↸∖≼↧ ↸⊳∪∐↴∖↴↕↴∖↴↑↴∖↴∪↕∶↴∙↸∖∐↸∖↥⋅⋜↧↑∐↕∶↴∙↴∖↴↸∖↑↴∖↴∪↕↴∖↴↸⊳↥⋅⋜⋯∐⋝↕↸∖≼↧↥⋅↸∖⋜↧∐∑⋜↧⊓∪∐↴∖↴∪↕ the dataset and determining∙⋅ the maxima: periodogram: power for: each (e.g.Marewetal.onnr.2005).," For the highest peak, the quoted FAP is estimated using a more robust Monte Carlo approach, which consists of generating sets of scrambled realizations of the dataset and determining the maximum periodogram power for each \citep[e.g.][]{Marcy05}."
+". For all the datasetsateHR presented“OSC:, m this* paper.éX. we ὉanalyzeANUZO ο3«M10 scrambled datasets."," For all the datasets presented in this paper, we analyze $3\times 10^5$ scrambled datasets."
+ The computed FAP for the stroug Keplerian signal at DP= 160.32 davs iu the RV dataset indicates an estimated FAP =1ν10ut, The computed FAP for the strong Keplerian signal at $P = $ 460.32 days in the RV dataset indicates an estimated FAP $\approx 4 \times 10^{-5}$.
+ Finally. the lower panel of Figure 1 shows the spectral window.," Finally, the lower panel of Figure \ref{fig:data_HD31253} shows the spectral window."
+" A peak at frequency. f£; in the spectral window fiction can be associated with aliases ocemnring at |f,zxf|. where is a true periodicity of the input signal."," A peak at frequency, $f_s$ in the spectral window function can be associated with aliases occurring at $\left| f_p \pm f_s\right|$, where $f_p$ is a true periodicity of the input signal."
+" For more details.f, see Dawson&Fab-rvcky (2010)."," For more details, see \citet{DawsonFabrycky10}."
+. Peaks in the spectral window function are often associated with relatively immutable periodicitics in the observational cadence. such as those arising from ↑↕∐∖↴∖↴↕≺∐∖↥⋅↸∖⋜↧↕⋜⋯≼↧↴∖↴∪↕⋜∐⋅≼↧⋜↧⋅↖↽∙↑↕∐∖↕∏∐⋜∐⋅↴∖↴⋅↖↽∐∪≼∐↸⊳↕⊔∪∐↑∐⋜⋯≼↧ the solar vear.," Peaks in the spectral window function are often associated with relatively immutable periodicities in the observational cadence, such as those arising from the sidereal and solar day, the lunar synodic month and the solar year."
+ The strongest peak in the periodograim is wellfit bv a Kepleriau orbit of period. 165.51 days aud seiü-unplitule A—12.22|., The strongest peak in the periodogram is well-fit by a Keplerian orbit of period 465.54 days and semi-amplitude $K = 12.22$.
+ Together with the ⋜↧↴∖↴↴∖↴⋯⊔↸∖≼↧↴∖↴↑↸∖∐⋜∐⋅⊔⋜↧↴∖↴↴∖↴∪↕⋡↕⋅⊇∶≩↜∖∕↿⋅∙↑↕∐↴∖↴⋜⋯∏≻∐↑⋯∐∖↕∐∏≻↕↸∖↴∖↴ ⋜↧∐∏∐↕∐∐∐⊔⋯⋜↧↴∖∷∖↴∪↕⋟↜∖∕↿↴∖↴↕∐∣⋟∶∪∙⋅," Together with the assumed stellar mass of 1.23 $\msun$ , this amplitude implies a minimum mass of $\mass \sin i = 0.50 \mjup$."
+↱↗∪↜∖∕↿⋮∣∙⊺↕∐∖↴⋈∖↴∖↴↑≓∐↑∪↥⋅↴⋝⇈ ↕≯∪↥⋅↑∐↸∖↻↕⋜⊔∐∖⊓↴∖↴⊔∐≼∐⋅↖↽↸∖↸⊳∩∖∐⊓⋅↕↸⊳≺↙≈∩∙, The best-fit orbit for the planet is mildly eccentric $e \approx 0.34$ ).
+∶≩⊔∙⊺↕∐↴∖↴↕≓↻↕⋜⊔↸∖↑ achieves⋅ a⋅ reduced 47 =2S87.amwith au4 RAIS of↽⊓ 1.23 1.," This 1-planet fit achieves a reduced $\chi^2 = 8.87 $, with an RMS of 4.23."
+59of 31253expected jitter of ΠΟ 31253 (that: is. theana 36.37 The iionjitterrequire to bring the reduced 4? of the beil-nt fo LU)1 Dis 3..02," The expected jitter of HD 31253 (that is, the amount of jitter required to bring the reduced $\chi^2$ of the best-fit solution to 1.0) is 3.92."
+ The top panel of Figure 2. UTshows the phased↴↴ stellar⋅ reflexflox velocity οἱof Hb 51350030253 compared:care to the> ΗΝRY dataset., The top panel of Figure \ref{fig:bestfit_HD31253} shows the phased stellar reflex velocity of HD 31253 compared to the RV dataset.
+ The velocitymiddle pancl shows the residuals to the ↕≓↻↕⋜∐∐∖↑↴∖↴∪↕∏↑↕∪, The middle panel shows the residuals to the 1-planet solution.
+↕∙⊟∐⋜↧∐↖↽∙↑∐↸∖↴⋝∪↑↑∪⋯↻⋜⋯↸∖↕↴∖↴∐∪↖↖↽↴∖↴∐↸∖ periodogram. of. the residualsDl of. the best-fit. solution.," Finally, the bottom panel shows the periodogram of the residuals of the best-fit solution."
+". No- interestingS peaks are evideut. indicatiug ≽⊓⋅⋅that the ]present data setprovidesPC no strong support for additional plauets iu. theκ. system,"," No interesting peaks are evident, indicating that the present data setprovides no strong support for additional planets in the system."
+population and have been reobserved.,population and have been reobserved.
+ 46 binary. candidates with significant RV shifts have been found., $46$ binary candidates with significant RV shifts have been found.
+ From the analysis of individual SDSS spectra. 81 additional stars with RV shifts on short timescales have been found.," From the analysis of individual SDSS spectra, $81$ additional stars with RV shifts on short timescales have been found."
+ Targets for follow-up spectroscopy have been selected using numerical simulations based on the properties of the known sdB close binary population and theoretical predictions about the relative fraction of massive compact companions., Targets for follow-up spectroscopy have been selected using numerical simulations based on the properties of the known sdB close binary population and theoretical predictions about the relative fraction of massive compact companions.
+ 69 binaries with high RV shifts as well as significant RV shifts on short timescales have been selected as good candidates for massive compact companions., $69$ binaries with high RV shifts as well as significant RV shifts on short timescales have been selected as good candidates for massive compact companions.
+ Atmospheric parameters. spectroscopic distances ard population memberships have been determined.," Atmospheric parameters, spectroscopic distances and population memberships have been determined."
+ The multi-site follow-up campaign started in. 2009 and is being conducted with nedium resolution. spectrographs mounted at several ditferert telescopes of mostly 4-m-class., The multi-site follow-up campaign started in 2009 and is being conducted with medium resolution spectrographs mounted at several different telescopes of mostly 4-m-class.
+ First results are presented in Geter et al. (20110)., First results are presented in Geier et al. \cite{geier11}) ).
+we found only a small difference. which was much too small to explain the discrepaney shown in Fig.,"we found only a small difference, which was much too small to explain the discrepancy shown in Fig."
+ 4+., 4.
+ Also acceptable departures from the mixing length parameter adopted in the evolutionary models only have a minor effect on the dependence of the vi—v» difference on v»., Also acceptable departures from the mixing length parameter adopted in the evolutionary models only have a minor effect on the dependence of the $\nu_3-\nu_2$ difference on $\nu_2$.
+ The problem with interpreting the b» and b; sequences in terms of the first two overtones could have been noticed earlier with a closer look at the Petersen Diagrams shown by Soszyfisski et al. (, The problem with interpreting the $_2$ and $_3$ sequences in terms of the first two overtones could have been noticed earlier with a closer look at the Petersen Diagrams shown by Soszyńsski et al. (
+2004).,2004).
+ In such diagrams. period ratios of all significant peaks detected in individual objects are plotted against the longer periods.," In such diagrams, period ratios of all significant peaks detected in individual objects are plotted against the longer periods."
+ In Fig., In Fig.
+ 11 of that paper. we may see a prominent cloud of points corresponding to the b» and b; sequences near P;/Pi=0.7.," 11 of that paper, we may see a prominent cloud of points corresponding to the $_2$ and $_3$ sequences near $P_2/P_3\approx0.7$."
+ However. most of them are located above this value. whereas in the same period range. the models yield values that are always below 0.7.," However, most of them are located above this value, whereas in the same period range, the models yield values that are always below 0.7."
+ The mean value for the cloud of 0.73 is not reached until v»= IjHz., The mean value for the cloud of 0.73 is not reached until $\nu_2\approx1\mu$ Hz.
+ We focused on type b OSARGs because they are. as we believe. RGB objects. and we lack needed models of AGB stars.," We focused on type b OSARGs because they are, as we believe, RGB objects, and we lack needed models of AGB stars."
+ Nonetheless. it is still instructive to compare frequencies of type a OSARGs with the values calculated along the isochrones in the same range of the W; magnitudes.," Nonetheless, it is still instructive to compare frequencies of type a OSARGs with the values calculated along the isochrones in the same range of the $W_I$ magnitudes."
+ The essential difference between AGB and RGB stars concerns the deep interior. which has no effect on radial node frequencies.," The essential difference between AGB and RGB stars concerns the deep interior, which has no effect on radial mode frequencies."
+ The former objects have helium-exhausted cores. while the latter are either in the phase of core helium burning or still ahead of it.," The former objects have helium-exhausted cores, while the latter are either in the phase of core helium burning or still ahead of it."
+ In Fig., In Fig.
+ 5. we compare data or type a OSARGs with results of the the same calculations as were used for type b. Points in the upper panel seem to follow a single isochrone. which should be somewhat younger and/or have higher Z than that for Z=0.008.," 5, we compare data on type a OSARGs with results of the the same calculations as were used for type b. Points in the upper panel seem to follow a single isochrone, which should be somewhat younger and/or have higher $Z$ than that for $Z=0.008$."
+ This suggests that at least some of type a OSARGs located below TRGB are in fact the first ascent objects before helium ignition., This suggests that at least some of type a OSARGs located below TRGB are in fact the first ascent objects before helium ignition.
+ The data on frequency differences. shown in the lower panel reveal the same problem with interpretation in terms of the frequency difference between the first two radial overtones., The data on frequency differences shown in the lower panel reveal the same problem with interpretation in terms of the frequency difference between the first two radial overtones.
+ The discrepancy is only somewhat smaller but there is much less spread in the data., The discrepancy is only somewhat smaller but there is much less spread in the data.
+ So the problem is not less severe., So the problem is not less severe.
+ We regard the discrepancy revealed in Figs., We regard the discrepancy revealed in Figs.
+ 4 and 5 às serious., 4 and 5 as serious.
+ Ány other assignment of radial mode orders makes the discrepancy worse., Any other assignment of radial mode orders makes the discrepancy worse.
+ If no modification in. stellar. models can remove the problem with interpretation in. terms of the radial modes. the option is to consider non-radial modes trapped in the envelope.," If no modification in stellar models can remove the problem with interpretation in terms of the radial modes, the option is to consider non-radial modes trapped in the envelope."
+ Such modes may avoid large radiative losses in the interior (Dziembowski et al., Such modes may avoid large radiative losses in the interior (Dziembowski et al.
+ 2001. Dupret et al 2009).," 2001, Dupret et al 2009)."
+ Though it is difficult to explain why such nodes could be preferentially excited. we did consider this possibility.," Though it is difficult to explain why such modes could be preferentially excited, we did consider this possibility."
+ We limited ourselves to dipolar modes because they exhibit relatively small amplitude reduction caused by cancelation of, We limited ourselves to dipolar modes because they exhibit relatively small amplitude reduction caused by cancelation of
+ N-body (e.g.22??).. (e...273: (e.c.277): 7): (0...2): (e.g...2," $N$ \citep[e.g.,][]{Frenk_etal88,
+ Dubinski_Carlberg91,Jing_Suto02,Bett_etal07}. \citep[e.g.,][]{Hoekstra_etal04,
+ Mandelbaum_etal06}; \citep[e.g.,][]{Ibata_etal01, Helmi04, Law_etal09}; \citep[e.g.,][]{Sackett_Sparke90}; \citep[e.g.,][]{Olling_Merrifield00}; \citep[e.g.,][]{Kolokotronis_etal01,Buote_etal02}."
+?).. e.c.?7).. eal," \citep[see,
+e.g.,][]{Dubinski94,Kazantzidis_etal04a,Abadi_etal10,Tissera_etal10}
+ \citep[e.g.,][]{Debattista_etal08, Valluri_etal10}."
+axy. such as its mass. size. or the timescale and mode of its asseniblv.," galaxy, such as its mass, size, or the timescale and mode of its assembly."
+ Tere we explore these issues using a series of coutrolle nunierical simulations where a triaxia] halo is evolve under the influence of a central «ink galaxy., Here we explore these issues using a series of controlled numerical simulations where a triaxial halo is evolved under the influence of a central disk galaxy.
+ Our muunerical experiments are simular iu nature to those of 7 and ?.., Our numerical experiments are similar in nature to those of \citet{Dubinski94} and \citet{Debattista_etal08}.
+" Tlowever. we exteud this work in severa respects,"," However, we extend this work in several respects."
+ For example. we focus on he shape of the eravitational potential rather than fiit of isodeusity contours.," For example, we focus on the shape of the gravitational potential rather than that of isodensity contours."
+ The poteutialis less sensitive o the influence of substructures. which can iuduce substaiial but transicut changes in the local deusitv whilst affectiug little the overall potential.," The potential is less sensitive to the influence of substructures, which can induce substantial but transient changes in the local density whilst affecting little the overall potential."
+" The eravitational potential is also a quantity of more direct relevance aux applicabilitv to many observational studies of halo sha205,", The gravitational potential is also a quantity of more direct relevance and applicability to many observational studies of halo shapes.
+ Further. we explore systematically several aspects of the erowth of the galaxy that may 1i principle affect the shape of its surrounding halo.," Further, we explore systematically several aspects of the growth of the galaxy that may in principle affect the shape of its surrounding halo."
+ We consider not only the eravitational portance of the «isk. but also its orientation relative to the halo princiod axes. as well as the timescale aud mode of its as«4ublv.," We consider not only the gravitational importance of the disk, but also its orientation relative to the halo principal axes, as well as the timescale and mode of its assembly."
+ Lastly. we use halo parameters in agrecuent wih the results of cosinological V-ody simmlations aud galaxy parameters which are cousiseut with observed scaling laws aud span the wide rauge iu surface xiehtuess of observed galaxy disks.," Lastly, we use halo parameters in agreement with the results of cosmological $N$ -body simulations and galaxy parameters which are consistent with observed scaling laws and span the wide range in surface brightness of observed galaxy disks."
+ We investigate changes in the shapes of triaxial dark matter (DM) halos induced by the erowth of central disk ealaxies modeled as rigid potentials (s 7 for details)., We investigate changes in the shapes of triaxial dark matter (DM) halos induced by the growth of central disk galaxies modeled as rigid potentials (see \citealt{Villalobos_etal10} for details).
+ We measure the shapeofthe eravitationalpotentialbv approximating the ixopoteutial surfaces by ellipsoids aud, We measure the shapeofthe gravitationalpotentialby approximating the isopotential surfaces by ellipsoids and
+ otal.2009).. (Mg) Mg=>10 Mg>10 (cf.Fever1999).. (c.g..Woosley1993:Burrowsctal.—2007).. (¢.@.," \citep{Meynet03, Meynet05, Eldridge06, Limongi06, Georgy09}, $M_\mathrm{f}$ $M_\mathrm{f} > 10$ $M_\mathrm{f} > 10$ \citep[cf.][]{Fryer99}. \citep[e.g.,][]{Woosley93, Burrows07}, \citep[e.g.,][]{Podsiadlowski04, Guetta07}."
+.2001:Caretta&DellaValle2007).. (2.0NM.XAe25 M.) cau be made from 12...60 primary coniponents in close binary svstenis via the so-called Case A/D mass λίαν of them may end their life as bright SNe [b/c leaving neutron stars as relates. if their final masses are less than about 7 10.," $2.0~\mathrm{M_\odot}
+\la M_\mathrm{He} \la 25~\mathrm{M_\odot}$ ) can be made from $12...60$ primary components in close binary systems via the so-called Case A/B mass Many of them may end their life as bright SNe Ib/c leaving neutron stars as remnants, if their final masses are less than about 7 – 10."
+... Population studies indeed show that close binary stars can produce a sufficient nuniber of SNe Ibe to explain their observed rate. without the need of invoking suele star progenitors (c.¢..Podsiadlowski.Joss&TsuTout 2008).," Population studies indeed show that close binary stars can produce a sufficient number of SNe Ibc to explain their observed rate, without the need of invoking single star progenitors \citep[e.g.,][]{Podsiadlowski92, deDonder98, Eldridge08}."
+. Therefore. it is most Likely that the majority of typical SNe Tbe are produced iu binary svsteus.," Therefore, it is most likely that the majority of typical SNe Ibc are produced in binary systems."
+ The observational evidence for the connection between SNe Tbe aud long sgauuuarav bursts (CRBs) has particularly motivated may observational studies to better understand SNe Ibe since the last decade (seeWoosley&Bloom2006.fora review)...," The observational evidence for the connection between SNe Ibc and long gamma-ray bursts (GRBs) has particularly motivated many observational studies to better understand SNe Ibc since the last decade \citep[see][for a
+review]{WB06}."
+ Theoretical stellar 1nodels of SNe Tbc progenitors are thus highly required nowadavs., Theoretical stellar models of SNe Ibc progenitors are thus highly required nowadays.
+ The most comprehensive studies on the detailed characteristics ofSNeTbe progenitorsinbinary svstenis were conducted by Woosley.Langer& (hereafter. WIN95) using miass-losing," The most comprehensive studies on the detailed characteristics ofSNeIbc progenitorsinbinary systems were conducted by \citet{Woosley95} (hereafter, WLW95) using mass-losing"
+"This spiral galaxy appears distorted in theGALEX image, showing broad streaks that point south-east away from the cluster centre.","This spiral galaxy appears distorted in the image, showing broad streaks that point south-east away from the cluster centre."
+ Some faint knots are visible in the MegaCam images., Some faint knots are visible in the MegaCam images.
+ Miller et al. (, Miller et al. (
+"2009) note the radio emission is offset from the galaxy and suggestive of ram-pressure stripping, while Finoguenov et al. (","2009) note the radio emission is offset from the galaxy and suggestive of ram-pressure stripping, while Finoguenov et al. ("
+2004) found that the X-ray source is offset to the east of the galaxy.,2004) found that the X-ray source is offset to the east of the galaxy.
+ Gavazzi et al. (, Gavazzi et al. (
+2006) detected the galaxy in HI but indicate that it is strongly gas-deficient (Defy;= 0.76).,2006) detected the galaxy in HI but indicate that it is strongly gas-deficient $_{\rm HI}=0.76$ ).
+ The Ha data show an offset in the emission to the east of the galaxy centre., The $\alpha$ data show an offset in the emission to the east of the galaxy centre.
+" This appears to be an edge-on SO or spiral galaxy, which presents a very long trail in theGALEX imaging, extending aarcmin kkpc) north towards the cluster centre."," This appears to be an edge-on S0 or spiral galaxy, which presents a very long trail in the imaging, extending arcmin kpc) north towards the cluster centre."
+" In the MegaCam data, the trail is faintly visible and the inner morphology is clearly disturbed."," In the MegaCam data, the trail is faintly visible and the inner morphology is clearly disturbed."
+ The galaxy is an HI non-detection in Bravo-Alfaro et al. (, The galaxy is an HI non-detection in Bravo-Alfaro et al. (
+"2000), with an upper limit which implies that it is gas deficient (Ὀείητ> 0.7).","2000), with an upper limit which implies that it is gas deficient $_{\rm HI}>0.7$ )."
+" In the INT data, we do not detect any Ha emission, either from the galaxy itself or from the trail."," In the INT data, we do not detect any $\alpha$ emission, either from the galaxy itself or from the trail."
+" Close in projection are an elliptical (GMP 2670 at aarcmin or kkpc, with rradial velocity difference) and a spiral (GMP 2601 at aarcmin or kkpc, with rradial velocity difference); hence this could possibly be a case of tidal stripping by a neighbouring galaxy, rather interaction with the cluster itself."," Close in projection are an elliptical (GMP 2670 at arcmin or kpc, with radial velocity difference) and a spiral (GMP 2601 at arcmin or kpc, with radial velocity difference); hence this could possibly be a case of tidal stripping by a neighbouring galaxy, rather interaction with the cluster itself."
+" This irregular or spiral galaxy has a post-starburst spectrum in the disk region, with the burst age estimated at MMyr and ongoing star formation in the nucleus (Caldwell, Rose Dendy 1999)."," This irregular or spiral galaxy has a post-starburst spectrum in the disk region, with the burst age estimated at Myr and ongoing star formation in the nucleus (Caldwell, Rose Dendy 1999)."
+" In ourGALEX image, we observe narrow tail of length ~0.5 aarcmin kkpc), extending to the north-east."," In our image, we observe a narrow tail of length $\sim$ arcmin kpc), extending to the north-east."
+a Stripping in this galaxy was first discussed by Yagi et al. (, Stripping in this galaxy was first discussed by Yagi et al. (
+"2007), who reported a kkpc Ha tail, which is also seen in our INT Ha imaging.","2007), who reported a kpc $\alpha$ tail, which is also seen in our INT $\alpha$ imaging."
+" This remarkably narrow and straight feature is also clearly seen in the deep Adami u-band image, and is co-located with the UV trail."," This remarkably narrow and straight feature is also clearly seen in the deep Adami $u$ -band image, and is co-located with the UV trail."
+" The presence of continuum emission suggests that star formation is taking place in the stripped material, not merely ionization of a purely gaseous tail as proposed by Yagi et al."," The presence of continuum emission suggests that star formation is taking place in the stripped material, not merely ionization of a purely gaseous tail as proposed by Yagi et al."
+" The galaxy is seen close in projection to an elliptical (GMP 2897, at aarcmin or kkpc) further away is another early-type galaxy (GMP 2852, at aarcmin or kkpc), but both have large differences in radial velocity aand rrespectively) and are unlikely to be physically associated with GMP 2910."," The galaxy is seen close in projection to an elliptical (GMP 2897, at arcmin or kpc) further away is another early-type galaxy (GMP 2852, at arcmin or kpc), but both have large differences in radial velocity and respectively) and are unlikely to be physically associated with GMP 2910."
+50.,50.
+ This increase could have a very significant effect. on magnification bias in galaxy surveys., This increase could have a very significant effect on magnification bias in galaxy surveys.
+" The dependence of a, on the scale length and. surface densitv of the disk is shown in 55.", The dependence of $\sigma_{\mu}$ on the scale length and surface density of the disk is shown in 5.
+ As in the case of the multiple-inaging cross section. the more massive disks. produce. a more significant οσο.," As in the case of the multiple-imaging cross section, the more massive disks produce a more significant effect."
+ Again. if the central surface mass density is increased. then the relative significance. of disks with smaller inclinations increases.," Again, if the central surface mass density is increased, then the relative significance of disks with smaller inclinations increases."
+ Adding a central bulge does not have a significant elfect on the high-magnification cross section: for a spherical »ulge with a à Iw density profile. extending out. to one kiloparsec from the centre and containing a total of 10 per cent of the disk mass. we find that the high-magnilication Cross section is increased by no more than 5r50 per cent.," Adding a central bulge does not have a significant effect on the high-magnification cross section; for a spherical bulge with a $r^{1/4}$ -law density profile, extending out to one kiloparsec from the centre and containing a total of 10 per cent of the disk mass, we find that the high-magnification cross section is increased by no more than 50 per cent."
+ The ojected surface mass density profile for an elliptical halo is similar to that for an inclined disk., The projected surface mass density profile for an elliptical halo is similar to that for an inclined disk.
+ Phus. we expect that a non-spherical halo which is Hattened along the disk is likely o enhance the lensing elfect of an inclined disk.," Thus, we expect that a non-spherical halo which is flattened along the disk is likely to enhance the lensing effect of an inclined disk."
+ However. his effect is not expected to be significant.," However, this effect is not expected to be significant."
+ The high-magnification cross section does not depend stronely on the redshift of the lens., The high-magnification cross section does not depend strongly on the redshift of the lens.
+ This suggests that the ligh-maenification cross section in both. the disk plus SIS ido model and the pure SIS halo model scales with the angular cliaunater clistance in a similar way.," This suggests that the high-magnification cross section in both, the disk plus SIS halo model and the pure SIS halo model scales with the angular diamater distance in a similar way."
+ Phus. we do not ⋖⊾⇀∖↓≻⋯↛⊓⊤∕↙∫⊤DIs ∕↙↿∪∣⋡⋖⋅≱∖↿↓⋅∪⊔⋏∙≟↓∙∖⇁∠⇂∢⋅↓≻⋖⋅⊔∠⇂⋖⋅⊔↿∪⊔⇂↓↕∢⊾∖⇁⋜↧↓⋯⊾∪⇂↿↓↕∢⊾ ⋅ ∠⇂∢⊾⊔⊳∖⊀↓↿∙∖⇁↓≻⋜⊔⋅⋜⋯↓⋖⋅∩⊾↓⋅≤≥⊔∪↓⋅↿↓↕∢⊾≼⇍∪⊔↓∪↓∪⋏∙≟⊲⊔∼⋜↧↓≼∼∪⊔⊳∖⋯⊔⇂∕∖⋡," Thus, we do not expect $\bar{\sigma}_{\mu}/\sigma^{\mathrm{SIS}}_{\mu}$ to be strongly dependent on the value of the density parameter $\Omega_0$ or the comological constant $\lambda$."
+" We have studied the effects of lensing by an exponential disk inside a SIS on two important parameters: the multiple- cross section σιµ. and the high-magnilication section e,CÀ)."," We have studied the effects of lensing by an exponential disk inside a SIS on two important parameters; the multiple-imaging cross section $\sigma_{\rm m}$, and the high-magnification cross-section $\sigma_\mu(A)$."
+ The lens model we use has a density. profile similar to that of the Alilky Way., The lens model we use has a density profile similar to that of the Milky Way.
+" Both ay and e,C) are sensitive to the lens parameters. in. particular to the inclination angle and the disk-to-halo mass ratio."," Both $\sigma_{\rm m}$ and $\sigma_\mu(A)$ are sensitive to the lens parameters, in particular to the inclination angle and the disk-to-halo mass ratio."
+ The results indicate that disks. ancl especially almost edge-on disks. have a significant ellect on the image configurations produced. by individual lenses. which agrees with the results of previous work bv. Maller et al. (," The results indicate that disks, and especially almost edge-on disks, have a significant effect on the image configurations produced by individual lenses, which agrees with the results of previous work by Maller et al. ("
+1997) and Wang Turner (1997). but that the statistical elfect on σι averaged. over all disk inclinations is unlikely to be very important.,"1997) and Wang Turner (1997), but that the statistical effect on $\sigma_{\rm m}$ averaged over all disk inclinations is unlikely to be very important."
+" The more signilicante result is. that includingὃν an exponential disk at an inclination greater than about ddeg increases 7, al elο10 by an order of magnitude as compared with an SIS lens. and that the inclination-averaged cross section. a, is increased. by a similar factor."," The more significant result is that including an exponential disk at an inclination greater than about deg increases $\sigma_\mu$ at $A>10$ by an order of magnitude as compared with an SIS lens, and that the inclination-averaged cross section $\bar \sigma_\mu$ is increased by a similar factor."
+ llence. disks could have a large elfect on the. predicted abundances of strongly magnified galaxy lenses.," Hence, disks could have a large effect on the predicted abundances of strongly magnified galaxy lenses."
+ This implies that in surveys for which dust extinction is unimportant. such as those in the radio and submillimetre wavebancds. the number of detected. strongly lensed sources. will increase (Blain 1996. 1997). and that they will tend to be preferentially associated with edge-on spiral galaxies (Blain etal.," This implies that in surveys for which dust extinction is unimportant, such as those in the radio and submillimetre wavebands, the number of detected, strongly lensed sources will increase (Blain 1996, 1997), and that they will tend to be preferentially associated with edge-on spiral galaxies (Blain et al.,"
+ in preparation)., in preparation).
+ Vhe preliminary results of the CLASS survey suggest that this is indeed the case (Browne ct al., The preliminary results of the CLASS survey suggest that this is indeed the case (Browne et al.
+ 1997)., 1997).
+ A recent preprint by Dartelmann Loch (1905) discussed. the statistical significance of lensing bw spira galaxies using a technique similar to ours and included the ellect of dust extinction in the lensing galaxy., A recent preprint by Bartelmann Loeb (1998) discussed the statistical significance of lensing by spiral galaxies using a technique similar to ours and included the effect of dust extinction in the lensing galaxy.
+" Vhey did no discuss the elfects of dilferent scale lengths or disk masses on the high-magnification cross sections and so it is ilicult to compare our results in detail. but they appear to be in broac agreement. bearing in mind that they assume £34=0.3 anc Il,=50kms1Alpe and that their disk model deviates [rom our exponential profile."," They did not discuss the effects of different scale lengths or disk masses on the high-magnification cross sections and so it is difficult to compare our results in detail, but they appear to be in broad agreement, bearing in mind that they assume $\Omega_0=0.3$ and $H_0=50\,\mathrm{km\,s^{-1}\,Mpc^{-1}}$ and that their disk model deviates from our exponential profile."
+ We thank Aviveh Maller for useful and interesting conversations. one of which motivated this paper. ancl Priva Natarajan. Malcolm: Longair and an anonvmous referee for providing helpful comments.," We thank Ariyeh Maller for useful and interesting conversations, one of which motivated this paper, and Priya Natarajan, Malcolm Longair and an anonymous referee for providing helpful comments."
+ (GCALEX. 1)). Τα (Ctavazzietal.2006) (Corteseet al.2009) 2007.," $M_r=-16.6$ \citep[\textit{GALEX}, \ref{UV}) $\alpha$ \citep{Gavazzi06} \citep{Cortese07, Yoshida08, Sun06, Sun07}."
+CO7).. Usronbereeretal.2008:I&apferer2009," \citep{Yoshida08, Kronberger08, Kapferer09, Chung09} \citep[][C07]{Cortese07}. \citep{Kronberger08, Kapferer09}."
+ medium. high deusity clouds are not stripped CTounueseu&Bryan 2009).," medium, high density clouds are not stripped \citep{Tonnesen09}."
+. For au observed example. isolated stax forming clouds remain in the disk of NGC 1102. another Vireo spiral. despite stripping of the IT (Crowlet2005).," For an observed example, isolated star forming clouds remain in the disk of NGC 4402, another Virgo spiral, despite stripping of the H \citep{Crowl05}."
+. We sugecst that the tail of IC. 3118 represents iu situ molecular cloud aud. star formation iu the galaxvs urbuleut wake following RPS., We suggest that the tail of IC 3418 represents in situ molecular cloud and star formation in the galaxy's turbulent wake following RPS.
+ Iu Sections 2 aud ?? we discuss the observations of IC 3118 aud its star-forming ail., In Sections \ref{ic3418} and \ref{knots} we discuss the observations of IC 3418 and its star-forming tail.
+ Iu Section { we show that while RPS of the IT is expected. theory distavors removing IH» from the galaxy via cither stripping or tices.," In Section \ref{env} we show that while RPS of the H is expected, theory disfavors removing $\rm{H_2}$ from the galaxy via either stripping or tides."
+ In Section 5 we discuss the norphologics of IC 3118 and simular star forming tails: demonstrating that they favor in situ molecular cloud ormation., In Section \ref{morphology} we discuss the morphologies of IC 3418 and similar star forming tails; demonstrating that they favor in situ molecular cloud formation.
+ IC 3115 is begiuniug the trausition between the blue and red sequences CNUVproο το) aud is thus interesting in its own right., IC 3418 is beginning the transition between the blue and red sequences ${\rm NUV}-r=2.75$ ) and is thus interesting in its own right.
+ The stars in the tail αν eventually contribute to the intracluster light (CL). warranting a study of the frequency of these tails.," The stars in the tail may eventually contribute to the intracluster light (ICL), warranting a study of the frequency of these tails."
+ We DIuuit our study to the tails formation., We limit our study to the tail's formation.
+" observed IC 3118 in the FUV (1350-1750 Aj} and the NUV (1750-2750 Aj) between 2001 March aud 2006 May with effective exposure times of 2 ks and 15 ks. respectively,"," observed IC 3418 in the FUV (1350-1750 ) and the NUV (1750-2750 ) between 2004 March and 2006 May with effective exposure times of 2 ks and 15 ks, respectively."
+ Au 15005 Ta observation was taken on the ESO 3.6 τα in 2005 and is publicly available (Cavazzietal.2003. 2006).," An 1800 s $\alpha$ observation was taken on the ESO 3.6 m in 2005 and is publicly available \citep{Gavazzi03, Gavazzi06}."
+. The ealaxy was observed by the Sloan Digital Sky Survey (SDSSYorketal.2000)... but optical colors cannot be measured for the tail.," The galaxy was observed by the Sloan Digital Sky Survey \citep[SDSS][]{York00}, but optical colors cannot be measured for the tail."
+ We adopt a distance to Virgo of 16.5 Mpe (Meietal2007).. at which distance GALENs resolution is 100 pe (5).," We adopt a distance to Virgo of 16.5 Mpc \citep{Mei07}, at which distance 's resolution is 400 pc $5\arcsec$ )."
+ IC 3118 is aamember of the Virgo Cluster at a projected distance of 150 kpc (91) from the cluster's core with a recessional velocity of 38kuis (Ciavazzietal.2001)., IC 3418 is a member of the Virgo Cluster at a projected distance of 450 kpc$94\arcmin$ ) from the cluster's core with a recessional velocity of $38~\rm{km~s^{-1}}$ \citep{Gavazzi04}.
+ Virgo is a dynamically voung cluster. as evidenced bysubstructure in the ealaxy distribution and the N-rav cluission (Bineechetal.1987:Schindler 1999)..," Virgo is a dynamically young cluster, as evidenced bysubstructure in the galaxy distribution and the X-ray emission \citep{Bingelli87, Schindler99}. ."
+ Galaxies with IT tails which are presumably uudergomg RPSare observed iu Virgo’s outskirts, Galaxies with H tails which are presumably undergoing RPSare observed in Virgo's outskirts
+We search for archival HETEG observations lor AGNs. asChandra is the only instrument currently capable of resolving narrow iron Ixo lines.,"We search for archival HETG observations for AGNs, as is the only instrument currently capable of resolving narrow iron $\alpha$ lines."
+ We ouly choose the objects with narrow Fe Ka lines detected al a confidence level > 3 0., We only choose the objects with narrow Fe $\alpha$ lines detected at a confidence level $>$ 3 $\sigma$.
+ We also restrict our sample wilh BLR reverberation mapping data available., We also restrict our sample with BLR reverberation mapping data available.
+ The resulting sample consists of ten Sevfert 1] galaxies., The resulting sample consists of ten Seyfert 1 galaxies.
+ All the archival WETG observations of these AGNs are listed in Table 1., All the archival HETG observations of these AGNs are listed in Table 1.
+ The process of data reduction and spectrum fitting is described in [11]., The process of data reduction and spectrum fitting is described in [11].
+ To improve the measurement of line widths. for sources with more than one IIIEZTG exposure. we litted multiple spectra sinmultaneouslv with invariable line center enerey and FWIIM.," To improve the measurement of line widths, for sources with more than one HETG exposure, we fitted multiple spectra simultaneously with invariable line center energy and FWHM."
+ The fitting results are shown in Table 2., The fitting results are shown in Table 2.
+ From literature. we find infrared time lags for six of the sources (see table 2).," From literature, we find infrared time lags for six of the sources (see table 2)."
+ From Fig., From Fig.
+ 30 in [4]. we see a tight correlation between the infrared lag and the V. band luminosity. with a scatter of 0.2 dex.," 30 in [4], we see a tight correlation between the infrared lag and the $V$ band luminosity, with a scatter of 0.2 dex."
+ Such a relationship was adopted to estimate the inlrared lags for the remaining four sources without infrared reverberation mapping observations in our sample., Such a relationship was adopted to estimate the infrared lags for the remaining four sources without infrared reverberation mapping observations in our sample.
+ Peterson et al., Peterson et al.
+ presented black hole masses lor 35 AGNs based on broad emission-line reverberation mapping data [20]., presented black hole masses for 35 AGNs based on broad emission-line reverberation mapping data [20].
+ The black hole masses are derived as where /5 (the scaling factor for BL reverberation mapping) is a zero-point calibration determined using the Mj;—σι relationship [19.24.25].," The black hole masses are derived as where $f_B$ (the scaling factor for BLR reverberation mapping) is a zero-point calibration determined using the $M_{BH}-\sigma_{\ast}$ relationship [19,24,25]."
+ since (he narrow Fe Ίνα line is always modeled by a single Gaussian component. in this research we caleulate the virial products as where fy is the scaling laetor. 7 the infrared lag and Veyyay the FWIIAI of the narrow iron Ίνα line.," Since the narrow Fe $\alpha$ line is always modeled by a single Gaussian component, in this research we calculate the virial products as where $f_T$ is the scaling factor, $\tau$ the infrared lag and $V_{FWHM}$ the FWHM of the narrow iron $\alpha$ line."
+" For a Gaussian. FWIIM/9j;;,,=2.355."," For a Gaussian, $\sigma_{line}=2.355$."
+" Note although the Fe Ίνα line consists of two components (xa, and Was). with the spectral resolution of the ΤΗΤα. its impact on the measurement of the line width with a single Gaussian is negligible [26]."," Note although the Fe $\alpha$ line consists of two components $\alpha_1$ and $\alpha_2$ ), with the spectral resolution of the HETG, its impact on the measurement of the line width with a single Gaussian is negligible [26]."
+ As shown in Figure 1. nine of the derived Ας are consistent with A/ps. but a scaling factor is requirecl.," As shown in Figure 1, nine of the derived $M_{T}$ s are consistent with $M_{B}$ s, but a scaling factor is required."
+ The solid line is (he best-fit line (slope fixed to 1) for all ten sources. derived using the orthogonal regression program GAUSSFIT (version 3.55: [27]).," The solid line is the best-fit line (slope fixed to 1) for all ten sources, derived using the orthogonal regression program GAUSSFIT (version 3.55; [27])."
+ The asymmetric statistical errors in (he masses were svinnmnetrized as (he mean of the positive ancl negative errors since GAUSSFIT can not work with asymmetric errors., The asymmetric statistical errors in the masses were symmetrized as the mean of the positive and negative errors since GAUSSFIT can not work with asymmetric errors.
+ The best-fit scaling factor. fy , The best-fit scaling factor $f_T$ 
+density defines the so-called opacity limit for fragmentation. which sets a minimum fragment mass of a few Jupiter masses (Low Lynden-Dell 1976: Boss 1988: BBB).,"density defines the so-called opacity limit for fragmentation, which sets a minimum fragment mass of a few Jupiter masses (Low Lynden-Bell 1976; Boss 1988; BBB)."
+ To model the opacity limit. for fragmentation without performing full radiative transfer. we use an equation of state given by p — typ”. where p is the pressure ancl A is a measure of the entropy of the gas.," To model the opacity limit for fragmentation without performing full radiative transfer, we use an equation of state given by $p$ = $K$ $\rho^{\eta}$, where $p$ is the pressure and $K$ is a measure of the entropy of the gas."
+ The value of 7 changes with density as: The gas is assumed to consist. of pure molecular hvdrogen (jj = 2). and the value of A. is such that. when the gas is isothermal. ἐν = cz. with the sound speed e=10tems tat T=10 WK. The pressure is continuous when the value of η changes.," The value of $\eta$ changes with density as: The gas is assumed to consist of pure molecular hydrogen $\mu$ = 2), and the value of $K$ is such that when the gas is isothermal, $K$ = $c_{\rm s}^2$, with the sound speed $c_{\rm s} = 1.85 \times 10^4 {\rm
+cm s}^{-1}$ at $T = 10$ K. The pressure is continuous when the value of $\eta$ changes."
+" A sink particle is inserted when the central density of a fragment exceeds p,=10MoIcm37. well above the critical""m density pi."," A sink particle is inserted when the central density of a fragment exceeds $\rho_{\rm s} = 10^{-10} {\rm g~cm}^{-3}$, well above the critical density $\rho_{\rm c}$."
+ Sink particles are point masses with an accretion radius. so that any gas particle that falls into it and is bound to the point mass is accreted.," Sink particles are point masses with an accretion radius, so that any gas particle that falls into it and is bound to the point mass is accreted."
+ In the present calculations. the aceretion radius f is Constant ancl equal to 5 AU.," In the present calculations, the accretion radius $R_{\rm sink}$ is constant and equal to 5 AU."
+ Therefore. disces around sink particles will be resolved only if their radii z10 AU.," Therefore, discs around sink particles will be resolved only if their radii $\gsim 10$ AU."
+ Sink particles interact with the gas only via gravity and accretion., Sink particles interact with the gas only via gravity and accretion.
+ The gravitational. acceleration between two sink particles is Newtonian for r.zc4 AU. but. is smoothed within this raclius using spline softening (Benz 1990).," The gravitational acceleration between two sink particles is Newtonian for $r \geq 4$ AU, but is smoothed within this radius using spline softening (Benz 1990)."
+ The maximum acceleration occurs at à~1 AU: therefore. this is the minimum binary separation that can be resolved.," The maximum acceleration occurs at $r \sim 1$ AU; therefore, this is the minimum binary separation that can be resolved."
+ We have performed. 19 cilferent calculations. of the fragmentation of a small-scale. turbulent molecular cloud. each of them under almost οκασ the same initial conditions.," We have performed 10 different calculations of the fragmentation of a small-scale, turbulent molecular cloud, each of them under almost exactly the same initial conditions."
+ Each cloud: core is spherical. has a mass of 5 AL. radius of z10! AU. and an initial uniform density of 10Seem7.," Each cloud core is spherical, has a mass of 5 $_\odot$, a radius of $\approx 10^4$ AU, and an initial uniform density of $10^{-18} {\rm g~cm}^{-3}$."
+ At the initial temperature of 10 Ix. the mean thermal Jeans mass is 0.5 M... Le. the Jeans number of the cloud is LO.," At the initial temperature of 10 K, the mean thermal Jeans mass is 0.5 $_\odot$, i.e. the Jeans number of the cloud is 10."
+ Phe global free-fall time of the cloud ἐς is z107 Ve., The global free-fall time of the cloud $t_{\rm ff}$ is $\approx 10^{5}$ yr.
+ We have imposed an initial supersonic turbulent velocity Ποιά on the gas. in the same manner as Ostriker. Stone Gammic (2001) and BBB.," We have imposed an initial supersonic turbulent velocity field on the gas, in the same manner as Ostriker, Stone Gammie (2001) and BBB."
+" We generate a divergence-free random Gaussian velocity field with a power spectrum PCA)xA"". where & is the wavenumber ancl à is the power index.half."," We generate a divergence-free random Gaussian velocity field with a power spectrum $P(k) \propto k^{\alpha}$, where $k$ is the wavenumber and $\alpha$ is the power index,."
+ Phe velocity [old is normalised so that initially it is in equipartition with the eravitational potential energy of the cloud core., The velocity field is normalised so that initially it is in equipartition with the gravitational potential energy of the cloud core.
+ These simulations do not include magnetic fields. as we have tried to isolate a particular. hyelrodvnanical fragmentation problem to characterise the properties of the resulting stellar systems.," These simulations do not include magnetic fields, as we have tried to isolate a particular hydrodynamical fragmentation problem to characterise the properties of the resulting stellar systems."
+ We have not included in our mocels any mechanical or raciative feedback mechanism., We have not included in our models any mechanical or radiative feedback mechanism.
+ This may be an appropriate choice. since the maximum stellar mass in these simulations does not exceed | ALL. whereas the most powerful winds ancl photoionisation fronts in. star-forming regions are produced by much more massive stars.," This may be an appropriate choice, since the maximum stellar mass in these simulations does not exceed 1 $_\odot$, whereas the most powerful winds and photoionisation fronts in star-forming regions are produced by much more massive stars."
+ The local Jeans mass must be resolved. throughout the calculation (Bate Burkert LOOT: Truclove et al., The local Jeans mass must be resolved throughout the calculation (Bate Burkert 1997; Truelove et al.
+ LOOT: Whitworth 1998). otherwise some of the fragmentation might be artificially enhanced. or suppressed.," 1997; Whitworth 1998), otherwise some of the fragmentation might be artificially enhanced or suppressed."
+" In. order to model a 5 M. cloud core with critical density. pe we need to use 3.510"" particles.", In order to model a 5 $_\odot$ cloud core with critical density $\rho_{\rm c}$ we need to use $3.5 \times 10^5$ particles.
+ Each of the bhwdrodvnamic calculations that are discussed. in this paper required z4000 CPU hours on the SCL Origin. 3800. Computer of the United. Ixingdom. Astrophysical Fluics Facility (CAPE)., Each of the hydrodynamic calculations that are discussed in this paper required $\approx 4000$ CPU hours on the SGI Origin 3800 Computer of the United Kingdom Astrophysical Fluids Facility (UKAFF).
+ Vhe hydrodynamic evolution of cach cloud is followed. until of the initial gas particles are acereted., The hydrodynamic evolution of each cloud is followed until of the initial gas particles are accreted.
+ At this point. the remaining eas is removed and thereafter the system. is evolved. as an N-body ensemble for LO Myr. using the numerical codeNBODYI. by Aarseth (Aarseth 1963: see Aarseth 1999 for an updated: description. of the code).," At this point, the remaining gas is removed and thereafter the system is evolved as an $N$ -body ensemble for 10 Myr, using the numerical code, by Aarseth (Aarseth 1963; see Aarseth 1999 for an updated description of the code)."
+ is a simple N-body algorithm which computes the gravitational forces between particles by direct summation., is a simple $N$ -body algorithm which computes the gravitational forces between particles by direct summation.
+ The numerical scheme is based upon the expansion of the gravitational force in a fourth-order polynomial with cliviced differences., The numerical scheme is based upon the expansion of the gravitational force in a fourth-order polynomial with divided differences.
+ Individual timesteps are also implemented., Individual timesteps are also implemented.
+ The ewtonian potential is softened in a manner that mimics a ΠΙΟ sphere., The Newtonian potential is softened in a manner that mimics a Plummer sphere.
+ lacks any regularization scheme [or the reatment of close binaries. in contrast with subsequent more sophisticated codes.," lacks any regularization scheme for the treatment of close binaries, in contrast with subsequent more sophisticated codes."
+ Close binaries can nevertheless ος accurately integrated. provided that a sullicicnthy small softening length and 5 parameter (which controls the size of the timesteps) is chosen., Close binaries can nevertheless be accurately integrated provided that a sufficiently small softening length and $\eta$ parameter (which controls the size of the timesteps) is chosen.
+ We took softening lengths. of Ll. the softening length of the SPIEL simulations. ancl 1 xwameters smaller than LO7. and found that energy. and angular momentum were conserved to an οσοαν) better han one part in one million.," We took softening lengths of $0.1 \times$ the softening length of the SPH simulations, and $\eta$ parameters smaller than $10^{-2}$, and found that energy and angular momentum were conserved to an accuracy better than one part in one million."
+ Each N-body simulation required only =50 CPU hours to be completed., Each $N$ -body simulation required only $\approx 50$ CPU hours to be completed.
+ The hyvdrodynamical evolution of the eloud produces shocks which decrease the turbulent kinetic energy. that initially supported the eloud., The hydrodynamical evolution of the cloud produces shocks which decrease the turbulent kinetic energy that initially supported the cloud.
+ In parts of the system. gravity begins to dominate and dense sell-geavitating cores form and collapse.," In parts of the system, gravity begins to dominate and dense self-gravitating cores form and collapse."
+ ‘These dense cores are the sites where the formation of stars and brown cdwarfs occurs., These dense cores are the sites where the formation of stars and brown dwarfs occurs.
+ The turbulence decays on the dynamical timescale of the cloud (as found by Mac Low οἱ al., The turbulence decays on the dynamical timescale of the cloud (as found by Mac Low et al.
+ 1998: Stone. OstrikerCGammic 1998: ancl BBB. among others). and star formation begins just after 1 to 1.5 global free-fall times Zip.," 1998; Stone, OstrikerGammie 1998; and BBB, among others), and star formation begins just after 1 to 1.5 global free-fall times $t_{\rm ff}$."
+ As mentioned before. the hivdrodynamical calculations were stopped. when of the gas has been accreted.," As mentioned before, the hydrodynamical calculations were stopped when of the gas has been accreted."
+ In terms of fy this means that. for the a = 3 and a = 5 ealeulations (heneclorth a3 and a5). we follow the evolution of the eloud for z-Ε ancl 5.5/7. respectively (i.c. an average of zz0.5 Myr).," In terms of $t_{\rm ff}$ this means that, for the $\alpha$ = $-3$ and $\alpha$ = $-5$ calculations (henceforth $\alpha3$ and $\alpha5$ ), we follow the evolution of the cloud for $\approx 4 t_{\rm ff}$ and $5.5
+t_{\rm ff}$ , respectively (i.e. an average of $\approx 0.5$ Myr)."
+ Altogether. the calculations," Altogether, the calculations"
+the disk is not expected to be altered sienilicantly: gas would therefore fill the region withinHye.,the disk is not expected to be altered significantly; gas would therefore fill the region within$\Rhole$.
+" Lea giant planet has formed with a mass sufficient to open a gap. both the gas and dust would be expected to be cleared dvnamically from the vicinity of the orbit of the planet. creating an inner disk (within A,<pope) Chat is fed by accretion streams [rom an outer disk (bevond Ryo: Lubow. Seibert. Artvmowicz 1999: Bryden οἱ 11999: lev 1999: D'Angelo et 22003: Lubow D'Angelo 2006)."," If a giant planet has formed with a mass sufficient to open a gap, both the gas and dust would be expected to be cleared dynamically from the vicinity of the orbit of the planet, creating an inner disk (within $\Rinner < \Rhole$ ) that is fed by accretion streams from an outer disk (beyond $\Rhole$; Lubow, Seibert, Artymowicz 1999; Bryden et 1999; Kley 1999; D'Angelo et 2003; Lubow D'Angelo 2006)."
+ Accretion streams are not expected in the case of a massive giant. planet (~10: e.g.. Lubow et 11999). and no significant gas or dust is expected anvwhere within Ay in this case.," Accretion streams are not expected in the case of a massive giant planet $\sim 10 M_J$; e.g., Lubow et 1999), and no significant gas or dust is expected anywhere within $\Rhole$ in this case."
+ A lack οἱ gas or dust within yee is also expected in the photoevaporation case: such svstenis are further expected to have a very low disk mass (e0.00137.: e.g.. Alexander Armitage 2007).," A lack of gas or dust within $\Rhole$ is also expected in the photoevaporation case; such systems are further expected to have a very low disk mass $\sim 0.001 \Msun$; e.g., Alexander Armitage 2007)."
+ Our results for V836 Tau are intriguing in this context., Our results for V836 Tau are intriguing in this context.
+ Compared to the SEDs of transition objects such as GAL Aur ancl DM Tan. where a strong infrared. excess only bevond ~I0jmn indicating an optically thin region AAU in size (Calvet et wpe22005). the SED of V336 Tan shows a significant infrared excess at a shorter wavelength ~5—Ίθμαι.," Compared to the SEDs of well-studied transition objects such as GM Aur and DM Tau, where a strong infrared excess appears only beyond $\sim 10\micron$ indicating an optically thin region AU in size (Calvet et 2005), the SED of V836 Tau shows a significant infrared excess at a shorter wavelength $\sim 5-10\micron$."
+" Therefore. if the V836 Tau svstem has an optically thin inner region. il is comparatively small ancl plausibly within ""range of disk radii probed by CO fundamental emission (within S 2AAU: opticallyNajita et ""m3)."," Therefore, if the V836 Tau system has an optically thin inner region, it is comparatively small and plausibly within the range of disk radii probed by CO fundamental emission (within $\lesssim 2$ AU; Najita et 2003)."
+ The SED of V836 Tau can be fit with a simple model of a flarecl (hick wilh an inner hole ~1 XAU in radius (see Appendix)., The SED of V836 Tau can be fit with a simple model of a flared optically thick disk with an inner hole $\sim 1$ AU in radius (see Appendix).
+ In comparison. the SED of the transition object 115 can be interpreted as indicating an optically thin region within ~ 3AAU (Berein et 22004). or a radial gap extending from AAU (Espaillat et 22007).," In comparison, the SED of the transition object 15 can be interpreted as indicating an optically thin region within $\sim 3$ AU (Bergin et 2004), or a radial gap extending from AU (Espaillat et 2007)."
+ A planet orbiting at such laree distances c ΛΑΔΙ; or e40 AAU) may create a gap in the disk. but not racially truncate the gaseous disk significantly within AAU.," A planet orbiting at such large distances $\sim 3$ AU or $\sim 40$ AU) may create a gap in the disk, but not radially truncate the gaseous disk significantly within AU."
+ In contrast. svstems like V836 Tau. in which the SED max indicate an optically thin region at much smaller radii (<1 AAU). ave the ones for which the CO hiundamental emission wouldin principle be capable of diagnosing a physically truncated inner disk if an orbiting companion is present.," In contrast, systems like V836 Tau, in which the SED may indicate an optically thin region at much smaller radii $< 1$ AU), are the ones for which the CO fundamental emission would in principle be capable of diagnosing a physically truncated inner disk if an orbiting companion is present."
+ The radially iruncated CO emission that we observe nav support this picture., The radially truncated CO emission that we observe may support this picture.
+ As an alternative interpretation of the available data. the short wavelength SED of V836 Tau (A< I0jan) can also be fit with a simple model of an inclined (7= 60). geometrically flat disk (see Appendix) that possibly results from significant grain growth and settling.," As an alternative interpretation of the available data, the short wavelength SED of V836 Tau $\lambda < 10\micron$ ) can also be fit with a simple model of an inclined $i=60$ ), geometrically flat disk (see Appendix) that possibly results from significant grain growth and settling."
+ In such a situation. the gaseous disk would be racially continuous and we might expect to observe a CO prolile like those of other non-transition T Tauri stars.," In such a situation, the gaseous disk would be radially continuous and we might expect to observe a CO profile like those of other non-transition T Tauri stars."
+ Such a profile. twpically centrally peaked. is not observed.," Such a profile, typically centrally peaked, is not observed."
+ In the photoevaporation and massive giant planet scenarios for the origi of a transitional SED. little gas is expected to be present within ο. in contrast to the observed situation where a gaseous disk is present at AAU and ongoing (possibly intermittent) stellar," In the photoevaporation and massive giant planet scenarios for the origin of a transitional SED, little gas is expected to be present within $\Rhole$ , in contrast to the observed situation where a gaseous disk is present at AU and ongoing (possibly intermittent) stellar"
+a strongly gravitationally focused encounter with a binary of total mass A/ and semimajor axis a the cross section is X—za(26M/67) and the timescale of interaction is 7=l1/(nXo). which is much less than a vear for (vpical masses. velocities. ancl densities.,"a strongly gravitationally focused encounter with a binary of total mass $M$ and semimajor axis $a$ the cross section is $\Sigma=\pi a (2GM/\sigma^2)$ and the timescale of interaction is $\tau=1/(n\Sigma\sigma)$, which is much less than a year for typical masses, velocities, and densities."
+ llowever. stars (hat interact with the IMDBII-SMDBII binary are eventually thrown out of the svstem. so the bottleneck is the (time needed lor other stars to diffuse into (he required orbital phase space.," However, stars that interact with the IMBH-SMBH binary are eventually thrown out of the system, so the bottleneck is the time needed for other stars to diffuse into the required orbital phase space."
+" This ""loss cone"" of stars could cause supermassive black hole binaries {ο stall in their inspiral. before (μον get close enough lor gravitational radiation to be significant (Begelman. Blandford. Rees 1980: see Sigurdsson Rees 1997. Milosavljevie Merritt 2003. Sigurdsson 2003. ancl Makino Funato 2004 [or recent discussions)."," This “loss cone"" of stars could cause supermassive black hole binaries to stall in their inspiral, before they get close enough for gravitational radiation to be significant (Begelman, Blandford, Rees 1980; see Sigurdsson Rees 1997, Milosavljevic Merritt 2003, Sigurdsson 2003, and Makino Funato 2004 for recent discussions)."
+ For IAIBIIs. this is not likely to be a problem.," For IMBHs, this is not likely to be a problem."
+ As discussed by Yu Tremaine (2003). once the original contingent of stars is ejected [rom the loss cone. the svstem will settle into à state in which the rate of diffusion of stars into the loss cone is balanced by the rate at which thev are ejected by interaction with the IMDBII-SMDII binary.," As discussed by Yu Tremaine (2003), once the original contingent of stars is ejected from the loss cone, the system will settle into a state in which the rate of diffusion of stars into the loss cone is balanced by the rate at which they are ejected by interaction with the IMBH-SMBH binary."
+" From equation (38) of Yu Tremaine (2003). the hardening timescale for a black hole binary of total mass M is where m, is the typical stellar mass in the central regions and 2x10|vr tis a characteristic value for the diffusion rate ray into the loss cone."," From equation (38) of Yu Tremaine (2003), the hardening timescale for a black hole binary of total mass $M$ is where $m_*$ is the typical stellar mass in the central regions and $2\times 10^{-4}~{\rm yr}^{-1}$ is a characteristic value for the diffusion rate $n_{\rm diff}$ into the loss cone."
+ Therefore. depending on the details of the stellar distribution. the orbital radius of the IMDII could be reduced by several e-Ioldings in a IIubble time. especially if the SMDII has M.zi105M...," Therefore, depending on the details of the stellar distribution, the orbital radius of the IMBH could be reduced by several e-foldings in a Hubble time, especially if the SMBH has $M\lta 10^6\,M_\odot$."
+ This process could be enhanced slightly because stus that interact with the IMDBII will typically not be ejected entirely [rom the core. hence they will retium for several interactions (Milosavljevie Merritt 2003).," This process could be enhanced slightly because stars that interact with the IMBH will typically not be ejected entirely from the core, hence they will return for several interactions (Milosavljevic Merritt 2003)."
+ In addition. gas dvnaimical friction from molecular cloucls (see. e.g... Ostriker 1999) can shrink the orbit further. ," In addition, gas dynamical friction from molecular clouds (see, e.g., Ostriker 1999) can shrink the orbit further. —"
+By the time a<10* pe. gravitational radiation can be important for an IMDBIL-SMDII binary.,"By the time $a<10^{-3}$ pc, gravitational radiation can be important for an IMBH-SMBH binary."
+ The timescale to merger is (from. e.g.. Peters 1964) where e is the orbital eccentricity.," The timescale to merger is (from, e.g., Peters 1964) where $e$ is the orbital eccentricity."
+ Thus. ide«0.0003 pe or the eccentricityv is high. merger can happen within a Hubble time.," Thus, if $a<0.0003$ pc or the eccentricity is high, merger can happen within a Hubble time."
+ Therefore. in contrast what might be the case for (wo supermassive black holes in a binary (Degelman οἱ al.," Therefore, in contrast what might be the case for two supermassive black holes in a binary (Begelman et al."
+ 1950). it is unlikely that there is a hang-up problem for an binary.," 1980), it is unlikely that there is a hang-up problem for an IMBH-SMBH binary."
+ The difference is that an IMDBIL-SMDII binary at a given separation has a much smaller binding energv than a binary with (wo supermassive black holes., The difference is that an IMBH-SMBH binary at a given separation has a much smaller binding energy than a binary with two supermassive black holes.
+ Hence. the stars that are ejected or displaced in (he process of hardening the binary come from a," Hence, the stars that are ejected or displaced in the process of hardening the binary come from a"
+he other in which the magnetic axis is the z/-axis (primed co-ordinates).,the other in which the magnetic axis is the $z^\prime$ -axis (primed co-ordinates).
+ The naenetic dipole lies iu the y+ plane at an angle à to the rotation axis., The magnetic dipole lies in the $y-z$ plane at an angle $\alpha$ to the rotation axis.
+ An cussion point will have co-ordinates ΕΛ where r ds the emission altitude.," An emission point will have co-ordinates $(r, \theta_e, \varphi_e)$ where $r$ is the emission altitude."
+ The radiation is supposed to be emitted in a direction (μι) tangeut o the magnetic field lines.," The radiation is supposed to be emitted in a direction $(\theta_r, \varphi_r)$ tangent to the magnetic field lines."
+" The latter is obtained by appropriately correcting (νο). for the tangential direction at auv given point (the usual 5 factor or unall 6,3."," The latter is obtained by appropriately correcting $\theta_e, \varphi_e$ ) for the tangential direction at any given point (the usual $\frac{3}{2}$ factor for small $\theta_e$ )."
+ The polar cap boundary is defined by the locus of those angles (0.5. 5) Which describe the feet of last open field lines.," The polar cap boundary is defined by the locus of those angles $\theta_{eb}, \varphi_{eb}$ ) which describe the feet of last open field lines."
+" Therefore. the emissiou can emauate from the emission capwhose boundary is defined by the corresponding angles (0,5. 42,5)."," Therefore, the emission can emanate from the emission cap whose boundary is defined by the corresponding angles $\theta_{rb},\varphi_{rb}$ )."
+ Às shown in KS the effects due to the stellar gravitational field namely change in the dipole geometry and lisht bending are μιαν opposite and leave the Coldreich-Julian (1969) type of beam essentially unaltered., As shown in KS the effects due to the stellar gravitational field namely change in the dipole geometry and light bending are mutually opposite and leave the Goldreich-Julian (1969) type of beam essentially unaltered.
+ These effects rapidly become insignificant for altitudes beyond a few stellar radii., These effects rapidly become insignificant for altitudes beyond a few stellar radii.
+ More iuportautlv they affect the enüssion cap in a svnunetrical wav through an overall shrinking., More importantly they affect the emission cap in a symmetrical way through an overall shrinking.
+ We therefore ignore these ecucral relativistic effects here., We therefore ignore these general relativistic effects here.
+ The directions (0.47) after aberration are denoted bv angles (0.47).," The directions $\theta, \varphi$ ) after aberration are denoted by angles $\hat \theta, \hat \varphi$ )."
+" The aberrated values (0,..47,) of the angles (0,.47,) are given by (INS - Eqs."," The aberrated values $(\hat \theta_r, \hat \varphi_r )$ of the angles $(\theta_r, \varphi_r)$ are given by (KS - Eqs."
+ 37 aud 38). aud where +—(1(02)H8? is the Lorentz factor aud with is the corotation velocity of the emission point at altitude r aud rey is the light cvliuder radius.," 37 and 38), and where $\gamma= (1-v^2)^{-1/2}$ is the Lorentz factor and with is the corotation velocity of the emission point at altitude $r$ and $r_L$ is the light cylinder radius."
+" One can transform the 0, aud y, values to the magnetic coordinates through The pulses will be observed by us when the direction (0,.45.) comcides with our line of sight."," One can transform the $\hat \theta_r$ and $\hat \varphi_r$ values to the magnetic coordinates through The pulses will be observed by us when the direction $(\hat \theta_r, \,\hat \varphi_r)$ coincides with our line of sight."
+" If the direction of Lue of sight is specified by angles (05s.φις) then the sweep of the line of sight through the ciission cap 1s giveu by where the angle y, will change by an amount Of in time interval f but 0, remains constant."," If the direction of line of sight is specified by angles $(\theta_{LS}, \, \varphi_{LS})$ then the sweep of the line of sight through the emission cap is given by where the angle $\hat \varphi_{r}$ will change by an amount $\Omega \, t$ in time interval $t$ but $\hat \theta_r$ remains constant."
+ Here O is the pulsar augular velocity., Here $\Omega$ is the pulsar angular velocity.
+"The HLC is stored in eight FITS binary tables. one for each hour in RA from 20"" to 4.","The HLC is stored in eight FITS binary tables, one for each hour in RA from $^{\mbox{\small h}}$ to $^{\mbox{\small h}}$."
+ Each record in the HLC corresponds to a single light-motion curve and stores 229 derived quantities with tag names listed and described in Tables 4. and 5.., Each record in the HLC corresponds to a single light-motion curve and stores 229 derived quantities with tag names listed and described in Tables \ref{tab:hlc1} and \ref{tab:hlc2}. .
+ The FITS files (~1.9 Gb compressed) may be obtained by web download frompublic., The FITS files $\sim$ 1.9 Gb compressed) may be obtained by web download from.
+" The IDL Astronomy User's function ""mrdtits"" is a convenient way to read-in the HLC. storing the FITS binary table automatically in an IDL structure."," The IDL Astronomy User's function “mrdfits” is a convenient way to read-in the HLC, storing the FITS binary table automatically in an IDL structure."
+" The IDL ""where"" function then becomes a very powerful tool for accessing subsets of data with euse.", The IDL “where” function then becomes a very powerful tool for accessing subsets of data with ease.
+ The most relevant comparison of our catalogue photometry can be made by comparing our results to those of Ivezicetal.2007) (from now on IVO7) who use 58 pre-2005 SDSS-I imaging runs of Stripe 82. observed under mostly photometric conditions. to construct a standard star catalogue.," The most relevant comparison of our catalogue photometry can be made by comparing our results to those of \citet{ive2007} (from now on IV07) who use 58 pre-2005 SDSS-I imaging runs of Stripe 82, observed under mostly photometric conditions, to construct a standard star catalogue."
+ TVO7 match unsaturated point sources satisfying high signal-to-noise criteria (photometric uncertainties below 0.05 mag) between runs and choose sources with at least 4+ epochs., IV07 match unsaturated point sources satisfying high signal-to-noise criteria (photometric uncertainties below 0.05 mag) between runs and choose sources with at least 4 epochs.
+ Those objects that are identified as non-varying. by requiring a X7 per degree of freedom for the mean magnitudes of less than 3 in the g. i and 7 bands. are chosen as candidate standard stars.," Those objects that are identified as non-varying, by requiring a $\chi^{2}$ per degree of freedom for the mean magnitudes of less than 3 in the $g$, $r$ and $i$ bands, are chosen as candidate standard stars."
+ The standard star catalogue is then internally recalibrated. correcting for flat-tield errors and. time-variable extinction. (for example. due to clouds).," The standard star catalogue is then internally recalibrated, correcting for flat-field errors and time-variable extinction (for example, due to clouds)."
+ The main differences between the LMCC and the IV0O7 catalogue are that we use a different photometric recalibration method. but based on the same idea of spatially dependent photometric zeropoints. and that we include the generally non-photometric observations from the SDSS-II supernova runs.," The main differences between the LMCC and the IV07 catalogue are that we use a different photometric recalibration method, but based on the same idea of spatially dependent photometric zeropoints, and that we include the generally non-photometric observations from the SDSS-II supernova runs."
+ Not surprisingly. for the 886396 objects from the IVO7 catalogue in the overlap area with the LMCC. we find unambiguous matches in the LMCC for 886191 objects using a match radius of0.7”. where the missing fraction of ~0.02% can be explained by the different criteria used to construct the two catalogues.," Not surprisingly, for the 886396 objects from the IV07 catalogue in the overlap area with the LMCC, we find unambiguous matches in the LMCC for 886191 objects using a match radius of, where the missing fraction of $\sim$ can be explained by the different criteria used to construct the two catalogues."
+ Of the 886191 matching objects in the LMCC. 219627. 606658. 879766. 886028 and 856456 objects have at least 3 good epochs in the i. g. r. Pand z bands. respectively. and mean PSF magnitudes brighter than 21.5. 21.5. 21.5. 21.5 and 20.5 mag. respectively.," Of the 886191 matching objects in the LMCC, 219627, 606658, 879766, 886028 and 856456 objects have at least 3 good epochs in the $u$, $g$, $r$, $i$ and $z$ bands, respectively, and mean PSF magnitudes brighter than 21.5, 21.5, 21.5, 21.5 and 20.5 mag, respectively."
+ It is these objects that we use when constructing the histograms in Figure 8.., It is these objects that we use when constructing the histograms in Figure \ref{fig:comp_ive}. .
+ In Figure 8.. each row of panels corresponds to a different wave band and the order employed is 5. g. r. ¢ and z from the top row to the bottom row.," In Figure \ref{fig:comp_ive}, , each row of panels corresponds to a different wave band and the order employed is $u$, $g$, $r$, $i$ and $z$ from the top row to the bottom row."
+ The left hand column of plots are normalised histograms of 47 per degree of freedom for the mean PSF magnitudes in the HLC. and they serve to confirm that the photometric uncertainties in the LMCC (inherited from the SDSS pipelines. with some adjustment during photometrie recalibration) are correct. since the histograms peak at a value of yr.," The left hand column of plots are normalised histograms of $\chi^{2}$ per degree of freedom for the mean PSF magnitudes in the HLC, and they serve to confirm that the photometric uncertainties in the LMCC (inherited from the SDSS pipelines, with some adjustment during photometric recalibration) are correct, since the histograms peak at a value of $\chi^{2}~/~(N-1)~\sim$ 1."
+ The middle column of panels in Figure δ are normalised ustograms of the uncertainties on the mean PSF magnitudes in 1ο HLC (solid lines) and on the mean magnitudes from the IV07 catalogue (dashed lines)., The middle column of panels in Figure \ref{fig:comp_ive} are normalised histograms of the uncertainties on the mean PSF magnitudes in the HLC (solid lines) and on the mean magnitudes from the IV07 catalogue (dashed lines).
+ It is clear that the mean magnitudes cuoted in the HLC are approximately twice as precise as those from IVO7. which is due to the greater number of epochs included in ye LMCC.," It is clear that the mean magnitudes quoted in the HLC are approximately twice as precise as those from IV07, which is due to the greater number of epochs included in the LMCC."
+ However. systematic errors are still likely to dominate je mean magnitudes at the ~1%e level. which corresponds to the evel of the photometric recalibrations and to the systematic error introduced by the slightly different band passes of each column of detectors in the SDSS camera.," However, systematic errors are still likely to dominate the mean magnitudes at the $\sim$ level, which corresponds to the level of the photometric recalibrations and to the systematic error introduced by the slightly different band passes of each column of detectors in the SDSS camera."
+ We do not correct the LMCC magnitudes for the different bandpasses in contrast to IV07., We do not correct the LMCC magnitudes for the different bandpasses in contrast to IV07.
+ The right hand column of panels in Figure 8. are normalised histograms of the difference between the HLC mean PSF magnitudes and the IVO7 catalogue mean magnitudes., The right hand column of panels in Figure \ref{fig:comp_ive} are normalised histograms of the difference between the HLC mean PSF magnitudes and the IV07 catalogue mean magnitudes.
+ A Gaussian has been fitted to each histogram (dotted lines). and the fitted centre and sigma are quoted in each panel.," A Gaussian has been fitted to each histogram (dotted lines), and the fitted centre and sigma are quoted in each panel."
+ The histograms are centred around zero at the millimagnitude level. except for the 1 band where the IV07 photometry is slightly offset from the HLC photometry by —4 mmag.," The histograms are centred around zero at the millimagnitude level, except for the $u$ band where the IV07 photometry is slightly offset from the HLC photometry by $\sim$ 4 mmag."
+ The histograms have sigmas of 28. 13. HI. 10 and 16 mmag in ou. g. r. 7 and z. respectively. which are consistent with the scatter in the photometric zeropoints of the standard SDSS runs derived in Section 2.2. and with the scatter in the similar internal photometric recalibrations of IVO7.," The histograms have sigmas of 28, 13, 11, 10 and 16 mmag in $u$, $g$, $r$, $i$ and $z$, respectively, which are consistent with the scatter in the photometric zeropoints of the standard SDSS runs derived in Section 2.2, and with the scatter in the similar internal photometric recalibrations of IV07."
+ The data set used to construct the IVO7 standard star catalogue is à subset of the data used to construct the LMCC. and therefore we ean use oursuperior temporal coverage to identify standard star candidates that are actuallyphotometrically variable.," The data set used to construct the IV07 standard star catalogue is a subset of the data used to construct the LMCC, and therefore we can use oursuperior temporal coverage to identify standard star candidates that are actuallyphotometrically variable."
+ Watkinset use the HLC to systematically identify variable stars in Stripe 82. and we choose to follow all but one of their cuts," \citet{wat2008} use the HLC to systematically identify variable stars in Stripe 82, and we choose to follow all but one of their cuts"
+"And ""Transient. Source Experiment (BATSE) correlation reported by Foster et 11996.Erplover Proportional. Counter Array(RAPE PCA) correlation reported by PE97). and PROT were also the first to observe requeney-cepencdent delavs in the radio emission.","And Transient Source Experiment (BATSE) correlation reported by Foster et 1996, Proportional Counter Array PCA) correlation reported by PF97), and PF97 were also the first to observe frequency-dependent delays in the radio emission."
+ Quasi-»eriodie Daring in the infrared. reported by Fender Pooley (1998). was also observed by Eikenberry ct (10908). who racl simultaneous observations withANTE PCA.," Quasi-periodic flaring in the infrared, reported by Fender Pooley (1998), was also observed by Eikenberry et (1998), who had simultaneous observations with PCA."
+" Phe source of these flares is consistent with infrared emitting plasmons rommaterial, following rapid X-ray variability and an X-ray dip."," The source of these flares is consistent with infrared emitting plasmons from, following rapid X-ray variability and an X-ray dip."
+ The confirmation of such a hypothesis was presented by Mirabel et ((1998) with simultaneousRNTE PCA. United Ixingdom Infrared Telescope CUINLICE) and Very Large Array (VLA) data.," The confirmation of such a hypothesis was presented by Mirabel et (1998) with simultaneous PCA, United Kingdom Infrared Telescope (UKIRT) and Very Large Array (VLA) data."
+ Several X-ray dips and subsequent infrared. and racio Dares were observed., Several X-ray dips and subsequent infrared and radio flares were observed.
+ We obtained sub-mim data over. two. days on. MJD 51026 and 51027. when GRS 1915)105 was mocderatelv active. with periods of quiescence interspersed.," We obtained sub-mm data over two days on MJD 51026 and 51027, when GRS 1915+105 was moderately active, with periods of quiescence interspersed."
+Co-ordinated with this were quasi-simullancous observations at. racio wavelengths., Co-ordinated with this were quasi-simultaneous observations at radio wavelengths.
+ The radio data were taken using the Green Bank Interferometer (GBI) at 2.25 and 8.3 Ον. and the RP at 15 Giz.," The radio data were taken using the Green Bank Interferometer (GBI) at 2.25 and 8.3 GHz, and the RT at 15 GHz."
+ Sub-mim dita at 350 CLHIz was taken using the James Clark Maxwell Telescope (JCAL) using the SCUDA array. and X-ray data were obtained at 212 keV [rom theRNTE ASAI instrument.," Sub-mm data at 350 GHz was taken using the James Clark Maxwell Telescope (JCMT) using the SCUBA array, and X-ray data were obtained at 2–12 keV from the ASM instrument."
+ The GBI monitored CRS 1915|105 simultaneously. at 2.25 and 8.3 Gllz with2. integrations. 10.15 minutes. in: length several times. dailv., The GBI monitored GRS 1915+105 simultaneously at 2.25 and 8.3 GHz with integrations 10–15 minutes in length several times daily.
+: The calibration: procedure is. reported in detail in Waltman et ((1994)., The calibration procedure is reported in detail in Waltman et (1994).
+ The RTI is an 8-clement castwest array operating at. 15-Gllz: details the operation. can be found: in| PE97., The RT is an 8-element east–west array operating at 15-GHz; details of the operation can be found in PF97.
+ OeThe ΟΝΕ ISub-nim Common User⊽ Bolometer Array (SCUBA)⊽ observed CIUS 1915|105 over two nights on ALJD 51026 ancl 51027 (1998 Aug 01 and 02) in photometric mode at 450 ancl S50 pum wavelengths' lor a description of the instrument see Llollaned et ((1999)., The JCMT Sub-mm Common User Bolometer Array (SCUBA) observed GRS 1915+105 over two nights on MJD 51026 and 51027 (1998 Aug 01 and 02) in photometric mode at 450 and 850 $\umu$ m wavelengths – for a description of the instrument see Holland et (1999).
+)eule Reeularkv] skvedipsak were takeng throughoutfrequencyercessalin the observations.-- and 5.after focusing.MESHES aligning ANCand pointingQm the array. [lux calibration was performed. using ?either Uranus or G45.1 (ILA(1950)=1911. Dec.(1950)=10? 45).," Regular skydips were taken throughout the observations, and after focusing, aligning and pointing the array, flux calibration was performed using either Uranus or G45.1 $\rm{RA}(1950)=19^{\rm{h}}~11^{\rm{m}}$, $\rm{Dec.}~(1950)=10^{\circ}~45'$ )."
+ A total of 15 observations o£ GRS 1915|105 were mace over the two nights.: cach observation ..included 50 integrations. of 10 s (except For the observation at 51026.574 which included 26 integrations).," A total of 15 observations of GRS 1915+105 were made over the two nights, each observation included 50 integrations of 10 s (except for the observation at 51026.574 which included 26 integrations)."
+ ref1026.1tMM Iryshowspholometryoverlwonighls AM hl .JD510260nd51021T," \\ref{1026,1027_photometry} shows photometry over two nights: MJD 51026 and 51027."
+".Enthefigurethetoppanelshowsthe N rayflicr, ATE ΕΣ... πο mmfl'"," In the figure, the top panel shows the X-ray flux from the ASM instrument, whereas the bottom panel shows the radio and sub-mm flux."
+irablelpresentsthefliurrangeduringobservalions. al ideret etlobscrentione rrovForte foreach phoromelCveX-3 rte indicespotaweal divean aprojected fl," Table \ref{weighted_average} presents the flux range during observations, a typical observation error for each photometric point, a weighted mean, a projected flux (detailed below), and errors on these fluxes."
+ In all radio. sub-mun and. X-ray. bands our data show GIU 1915]105 to be variable.," In all radio, sub-mm and X-ray bands our data show GRS 1915+105 to be variable."
+ At all wavelengths. there is a general trend. to lower fluxes ancl smaller. variability over the timescale of our observations.," At all wavelengths, there is a general trend to lower fluxes and smaller variability over the timescale of our observations."
+ roleBhiecayshou sthephotome lry omtheC DIatsCH Esandlhebbat, \\ref{RT_GBI_decay} shows the photometry from the GBI at 8 GHz and the RT at 15 GHz for 7 d around the time of our SCUBA observations.
+ 1561," The flux is in general decaying, with some fluctuations around days 51023 and 51025."
+ mmobservalionsondagsb 10260 Τη, The decay becomes steadier around the time of the sub-mm observations on days 51026 and 51027.
+ We have modelled. the decay the last three cays shown in rolUeD Loecagatallibrecradiowarelengihswithatinear [it, We have modelled the decay of the last three days shown in \\ref{RT_GBI_decay} at all three radio wavelengths with a linear fit.
+ Lhisprovidesapl flireswheninterpolatedtothelimeswbhenthesub mimedalawerelaken.providinglherewerenoraresharpoulbursts.," This provides a plausible estimate to the cm-wave fluxes when interpolated tothe times when the sub-mm data were taken, providing there were no rare sharp outbursts."
+Aeeragedpro The, Averaged projected fluxes using this method are given in Table \ref{weighted_average}.
+ main dillieultv. in interpretation of these data is the validity of extrapolating a variable source using a simple constant decay., The main difficulty in interpretation of these data is the validity of extrapolating a variable source using a simple constant decay.
+ Lf data at all radio frequencies follows the simple decay law. ancl no additional Lares were produced at the time of the SCUBA observations. then the spectra over⊽⋅ the (wo epochs in⋠⊔ re[proj.pectrashousaclearsub mimercess.," If data at all radio frequencies follows the simple decay law, and no additional flares were produced at the time of the SCUBA observations, then the spectra over the two epochs in \\ref{proj_spectra} shows a clear sub-mm excess."
+ M hiletheselastlwoassumplionsmagbeoversimplif yinglherealilj din P E9Tshowsalargeshort flare).. ; preciousobservalionshaveshownahigh; ; frarcdeomparcdwilhradioobservations(⇁ hender& ⇁ ⇁ ↴ ," While these last two assumptions may be over simplifying the reality of GRS 1915+105 4 in PF97 shows a large short-lived flare), previous observations have shown a high-frequency excess at infrared compared with radio observations (Fender Pooley 1998; Mirabel et 1998)."
+Pootal. yandwithmullipleoscillalions. ⇁ . lhberadiospectrimatalimcwhenthe. . | mJv) is very similar to that in this paper 11 in Mirabel et 11998). and during times when the radio flux at S Gllz is around. 15 mJy. the 2.2 jun flux is 35-40 ms.," Photometric observations by these authors have shown that over a large range in radio flux (10--100 mJy), and with multiple oscillations, the radio spectrum at a time when the flux is low $\simeq 20$ mJy) is very similar to that in this paper 1 in Mirabel et 1998), and during times when the radio flux at 8 GHz is around 15 mJy, the 2.2 $\umu$ m flux is 35-40 mJy."
+ The flat spectrum (= high-frequency excess) presented here is seen in other X-ray. binaries where there exists significant “μμ sub sub-mm regimes., The flat spectrum $\equiv$ high-frequency excess) presented here is seen in other X-ray binaries where there exists significant emission at the mm and sub-mm regimes.
+"uM ;Both Cyg. N-1. ypic and, (Fende iglitei )glev et 998) have “spectra with ολη, àc deviating from the expected optically thin. πας 0.5. steep spectrum observed during an outburst (where spectral index à as S,x v7"")."," Both Cyg X-1 (Fender et 2000) and Cyg X-3 (Fender et 1995; Ogley et 1998) have spectra with indices $\alpha \simeq 0$, deviating from the expected optically thin, $-1 \leq \alpha \leq -0.5$ , steep spectrum observed during an outburst (where spectral index $\alpha$ as $S_{\nu} \propto \nu^{\alpha}$ )."
+ As discussed. by other. authors. a sub-mim excess ancl approximately [lat svnchrotron spectrum: are. likely," As discussed by other authors, a sub-mm excess and approximately flat synchrotron spectrum are likely"
+1997)).,\cite{forb97}) ).
+ Spectroscopic analysis of 18 GCs by Nissler-Patig etal. (1998)), Spectroscopic analysis of 18 GCs by Kissler-Patig etal. \cite{kiss98a}) )
+ shows a ictallicity range between 1.6 and 0.3 dex (with possible peaks at 1.3 and 0.6 dex). and two exceptional GCs at about 0.2 dex. located in the red (inetal rich) tail of the color distributions.," shows a metallicity range between $-1.6$ and $-0.3$ dex (with possible peaks at $-1.3$ and $-0.6$ dex), and two exceptional GCs at about 0.2 dex, located in the red (metal rich) tail of the color distributions."
+ The comparison of the line indices with theoretical evolutionary models sugecsts that most of the GCs are older than at least s Cir., The comparison of the line indices with theoretical evolutionary models suggests that most of the GCs are older than at least 8 Gyr.
+ If one fits the GC color distribution with two Gaussians. the nunboer ratio of metal rich (red) to metal poor (blue) GCs is about 1:1 (Forbes et al. 1997)).," If one fits the GC color distribution with two Gaussians, the number ratio of metal rich (red) to metal poor (blue) GCs is about 1:1 (Forbes et al. \cite{forb97}) )."
+ The radial extension of the GCS around NCC 1399 cau be trace out to about 10 (253 kpe}., The radial extension of the GCS around NGC 1399 can be traced out to about $10\arcmin$ $\simeq 53$ kpc).
+" The slope ofthe GC surface deusity profile. pXr. ijs about à=—1.5x0.2. ""heu taking the average of the published values;"," The slope of the GC surface density profile, $\rho \propto r^\alpha$, is about $\alpha = -1.5\pm0.2$, when taking the average of the published values."
+ Forbes et al. (1997)), Forbes et al. \cite{forb97}) )
+ found that the distribution of the blue GC subpopulation is even flatter (021.023:0.2). whereas the red GCs are more centrally couceutrated (02ντος 2). conrparable to the slope of the galaxy ight 0.1).," found that the distribution of the blue GC subpopulation is even flatter $\alpha \simeq 
+-1.0\pm0.2$ ), whereas the red GCs are more centrally concentrated $\alpha 
+\simeq -1.7\pm0.2$ ), comparable to the slope of the galaxy light $\alpha =
+-1.6\pm0.1$ )."
+ See Fig., See Fig.
+ 2 for a schematic overview., \ref{fsdslope} for a schematic overview.
+ Racial velocities of TL GCs around NCC 1399. have been measured (Ixissler-Patig et al. 1999.. ," Radial velocities of 74 GCs around NGC 1399 have been measured (Kissler-Patig et al. \cite{kiss99a}, ,"
+Minniti ct al. 1998..," Minniti et al. \cite{minn98},"
+ Ixissler-Patig et al. 1998))., Kissler-Patig et al. \cite{kiss98a}) ).
+" The velocity dispersio- for the whole sample is e,=373c20 kan 1", The velocity dispersion for the whole sample is $\sigma_v = 373\pm35$ km $^{-1}$.
+ No differences can be seen between the red aud blue subpopulations., No differences can be seen between the red and blue subpopulations.
+" However. there exists a radial depeudence of the velocity dispersion iu the seuse that o, rises from 263292 to LOS+LOT lan | between 2 and 8! Patig ct al. 1999))."," However, there exists a radial dependence of the velocity dispersion in the sense that $\sigma_v$ rises from $263\pm92$ to $408\pm107$ km $^{-1}$ between $2\arcmin$ and $8\arcmin$ (Kissler-Patig et al. \cite{kiss99a}) )."
+ The ealaxy light of NGC 1399 follows an extended ¢D profile (Schombert 1986.. Killeeu. Dicknell 1988)) out to a radial distance of about 31 arcu from the galaxy center (Xj=28 mag isopliotal surface brightuess level).," The galaxy light of NGC 1399 follows an extended cD profile (Schombert \cite{scho86}, Killeen Bicknell \cite{kill}) ) out to a radial distance of about 34 arcmin from the galaxy center $\Sigma_B = 28$ mag isophotal surface brightness level)."
+ This is abou 180 kpc in Fornax distance (18.2 Mpc) and conrparable to the exteut of the N-rav. envelope (sche et al. 1996..," This is about 180 kpc in Fornax distance (18.2 Mpc) and comparable to the extent of the X-ray envelope (Ikebe et al. \cite{ikeb},"
+ Joues et al. 19973)., Jones et al. \cite{jone97}) ).
+ The determination of the stellar population paraiucters of the outer ¢D halo. like accurate photometric colors. mnctallicity or velocity dispersion. is very difficult due to the low surface brightuess.," The determination of the stellar population parameters of the outer cD halo, like accurate photometric colors, metallicity or velocity dispersion, is very difficult due to the low surface brightness."
+ Long slit spectra have been taken for the stellar bulge population within a radius of about 1/5 frou the ceuter of NGC 1399 (Fraus et al. 1989..," Long slit spectra have been taken for the stellar bulge population within a radius of about $1\farcm5$ from the center of NGC 1399 (Franx et al. \cite{fran},"
+ Bickuell ct al. 1989))., Bicknell et al. \cite{bick}) ).
+ The velocity dispersion is about 200 kus + at 1/5 aud rises within the ceutral 1ff to a ceutral value of about 360 aus I., The velocity dispersion is about 200 km $^{-1}$ at $1\farcm5$ and rises within the central $10\arcsec$ to a central value of about 360 km $^{-1}$.
+" Besides the GCS, an useful tracer for the stellar »opulation at larger radii is the population of plauctary nebulae (PNe)."," Besides the GCS, an useful tracer for the stellar population at larger radii is the population of planetary nebulae (PNe)."
+ Arnaboldi et al. (199 1)., Arnaboldi et al. \cite{arna}) )
+ studied the kinematics of 37 PNe out to a radius of 1/5., studied the kinematics of 37 PNe out to a radius of $4\farcm5$.
+" They fouud au increase in the velocity dispersion with increasing radius from 269 kan | for ro«216 to 105 kins + for 2/6<rUS 8 of the 37 Όλο),", They found an increase in the velocity dispersion with increasing radius from 269 km $^{-1}$ for $r < 2\farcm6$ to 405 km $^{-1}$ for $2\farcm6 < r < 4\farcm5$ (18 of the 37 PNe).
+ Iu this subsection we divide the light profile of NGC 1399 into a cD halo aud a bulee component in order to compare their characteristics with those of the GCS aud the dwartgalaxy population., In this subsection we divide the light profile of NGC 1399 into a cD halo and a bulge component in order to compare their characteristics with those of the GCS and the dwarfgalaxy population.
+ We determined the absolute huninosity of the ¢D halo in the following wav: in the port’! plot (Fie. 1. ," We determined the absolute luminosity of the cD halo in the following way: in the $\mu$ $r^{1/4}$ plot (Fig. \ref{fcdprof}, ,"
+upper panels) one cau see that the SB profile of NGC 1399 (determined from the NE CCD, upper panels) one can see that the SB profile of NGC 1399 (determined from the NE CCD
+Photometric measurements of the αubined light of the svstemi have been made by a nuniber of authors over a broad range of colours.,Photometric measurements of the combined light of the system have been made by a number of authors over a broad range of colours.
+" They are summarised and homogenized in the critical compilation of Legectt (1992). from which we adopt: U-10.79. B=9.57 Ves.09, R=7.09. I-5.97. J=l75. II-L11. WK=3.93. L=3.77. L'23.67. with quoted errors of <3% on TIT."," They are summarised and homogenized in the critical compilation of Leggett (1992), from which we adopt: U=10.79, B=9.57 ,V=8.09, R=7.09, I=5.97, J=4.75, H=4.14, K=3.93, L=3.77, L'=3.67, with quoted errors of $<$ on IJHK."
+ The system has colours consistent with those of solar metallicitv stars (Legeettao 1992). as expected from the spectroscopic iietallicitv discussed below.," The system has colours consistent with those of solar metallicity stars (Leggett 1992), as expected from the spectroscopic metallicity discussed below."
+ Absolute magnitudes were derived from the parallax aud the apparent magnitudes of the individual stars., Absolute magnitudes were derived from the parallax and the apparent magnitudes of the individual stars.
+ The abxJute magnitudes of the brighter star lave 1awcertaimties which are dominated by those of the photonsctry of the system. while those for the secoudary have some contribution from the magnitude difference.," The absolute magnitudes of the brighter star have uncertainties which are dominated by those of the photometry of the system, while those for the secondary have some contribution from the magnitude difference."
+ The parallax doesut appreciably contribute in either case., The parallax doesn't appreciably contribute in either case.
+ The accuvate masses and absolute magnitudes that we jiwe obtained for the Cl 570DC system represent a rew beinluuark for model calculations (e.g. Daraffe oet al., The accurate masses and absolute magnitudes that we have obtained for the Gl 570BC system represent a new benchmark for model calculations (e.g. Baraffe et al.
+ 1998) and an iudepenudeut check of the empirical Ποηiuositv relations (ΗΛΙΟ)., 1998) and an independent check of the empirical mass-luminosity relations (HMcC).
+ The ουκταας which hey brinο to the models are largely complementary to hose conimg from the eclipsing binaries. whose absolute radi cau be determined very accurately but whose arecr distances on the other haud coutribute significant uncertainties to the absolute magnitudes.," The constraints which they bring to the models are largely complementary to those coming from the eclipsing binaries, whose absolute radii can be determined very accurately but whose larger distances on the other hand contribute significant uncertainties to the absolute magnitudes."
+ As |me ouphaszed by theoreticians. and bx Observers of mmere massive stars (c.g. Andersen 1991). there is however no such thing as oue single niasluninositv relation: stellar luminosities depend on chemical Cuposition as well as ou mass (in general they depend o rage too. though not iu the age aud mass rauge ciscussec here).," As long emphasized by theoreticians, and by observers of more massive stars (e.g. Andersen 1991), there is however no such thing as one single mass-luminosity relation: stellar luminosities depend on chemical composition as well as on mass (in general they depend on age too, though not in the age and mass range discussed here)."
+ Quantitative metallicity determinations however are notoriously difficult for AL dwarfs (e.g. Viti et al., Quantitative metallicity determinations however are notoriously difficult for M dwarfs (e.g. Viti et al.
+ 1997: Valeuti et al., 1997; Valenti et al.
+ 1998)., 1998).
+ Observers iu this field usuuly have to resort to photometric metallicity estimators (Legeett 1992) which are only approximately calibrated. or otherwise assume by default a solar," Observers in this field usually have to resort to photometric metallicity estimators (Leggett 1992) which are only approximately calibrated, or otherwise assume by default a solar"
+"Expressed in terms of their volume scalars. tlie deceleration parameters associated with the a, aud πμ [rames are respectively given by Our main question is whether d can take negative values while g is still positive.","Expressed in terms of their volume scalars, the deceleration parameters associated with the $u_a$ and $\tilde{u}_a$ frames are respectively given by Our main question is whether $\tilde{q}$ can take negative values while $q$ is still positive."
+ If so. the (πασά observers will experience accelerated expansion in a decelerating universe.," If so, the tilded observers will experience accelerated expansion in a decelerating universe."
+ To investigate this possibility. we first use clefinitious (1)) to recast expressious (3)) iuto respectively.," To investigate this possibility, we first use definitions \ref{eq:qs1}) ) to recast expressions \ref{eq:Rays1}) ) into respectively."
+ These already slow that q aud qd are generally different. but it helps to relate the two deceleration parameters directly.," These already show that $q$ and $\tilde{q}$ are generally different, but it helps to relate the two deceleration parameters directly."
+" Recall that 6=p and d,=6;4(0/3)0, to linear order in e.", Recall that $\tilde{\rho}=\rho$ and $\tilde{A}_a=\dot{v}_a+(\Theta/3)v_a$ to linear order in $v_a$ .
+ Then. employing definition 0= τομ. relation (2)) aud keeping up to c;-order terms. expressions (5aa) and (bb) combine to where Q. 0>O always.," Then, employing definition $\vartheta=\tilde{\rm D}^av_a$ , relation \ref{eq:Thetas}) ) and keeping up to $v_a$ -order terms, expressions \ref{eq:Rays2}a a) and \ref{eq:Rays2}b b) combine to where $\Theta$, $\vartheta>0$ always."
+" We may also involve the volume scalar of the peculiar motion further by using the (linear in e) relation )=D""6,—00/3 (Ellis&Tsagas2002)..", We may also involve the volume scalar of the peculiar motion further by using the (linear in $v_a$ ) relation $\dot{\vartheta}=\tilde{\rm D}^a\dot{v}_a -\Theta\vartheta/3$ \citep{ET}.
+ Then. Eq. (6))," Then, Eq. \ref{eq:tq1}) )"
+ leads to given that O=O41)., leads to given that $\tilde{\Theta}=\Theta+\vartheta$.
+ The above relates the deceleration parameter iu the tilded [rane to that oL the actual universe and it is our main result., The above relates the deceleration parameter in the tilded frame to that of the actual universe and it is our main result.
+ It should wow be clear that q and q are generally dillereut., It should now be clear that $q$ and $\tilde{q}$ are generally different.
+ Moreover. as loug as the right-hand side of (7)) remains below unity. positive values [or q do not a priori guarantee the same for q.," Moreover, as long as the right-hand side of \ref{eq:tq2}) ) remains below unity, positive values for $q$ do not a priori guarantee the same for $\tilde{q}$."
+ Iu other words. it is theoretically possible lor the tiklec observer to experience accelerated expansion in a decelerating universe.," In other words, it is theoretically possible for the tilded observer to experience accelerated expansion in a decelerating universe."
+ Putting it differently. oie could say that meastuing uegative deceleration parameter iu a frame drifting relative to theCMIB (like that of our Local Group for example) does not uecessarily imply an accelerating universe.," Putting it differently, one could say that measuring negative deceleration parameter in a frame drifting relative to theCMB (like that of our Local Group for example) does not necessarily imply an accelerating universe."
+ At this point it is worth noting that. accordingto (Lbb). condition —1«d<0 is equivalent to —07/3XDog.«»O 0.," At this point it is worth noting that, accordingto \ref{eq:qs1}b b), condition $-1<\tilde{q}<0$ is equivalent to $-\tilde{\Theta}^2/3<\dot{\tilde{\Theta}}<0$ ."
+ This⋅ meaus that both d~ aud OS canbe simultaneously. negative.., This means that both $\tilde{q}$ and $\dot{\tilde{\Theta}}$ canbe simultaneously negative.
+ Analogous, Analogous
+dduring the 2010 episode is fairly typical of BH outbursts 11 general.,during the 2010 episode is fairly typical of BH outbursts in general.
+ Our spectral analysis showed that the source transited from at initial LHS. with à spectrum dominated by Comptonizatior in the beginning of the outburst (~ 55217— 55281 MJD). to softer states where the power law cut-off was not needec anymore and the photon index of the power law was very soft (~ 2.3).," Our spectral analysis showed that the source transited from an initial LHS, with a spectrum dominated by Comptonization in the beginning of the outburst $\sim$ 55217– 55281 MJD), to softer states where the power law cut-off was not needed anymore and the photon index of the power law was very soft $\sim$ 2.3)."
+ In X-ray binaries. a cut-off power law spectrum is usually interpreted as the signature of inverse Comptonisatior of soft seed photons by a thermalised (re.. with velocities following a Maxwellian distribution) population of electrons.," In X-ray binaries, a cut-off power law spectrum is usually interpreted as the signature of inverse Comptonisation of soft seed photons by a thermalised (i.e., with velocities following a Maxwellian distribution) population of electrons."
+ Changes in the hard component can signal the presence of a compact jet. a corona. or reprocessed hard X-ray emission due to X-ray heating from an extended central source.," Changes in the hard component can signal the presence of a compact jet, a corona, or reprocessed hard X-ray emission due to X-ray heating from an extended central source."
+ The hard component therefore evolvec as expected., The hard component therefore evolved as expected.
+ Regardless of the caveats mentioned in Sect., Regardless of the caveats mentioned in Sect.
+ 3.3. about the components. and the exact value of the inner radius obtained from the fits. we see a trend in the disc parameters (Table 2 and Fig. 6))," \ref{xgamspec} about the components, and the exact value of the inner radius obtained from the fits, we see a trend in the disc parameters (Table \ref{tab:para} and Fig. \ref{param-boff}) )"
+ which is compatible with an increase of the mass accretior rate before the transition., which is compatible with an increase of the mass accretion rate before the transition.
+ Then the dise globally recedes whe the source reaches a low luminosity (see.e.g..?).. during the decay phase of the outburst.," Then the disc globally recedes when the source reaches a low luminosity \citep[see, e.g.,][]{tomsick09}, during the decay phase of the outburst."
+ Although we have then observed the dise slowly moving outwards during the hardening (??).. this is not necessarily true for all transient sources (see. e.g.. ?)).," Although we have then observed the disc slowly moving outwards during the hardening \citep{Chen:1997,CadolleBel:2004}, this is not necessarily true for all transient sources (see, e.g., \citealt{Miller:2006}) )."
+ This is still strongly debated (???) as. for example. i XTE J1817—330.," This is still strongly debated \citep{Done:2007,Rykoff:2007,Gierlinski:2008} as, for example, in XTE $-$ 330."
+ We do not address this question specifically in this work as the majority of our data were taken in the LHS and not long after the main transition into softer states., We do not address this question specifically in this work as the majority of our data were taken in the LHS and not long after the main transition into softer states.
+ In addition. with a lower energy limit often at 3 keV. few data are sensitive enough for this purpose.," In addition, with a lower energy limit often at 3 keV, few data are sensitive enough for this purpose."
+ However. one can remark that the rather low disc temperatures are consistent with the compact object being a BH and not a neutron star (e.g..?)..," However, one can remark that the rather low disc temperatures are consistent with the compact object being a BH and not a neutron star \citep[e.g.,][]{Tanaka:1995}."
+ The soft and hard X-rays evolved as seen in previous outbursts of339-4: first in the HIMS. the source transited to softer states where the dise dominated the emission. the inner disc moved closer to the BH and then gradually cooled down before returning to quiescence.," The soft and hard X-rays evolved as seen in previous outbursts of: first in the HIMS, the source transited to softer states where the disc dominated the emission, the inner disc moved closer to the BH and then gradually cooled down before returning to quiescence."
+ The spectral evolution is consistent with other works: mmonitoring showed softening and LFQPO evolutions (22?) as sstarted to leave the LHS (?)..," The spectral evolution is consistent with other works: monitoring showed softening and LFQPO evolutions \citep{shap10,motta10,yu10} as started to leave the LHS \citep{motta10}."
+ Fig., Fig.
+ 6 shows some differences with the Fig., \ref{param-boff} shows some differences with the Fig.
+ 6 of ?:: while the fluxes. dise temperature and power law indices show the same trend. our individual cut-off values are sometimes very high (e. g.. near 55260 MJD). like in the HSS of their Fig.," 6 of \citet{Motta09}: while the fluxes, disc temperature and power law indices show the same trend, our individual cut-off values are sometimes very high (e. g., near 55260 MJD), like in the HSS of their Fig."
+ 6. and no cut-off is needed.," 6, and no cut-off is needed."
+ The evolution in the rising phase of the LHS is therefore not as smooth as expected at higher energies (which are well constrained for the first time). but there is no clear explanation for that.," The evolution in the rising phase of the LHS is therefore not as smooth as expected at higher energies (which are well constrained for the first time), but there is no clear explanation for that."
+ The 2-20 keV flux (from 1.43 to 14.6 x 107? ere em? s!) and the bolometric flux (extrapolated from 0.01 to 1000 keV) variations (0.82 to 2.93K10 erg emo? s!) confirm this spectral evolution., The 2–20 keV flux (from 1.43 to 14.6 $\times$ $^{-9}$ erg $^{-2}$ $^{-1}$ ) and the bolometric flux (extrapolated from 0.01 to 1000 keV) variations (0.82 to $\times$ $^{-8}$ erg $^{-2}$ $^{-1}$ ) confirm this spectral evolution.
+ The source reached a high luminosity at the maximum of ~ 12.9 x (d/6 kpe)? 1075 erg s7!., The source reached a high luminosity at the maximum of $\sim$ 12.9 $\times$ (d/6 $^2$ $^{37}$ erg $^{-1}$.
+ This represents ~ of the Eddington luminosity for a 7 solar mass BH: stellar mass BHs acereting at or below 107? Legg are found in the LHS (?).., This represents $\sim$ of the Eddington luminosity for a 7 solar mass BH; stellar mass BHs accreting at or below $^{-2}$ $L_{\rm Edd}$ are found in the LHS \citep{McClintock:2006}.
+ Besides. the evolution of the reflection component from 0.19 to 0.51 indicates that the source becamesofter.," Besides, the evolution of the reflection component from 0.19 to 0.51 indicates that the source becamesofter."
+ Later than the observations presented in this paper. the source switched back to the SIMS. and subsequently to the HIMS.," Later than the observations presented in this paper, the source switched back to the SIMS, and subsequently to the HIMS."
+ This is not unusual: during the decay of their outbursts. XTE J1720—318. XTE J1650—500 and SWIFT J1753.5-0127 showed a slowly receding dise with a decreasing inner temperature. while at the same time the relative amount of the power law contribution mereased again.," This is not unusual: during the decay of their outbursts, XTE $-$ 318, XTE $-$ 500 and SWIFT $-$ 0127 showed a slowly receding disc with a decreasing inner temperature, while at the same time the relative amount of the power law contribution increased again."
+ A year after the beginning of its outburst (end of February 2011). ffinally switched back to the LHS (?).. and then to quiescence (?)..," A year after the beginning of its outburst (end of February 2011), finally switched back to the LHS \citep{Russ11c}, and then to quiescence \citep{Russ11d}."
+ The broadband spectral parameter evolution from hard to soft states. as well as our radio/NIR/optical studies. are consistent with the preliminary results of ??..," The broadband spectral parameter evolution from hard to soft states, as well as our radio/NIR/optical studies, are consistent with the preliminary results of \citet[]{Lew10,corb10a}."
+ A discrete radio ejection usually occurs around the hard to soft transition in a BH transient outburst: ? witnessed with the ATCA the interaction of a relativistic jet from with the interstellar medium. implying that a major ejection occurred earlier.," A discrete radio ejection usually occurs around the hard to soft transition in a BH transient outburst: \cite{corb10b} witnessed with the ATCA the interaction of a relativistic jet from with the interstellar medium, implying that a major ejection occurred earlier."
+ During the LHS. we detected the compact core jet in the optical/IR.," During the LHS, we detected the compact core jet in the optical/IR."
+ The source continued to rise in the LHS for a while (?).. therefore 1t was also rising at radio frequencies.," The source continued to rise in the LHS for a while \citep{Wu10}, therefore it was also rising at radio frequencies."
+ We added radio flux data to our SED (Fig. 8)), We added radio flux data to our SED (Fig. \ref{sed1bisfit}) )
+ but the entire campaign will be presented in a forthcoming paper (Corbel et al., but the entire campaign will be presented in a forthcoming paper (Corbel et al.
+ 2011. in prep.).," 2011, in prep.)."
+ The rapid drop in optical flux and colour change observed at the start of the state transition are reminiscent of previous outbursts (?).., The rapid drop in optical flux and colour change observed at the start of the state transition are reminiscent of previous outbursts \citep{coriat09}.
+ Since rapid variability probably originates in the synchrotron Jet compoent (??).. this is consistent with the jet no longer making a strong contribution to the optical emission and that it was fading (?)..," Since rapid variability probably originates in the synchrotron jet component \citep{casella10,gand10}, this is consistent with the jet no longer making a strong contribution to the optical emission and that it was fading \citep{Ru10}."
+ Such chromatic behaviour. also observed in the REM data. hints for the presence of a component mainly contributing 1n the NIR frequencies which then fades or disappears after the transition (see also Fig. 7)).," Such chromatic behaviour, also observed in the REM data, hints for the presence of a component mainly contributing in the NIR frequencies which then fades or disappears after the transition (see also Fig. \ref{sed4}) )."
+ The optical and infrared flux continued to drop in the days following the transition., The optical and infrared flux continued to drop in the days following the transition.
+ The chages In optical flux and spectrum over the transition have been reported before for and other BH LMXBs (see. ??)..," The changes in optical flux and spectrum over the transition have been reported before for \\citep{Homan:2005b,coriat09} and other BH LMXBs \citep[see, e.g.,][]{bux04,russ07}."
+ In this paper. we clearly detect the jet evolution and its dramatic quenching at all wavelengths.," In this paper, we clearly detect the jet evolution and its dramatic quenching at all wavelengths."
+ To compute the NIR/optical spectral energy distributions (SEDs) of the source. we first corrected our magnitudes for interstellar absorption.," To compute the NIR/optical spectral energy distributions (SEDs) of the source, we first corrected our magnitudes for interstellar absorption."
+ We assumed a colour excess of E(B—V)1.2+0.1 (2)... which is fully consistent with our measurements of Ny and of ?..," We assumed a colour excess of $E(B-V) = 1.2 \pm 0.1$ \citep{Zd98}, which is fully consistent with our measurements of $N_{\rm H}$ and of \citet{corbfen02}."
+ Using standard extinction curves from ? and ?.. we obtained the de-reddening parameters for each of our UV/optical/NIR filters (Table 4)).," Using standard extinction curves from \cite{Fi99} and \cite{Kata08}, , we obtained the de-reddening parameters for each of our UV/optical/NIR filters (Table \ref{tab:lambda_GX_339}) )."
+ Figure 7. was obtained with the FTS. SMARTS and REM data at four distinct epochs (see caption).," Figure \ref{sed4} was obtained with the FTS, SMARTS and REM data at four distinct epochs (see caption)."
+ The spectra clearly changed between the, The spectra clearly changed between the
+dex.,dex.
+ The second method is linearly fitting the data iu the [Fe/I] diagrams directly for stars from 1<|Z]<3 kpc. where|| ZZ]the contribution from the thin disk aud halo is not significant.," The second method is linearly fitting the data in the $\feh - |Z|$ diagrams directly for stars from $1< |Z| <3$ kpc, where the contribution from the thin disk and halo is not significant."
+ Five eroups in different directions can be separated to test for consistency. aud we found that they eive simular eradieuts of around 0.22540.07 dex spc in four eroups. with the ouly exception for the direction with L<X«IL," Five groups in different directions can be separated to test for consistency, and we found that they give similar gradients of around $-0.225\pm0.07$ dex $^{-1}$ in four groups, with the only exception for the direction with $-4 < X <4$."
+ Both incthods indicate the existence of a 1netallicity eradicut iu the thick disk., Both methods indicate the existence of a metallicity gradient in the thick disk.
+ Iu our opinion. the first method gives a lower eradieut because the mode predictions suppress the existence ofboth metal rich iux metal poor compoucuts in the thick disk.," In our opinion, the first method gives a lower gradient because the model predictions suppress the existence of both metal rich and metal poor components in the thick disk."
+ We thereforc5 favor the second method with a steeper gradieut becaus¢5 it comes froin the data directly aud thus does not depen ou the model assmuption., We therefore favor the second method with a steeper gradient because it comes from the data directly and thus does not depend on the model assumptions.
+ This work has been supported bv the Nationa Natural Science Fouudatiou of China uuder erauts No., This work has been supported by the National Natural Science Foundation of China under grants No.
+ 11073026. 10673015. I0821061. 11078019. the Nationa Basic Research Program of China (973 program) No.," 11073026, 10673015, 10821061, 11078019, the National Basic Research Program of China (973 program) No."
+ 2007CD815102/815103., 2007CB815103/815403.
+ Funding for SDSS-III has been provided by the Alfred P. Sloan Foundation. the Participating Institutions. the National Scicuce Foundation. aud the U.S. Departineut of Encreyv Office of Science.," Funding for SDSS-III has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, and the U.S. Department of Energy Office of Science."
+ The SDSS-HI we site ds lttp://www.sdss3.ore/. SDSS-III Is managed bv the Astrophysical Research Cousortiun for the Participating Iustitutious of the SDSS-III Collaboration iucludius the University of Avizona. the Brazilia Participation ποπ). Brookhaven National Laboratory. Universtv of Cambridge. University of Florida. the5 French Participation Group. the German Participation Croup. the Instituto de Astrofisica de Canarias. tle Michigan State/Notre Dame/JINA Participation Group. Johus IIopkius University. Lawrence Berkeley National Laboratory. Mas Plauck Tustitute for Astroplwsic New Mexico State University. New York Universi Ohio State University. Penusvlvauia. State Universi University of Portsmouth. Princeton University. f Spanish Participation Croup. University of Tolsvo. University of Utah. Vanderbilt University. Universitv Oof Virginia. University of Washington. aud Yale University.," The SDSS-III web site is http://www.sdss3.org/. SDSS-III is managed by the Astrophysical Research Consortium for the Participating Institutions of the SDSS-III Collaboration including the University of Arizona, the Brazilian Participation Group, Brookhaven National Laboratory, University of Cambridge, University of Florida, the French Participation Group, the German Participation Group, the Instituto de Astrofisica de Canarias, the Michigan State/Notre Dame/JINA Participation Group, Johns Hopkins University, Lawrence Berkeley National Laboratory, Max Planck Institute for Astrophysics, New Mexico State University, New York University, Ohio State University, Pennsylvania State University, University of Portsmouth, Princeton University, the Spanish Participation Group, University of Tokyo, University of Utah, Vanderbilt University, University of Virginia, University of Washington, and Yale University."
+Recent vears have seen dramatic increases in the size of the largest cosmological N-body simulations. with giga-particle runs now able to track the non-linear gravitational evolution ofa significant [fraction of the observable universe (Colbere et 1998).,"Recent years have seen dramatic increases in the size of the largest cosmological N-body simulations, with giga-particle runs now able to track the non-linear gravitational evolution of a significant fraction of the observable universe (Colberg et \markcite{C98}
+ 1998)."
+ Correspondinglv. it is becoming possible to make more and more detailed predictions for the non-linear properties of dark matter in any cosmological model. with convergent results being obtained for such features as the profiles ancl correlation functions of dark matter halos (Moore et 1999).," Correspondingly, it is becoming possible to make more and more detailed predictions for the non-linear properties of dark matter in any cosmological model, with convergent results being obtained for such features as the profiles and correlation functions of dark matter halos (Moore et \markcite{M99} 1999)."
+ In. spite of these advances. the one ercal challenge remaining for the simulators is to realistically track the evolution of cosmic gas through its complex eveles of star-Formation and feedback.," In spite of these advances, the one great challenge remaining for the simulators is to realistically track the evolution of cosmic gas through its complex cycles of star-formation and feedback."
+ Given the current state of the art. for. cosmological simulations. rich clusters of galaxies provide one of the most powerful tools with which to make detailed tests of cosmology.," Given the current state of the art for cosmological simulations, rich clusters of galaxies provide one of the most powerful tools with which to make detailed tests of cosmology."
+ Firstly. cluster potential wells are so deep that their formation ancl evolution is largely. unallectec by the uncertain history of the cosmic gas. making clissipationless N-bocly simulations an ideal place in which to study their properties.," Firstly, cluster potential wells are so deep that their formation and evolution is largely unaffected by the uncertain history of the cosmic gas, making dissipationless N-body simulations an ideal place in which to study their properties."
+ Second. clusters correspond to rare peaks in the primordial densitw field. so their statistics are extremely sensitive to detailed features of the cosmology.," Second, clusters correspond to rare peaks in the primordial density field, so their statistics are extremely sensitive to detailed features of the cosmology."
+ In this paper. E will show that a simple model proposed bv Sheth. Mo (1999) can give an extremely eood [it to the correlation function. of massive halos observed. in recent state of the art. numerical simulations.," In this paper, I will show that a simple model proposed by Sheth, Mo \markcite{SMT} (1999) can give an extremely good fit to the correlation function of massive halos observed in recent state of the art numerical simulations."
+ Using this model Lo will compare the predictions of a range of cosmological models. with observations of the cluster correlation length., Using this model I will compare the predictions of a range of cosmological models with observations of the cluster correlation length.
+ Lo particular. L will show that correlation length observations place strong Constraints on the power spectrum of mass fluctuations in the universe. constraints that are almost independent. of other cosmological parameters.," In particular, I will show that correlation length observations place strong constraints on the power spectrum of mass fluctuations in the universe, constraints that are almost independent of other cosmological parameters."
+" In section 2. I describe the Sheth. Alo ""Formen model for the cluster correlation. length. together with a similar semi-analvtic model due to Mo White (1996). and other bias mocdels cliscussecl in. the literature."," In section \ref{sec-model} I describe the Sheth, Mo Tormen model for the cluster correlation length, together with a similar semi-analytic model due to Mo White (1996), and other bias models discussed in the literature."
+ In section 3λ ] compare the predictions of these models with recent numerical simulations. and in section 4 1 show the constraints which can be derived when the models are compared. with observations.," In section \ref{sec-simulations} I compare the predictions of these models with recent numerical simulations, and in section \ref{sec-data} I show the constraints which can be derived when the models are compared with observations."
+ In section 5 ] discuss the relationship of this work to previous studies. and draw my conclusions.," In section \ref{sec-conclusions} I discuss the relationship of this work to previous studies, and draw my conclusions."
+"The primary scientific goal of neutron star (NS) mass and radius measuremeuts are to coustrain the NS equation of state (EOS) and thus internal composition οιος, ueutrons and protous vs. livperous or free quarks).","The primary scientific goal of neutron star (NS) mass and radius measurements are to constrain the NS equation of state (EOS) and thus internal composition (e.g., neutrons and protons vs. hyperons or free quarks)."
+ The NS EOS would be sharply coustrained by a firm determination of a NS inass above 1.6AZ., The NS EOS would be sharply constrained by a firm determination of a NS mass above 1.6.
+.. This would eliminate all EOSs with significaut softening (van WNerkwijk 2001)., This would eliminate all EOSs with significant softening (van Kerkwijk 2001).
+ Dose-Eiusteiu condensation of kaons. livperous. and pions in the core would ereatlv soften the equation of state. leading to NSs which reach a maximum mass of 1.1-1.5 ((Lattimer Prakash 2001).," Bose-Einstein condensation of kaons, hyperons, and pions in the core would greatly soften the equation of state, leading to NSs which reach a maximum mass of 1.4-1.5 (Lattimer Prakash 2001)."
+ Mass determinations of high accuracy have been made for twelve neutron stars in binary svstenus using general relativistic effects. ten of them in double NS systems.," Mass determinations of high accuracy have been made for twelve neutron stars in binary systems using general relativistic effects, ten of them in double NS systems."
+" Thorsett Chakrabarty (1999) fiud that a Gaussian distribution with aand 00.01, wwell describes the masses of these NSs. with none above 1.15AZ."," Thorsett Chakrabarty (1999) find that a Gaussian distribution with and $\sigma$ well describes the masses of these NSs, with none above 1.45."
+.. Tlowever. at least two NSs in accreting X-ray systenis show evidence for higher masses.," However, at least two NSs in accreting X-ray systems show evidence for higher masses."
+ Vela X-1 las a reported mass of 1.564093211. confidence). though svstematic deviations in radial velocity deteriuinations coutinue to make this result uucertain (Barziv et al.," Vela X-1 has a reported mass of $1.86\pm0.32$ confidence), though systematic deviations in radial velocity determinations continue to make this result uncertain (Barziv et al."
+ 2001 and refs therein)., 2001 and refs therein).
+" Cyeuus N-2 has a reported mass of 1.78+02241, ((lo). but this result depends upon the modeling of the secondary and disk light curves (Orosz aud duulkers 1999)."," Cygnus X-2 has a reported mass of $1.78\pm0.23$ $1\sigma$ ), but this result depends upon the modeling of the secondary and disk light curves (Orosz and Kuulkers 1999)."
+ Other lines of evidence also coustrain NS equatious of state., Other lines of evidence also constrain NS equations of state.
+ For iustauce. cooling rates of nuüddle-aged radio pulsars Πρίν a πια auge of masses around 1.1 aand a somewhat stiff EOS (saminker et 22001).," For instance, cooling rates of middle-aged radio pulsars imply a small range of masses around 1.4 and a somewhat stiff EOS (Kaminker et 2001)."
+ Quasi-periodic oscillations in some low-mass X-ray biuaries have soletimes been interpreted (using various models) to indicate niasses near 2 (sco van der Ίνα 2000 for a review). aud to force a radius less than 15.2 kin (vau Straaten ot 22000).," Quasi-periodic oscillations in some low-mass X-ray binaries have sometimes been interpreted (using various models) to indicate masses near 2 (see van der Klis 2000 for a review), and to force a radius less than 15.2 km (van Straaten et 2000)."
+ Also. calculation of the Crab pulsas moment of inertia (Dejeer Tacusel 2002) and modeling of X-ray burst oscillations (Nath ct al.," Also, calculation of the Crab pulsar's moment of inertia (Bejger Haensel 2002) and modeling of X-ray burst oscillations (Nath et al."
+ 2002) suesest that only stiff EOSs are plausible., 2002) suggest that only stiff EOSs are plausible.
+ Receut attempts to fit the X-ray aud. optical spectrum of the isolated neutron star RN J185635-3751 seein to have ruled out all physically well-uucderstoocd atmospheres. so nass and radius calculations for RX J185635-375 remain somewhat tentative (see Burwitz ct al.," Recent attempts to fit the X-ray and optical spectrum of the isolated neutron star RX J185635-3754 seem to have ruled out all physically well-understood atmospheres, so mass and radius calculations for RX J185635-3754 remain somewhat tentative (see Burwitz et al."
+ 2002 aud Ireferences therein)., 2002 and references therein).
+ Temperature anisotropies and relatively high (2> 1016) surface magnetic fields are general features of anodel predictions for pulsar surface| clusion. so transicutly accreting neutron stars (with ich lower magnetic fields. and uceligible teuiperature anisotropies) lay provide a simpler path to determing the NS mass and radius if the thermal component cau be cleauly fit.," Temperature anisotropies and relatively high $B>10^{11}G$ ) surface magnetic fields are general features of model predictions for pulsar surface emission, so transiently accreting neutron stars (with much lower magnetic fields, and negligible temperature anisotropies) may provide a simpler path to determining the NS mass and radius if the thermal component can be cleanly fit."
+ Several NSs that have heen observed in outburst as soft τας trausicuts have also been detected i quiescence (0.5-2.5 keV. N-vay luninosity. Ly.107773 ores 1). for example Cen Αι and Aquila N-1. (vau Paradijs et al.," Several NSs that have been observed in outburst as soft X-ray transients have also been detected in quiescence (0.5-2.5 keV X-ray luminosity, $L_X,\sim10^{32-34}$ ergs $^{-1}$ ), for example Cen X-4 and Aquila X-1 (van Paradijs et al."
+ 1987. Verbunut et al.," 1987, Verbunt et al."
+ 1991)., 1994).
+ See Campana et ((1998a) for a review of soft XN-rav frausients. also kuown as neutron star X-ray novae or quiesceut low-niss N-ray binaries (q(LAINBs}.," See Campana et (1998a) for a review of soft X-ray transients, also known as neutron star X-ray novae or quiescent low-mass X-ray binaries (qLMXBs)."
+ Spectral fits with a soft (kKT-0.2-0.3 keV) blackbody (BB) spectra. usually requiring a hard power-law tail of photon iudex 12. have been acceptable but imply au cussion area of ~1 kin radius. much snaller than appropriate for a NS surface.," Spectral fits with a soft (kT=0.2-0.3 keV) blackbody (BB) spectrum, usually requiring a hard power-law tail of photon index $1-2$, have been acceptable but imply an emission area of $\sim1$ km radius, much smaller than appropriate for a NS surface."
+ However. heavy eleiieuts settle out of the atinosphliere ou a timescale of seconds (Romani 1987). so pure hydrogen atmospheres (assuming accretion has occurred from a nonu- secoudarv) have been used receutlv to calculate," However, heavy elements settle out of the atmosphere on a timescale of seconds (Romani 1987), so pure hydrogen atmospheres (assuming accretion has occurred from a non-helium secondary) have been used recently to calculate"
+understaud to be |1.1].,"understand to be $[-1,1]$."
+ This choice is called an infinite. homogencous IE.S.. It has heen originally investigated by Elton aud Yan [6]..," This choice is called an infinite, homogeneous I.F.S.. It has been originally investigated by Elton and Yan \cite{elton}."
+" Using the maps 0o5,; we can define a stochastic process iu (X=|.1.1] via the following rule: given au initial point &C|1.1]. choose a value of 3 at random in| 1.1]. according to a distribution σσ) (whose support may contain an infinity of points) aud apply 05,5 to map. intoo5,46e)."," Using the maps $\phi_{\delta,\beta}$ we can define a stochastic process in $X=[-1,1]$ via the following rule: given an initial point $x \in [-1,1]$ , choose a value of $\beta$ at random in $[-1,1]$, according to a distribution $\sigma(\beta)$ (whose support may contain an infinity of points) and apply $\phi_{\delta,\beta}$ to map $x$ into$\phi_{\delta,\beta}(x)$."
+ Iterate the procedure., Iterate the procedure.
+ A general theorein. due to Meudivil |185].. guarantees that there exists a unique invariant measure µ for this stochastic process.," A general theorem, due to Mendivil \cite{mendi}, guarantees that there exists a unique invariant measure $\mu$ for this stochastic process."
+" In addition. this measure cau )o found. probability oue. bv the Cesaro average of atomic measures at the soints c; of a trajectory of the process: =ὃν,>pe Tw"," In addition, this measure can be found, probability one, by the Cesaro average of atomic measures at the points $x_j$ of a trajectory of the process: $\frac{1}{n} \sum \delta_{x_j} \rightarrow \mu$."
+a rather pictorial view. we nav describe the process above by saving that the point «is the location of a predator iu a rather peculiarchese of the prey. located at the point Jj: he predator moves towards the prev. but as soon as its distance from 3 is reduced by a factor 8 the prey disappears. to reappear nstautaneously at a new ocation. and the process is repeated.," In a rather pictorial view, we may describe the process above by saying that the point $x$ is the location of a predator in a rather peculiar of the prey, located at the point $\beta$: the predator moves towards the prey, but as soon as its distance from $\beta$ is reduced by a factor $\delta$ the prey disappears, to reappear instantaneously at a new location, and the process is repeated."
+ Therefore. σ is the distribution of the )oxition of the prey. that. along with the value of à vields the distribution ji of the predators positions.," Therefore, $\sigma$ is the distribution of the position of the prey, that, along with the value of $\delta$ yields the distribution $\mu$ of the predator's positions."
+ The measure jp can be equivaleutly defined by. eq. (1)):, The measure $\mu$ can be equivalently defined by eq. \ref{inva2}) ):
+ its approxination properties aud the related inverse problems have been discussed in [3.9.8. 16].," its approximation properties and the related inverse problems have been discussed in \cite{steve,gio1,forte,gio-nal}."
+" The trausfer operator T for infinite. affine homogeneous LE.S. takes the following form: (T f(x} = jJ) foo, (e) = fte | 303)).(1L) where fds any coutmmous function. aud where. to simplify the notation. we have introduced the svinbol 6:—1.6. to be used throughout the paper."," The transfer operator $T$ for infinite, affine homogeneous I.F.S. takes the following form: (T f)(x) = ) f (x) = ) x + ), where $f$ is any continuous function, and where, to simplify the notation, we have introduced the symbol $\bar{\delta}:= 1 - \delta$, to be used throughout the paper."
+ Eq. (1)), Eq. \ref{inva2}) )
+" now beconies: PAT poo foo m dpt foe) = [dptoey ο"".", now becomes: d(T^* )(x) f(x) := (x) (T f)(x) = (x) ) x + ).
+ This is equivaleut to sav that. if ji ds uviuriaut. then [dotethe equality Εμμ ως = μα fce) holds for auy continuous fiction f£.," This is equivalent to say that, if $\mu$ is invariant, then the equality (x) f(x) = (x) (T f)(x) holds for any continuous function $f$."
+ Iu this settiug. our problem cau therefore be formulated bv asking what characteristics of the coustaut 9 and of the nieasure σ iufluence the nature of jp.," In this setting, our problem can therefore be formulated by asking what characteristics of the constant $\delta$ and of the measure $\sigma$ influence the nature of $\mu$."
+ We start attacking this question in the nest section. iu the case when σ is theLebesgue measure.," We start attacking this question in the next section, in the case when $\sigma$ is theLebesgue measure."
+ Iu this section. we consider a case that can be treated analytically as well as nunericallv.," In this section, we consider a case that can be treated analytically as well as numerically."
+It is defined by letting σ be the Lebesque measure on |. 1.1].,"It is defined by letting $\sigma$ be the Lebesque measure on $[-1,1]$ ."
+ By, By
+review).,.
+.. It is reproduced im almost all sources aud outbursts with a striking level of stability., It is reproduced in almost all sources and outbursts with a striking level of stability.
+ There are more than twenty N-ray sources du our ealaxy iu which optical iieasurenmienuts show BIT masses lugher than the theoretical mass lint for a rotating neutron star (sceRemillard&AlcClintock2006.fordetails}..., There are more than twenty X-ray sources in our galaxy in which optical measurements show BH masses higher than the theoretical mass limit for a rotating neutron star \citep[see][ for details]{rm}.
+ These sources are therefore called confined— BUs., These sources are therefore called confirmed BHs.
+ There is also approximately the same nuuber of N-rav binaries for which optical mass measiureuents are unavailable but which otherwise show properties very suilar to the confined DIIs., There is also approximately the same number of X-ray binaries for which optical mass measurements are unavailable but which otherwise show properties very similar to the confirmed BHs.
+ These sources are referred to as DII candidates., These sources are referred to as BH candidates.
+ Better knowledge ofthe DII population is important for uuderstauding some aspects of stellar evolution., Better knowledge of the BH population is important for understanding some aspects of stellar evolution.
+ In this Paper we preseut DIT mass and distance nieasureimoeuts for NTE J1752-225 using the correlation scaling method., In this Paper we present BH mass and distance measurements for XTE J1752-223 using the correlation scaling method.
+ Shaposhuikov&Titarchuk(2009.hereafterSTO9) have shown that correlations diving state transitions can be used to estimate DII lnass and source distance., \citet[][hereafter ST09]{st09} have shown that correlations during state transitions can be used to estimate BH mass and source distance.
+ The scaling method relies ou the theoretically motivated aud observationally tested asstuuption that the QPO frequency for a particular accretion. state of the system is set bv the DIT mass., The scaling method relies on the theoretically motivated and observationally tested assumption that the QPO frequency for a particular accretion state of the system is set by the BH mass.
+ Tuformation on the spectra normalization allows us also to estimate the distance to the source., Information on the spectrum normalization allows us also to estimate the distance to the source.
+ During our monitoring program with we were able to observe a part of the LITS-to-ITSS. transition. which provided suffüiient data to apply the scaling method to ATE J1752-223.," During our monitoring program with we were able to observe a part of the LHS-to-HSS transition, which provided sufficient data to apply the scaling method to XTE J1752-223."
+ The scaling method vields for NTE 223 the DII mass of 9.5 AL. and the distance of about 3.5 kpec., The scaling method yields for XTE J1752-223 the BH mass of $\sim 9.5$ $_\odot$ and the distance of about 3.5 kpc.
+ The main goal ofthis Paper is to present the evolution of X-ray properties diving the NTE J1752-223 discovery outburst observed with aud to report the mass and distance estimates of the ΤΙ��compact object based ou the spectral aud. variability correlation scaling method., The main goal of this Paper is to present the evolution of X-ray properties during the XTE J1752-223 discovery outburst observed with and to report the mass and distance estimates of the compact object based on the spectral and variability correlation scaling method.
+ We analyzed all RNTE observations taken duriug the outburst active phase., We analyzed all RXTE observations taken during the outburst active phase.
+ We follow the properties of the ΟΠΟΙΟΥ spectrum by applying a Comptonization model to the data., We follow the properties of the energy spectrum by applying a Comptonization model to the data.
+ We also characterize the variability properties bv exanüniue periodic and aperiodic features iu the Fourier Power Deusity Spectrum., We also characterize the variability properties by examining periodic and aperiodic features in the Fourier Power Density Spectrum.
+ We also analyze the energv dependence of the variability aud its evolution diving the LIIS-to-IISS state transition., We also analyze the energy dependence of the variability and its evolution during the LHS-to-HSS state transition.
+ Specifically. we find strone evidence that fast (less than 100 sec) aperiodic variability is almost completely confined to he power-law spectral component. while the black ον component iutroduces no significant variability.," Specifically, we find strong evidence that fast (less than 100 sec) aperiodic variability is almost completely confined to the power-law spectral component, while the black body component introduces no significant variability."
+ The observed distribution of variability within the nou-hermal part of the ταν cussion. 1.6. a steep decrease of variability amplitude with energv. has muportaut iuplications for the plwsics of photom up-scatteriug iu he accreting DITs.," The observed distribution of variability within the non-thermal part of the X-ray emission, i.e. a steep decrease of variability amplitude with energy, has important implications for the physics of photon up-scattering in the accreting BHs."
+ The paper is structured as follows., The paper is structured as follows.
+ In the next section we describe the discovery observation with ZXTE/PCA nlee scans., In the next section we describe the discovery observation with /PCA bulge scans.
+ Details of our data analysis are prescuted in 8., Details of our data analysis are presented in \ref{data}.
+ In 81. we present the general evolution of the source as observed by withRATE audSusfft. aud wicover he source evolution through DII spectral states.," In \ref{evolution} we present the general evolution of the source as observed by with and, and uncover the source evolution through BH spectral states."
+ Iu 85 we present analysis of the energv depeudence of the varnabilitv and its evolution with the source spectral states., In \ref{rms_dist} we present analysis of the energy dependence of the variability and its evolution with the source spectral states.
+ Mass measurements and distance estimates withRNTE data using the scaling technique are described iu &6.., Mass measurements and distance estimates with data using the scaling technique are described in \ref{mass}.
+ In 87. woe discuss the source behavior aud possible inuplications of our observational results for the accretion reenucsran in DIIS during different spectral states aud transitions between them., In \ref{discussion} we discuss the source behavior and possible implications of our observational results for the accretion regimes in BHs during different spectral states and transitions between them.
+ Couchisious follow in 88.., Conclusions follow in \ref{summary}.
+ Ou October 23. 2009 17:52 UTRNTE performed a routine scan of the Galactic bulge using the Proportional Counter Array (PCA).," On October 23, 2009 17:52 UT performed a routine scan of the Galactic bulge using the Proportional Counter Array (PCA)."
+ The scan analysis showed residuals consistent with a new source at the position RA = 268.0540.08. DEC = -22.5140.02 (J2000 coordinate system).," The scan analysis showed residuals consistent with a new source at the position RA = $\pm$ 0.08, DEC = $\pm$ 0.02 (J2000 coordinate system)."
+ Search for fast time variability vielded uo significant pulsations., Search for fast time variability yielded no significant pulsations.
+ Ou October 21. 2009Swift BAT instrument trieeered on the source and determined a position consistent with the PCA error circle (ATarkwardtetal. 2009a).," On October 24, 2009 BAT instrument triggered on the source and determined a position consistent with the PCA error circle \citep{atel2258}."
+. On October 25. 06:08 UT performace a dedicated. follow-up observation. which confirmed the previously determined source position aud allowed to estimate an interstellar extinction toward the source of Ny=0.16«1072 7.," On October 25, 06:08 UT performed a dedicated follow-up observation, which confirmed the previously determined source position and allowed to estimate an interstellar extinction toward the source of $N_H=0.46\times 10^{-22}$ $^{-2}$."
+ The observation also revealec a very hard nou-thermal spectrum with a power law index of 1.2 (Alarkwardtetal.2009b)., The observation also revealed a very hard non-thermal spectrum with a power law index of 1.2 \citep{atel2261}.
+. On October 26. 15:03: UT.TE was able to make its first poiute« observation of NTE J1752-223.," On October 26, 15:03 UT was able to make its first pointed observation of XTE J1752-223."
+ Overall spectral anc timing properties are strougly reminiscent of the extreme low-hard states exhibited bv Cye N-1 and GX 330-1. which prompted a preliminary classification of the source asa DII candidate (Shaposhuikovetal.2009).," Overall spectral and timing properties are strongly reminiscent of the extreme low-hard states exhibited by Cyg X-1 and GX 339-4, which prompted a preliminary classification of the source as a BH candidate \citep{atel2269}."
+. Following this identification we trigeered ourZXTE Cycle 11 TOO observations for frequent monitoring of DII trausicut A-rav sources., Following this identification we triggered our Cycle 14 TOO observations for frequent monitoring of BH transient X-ray sources.
+ This program revealed source evolution completely consistent with the source being a BIT candidate., This program revealed source evolution completely consistent with the source being a BH candidate.
+ Namely. the source exhibited typical black hole spectral states; accompanied with corresponding fast variability properties.," Namely, the source exhibited typical black hole spectral states, accompanied with corresponding fast variability properties."
+ As described below. this," As described below, this"
+variants) the information on the physical state in the interior of a region can be represcuted on the xceons boundary and is Buited by the area of this boundary.,variants) the information on the physical state in the interior of a region can be represented on the region's boundary and is limited by the area of this boundary.
+ Two papers iu he influcutial journal[112]. raise the hope that seeiug spacetime-foai effects; and testing quantum eravity theories might not be as forbidding as usually assumed.," Two papers in the influential journal\cite{amelinocamelia} raise the hope that seeing spacetime-foam effects, and testing quantum gravity theories might not be as forbidding as usually assumed."
+ The idea is that there is a number of clifferent mstauces (he neutral kaou system. eanuna ray burst phenomenology. interferometers... yan which prescutly operating nieasuremieut or observation apparata. or apparaa that are going to be soon constructed. involve seusitivitv scales comparable tOo Or not too far from the Planck scale |113]..," The idea is that there is a number of different instances (the neutral kaon system, gamma ray burst phenomenology, interferometers ) in which presently operating measurement or observation apparata, or apparata that are going to be soon constructed, involve sensitivity scales comparable to –or not too far from– the Planck scale \cite{phenomenology}."
+ If this direction fails. τοσαis quantum eravitv might require the investigation of very early cosmology [11n," If this direction fails, testing quantum gravity might require the investigation of very early cosmology \cite{veneziano}."
+ For critical discussions of current direction of research iu quantum gravity. see for instance |115I.," For critical discussions of current direction of research in quantum gravity, see for instance \cite{qg}."
+ The lines of research that I have sumunarized in Section 2 have found many outs of contact in the course of heir deveoplent and have often intersected each other., The lines of research that I have summarized in Section 2 have found many points of contact in the course of their development and have often intersected each other.
+ For iustanece. there is a formal wav of deriving a suni over over ustorics formulation from a canonical theory aud viceversa: the perturbative expansion can also |© Obtained expaucing the sum over listorics: string theory oday faces the prollem of a finding its nouperturbative formulation. and tpus he typical problems of a canonkul theory. while loo) quantuni gravity las uutated into the spiu foam moeels. a sun over lisory formulation. usine echuiques that can be traced to a development of stiine theory of the carly ineties.," For instance, there is a formal way of deriving a sum over over histories formulation from a canonical theory and viceversa; the perturbative expansion can also be obtained expanding the sum over histories; string theory today faces the problem of a finding its nonperturbative formulation, and thus the typical problems of a canonical theory, while loop quantum gravity has mutated into the spin foam models, a sum over history formulation, using techniques that can be traced to a development of string theory of the early nineties."
+" Receutlv. Lce Sinolin has been developing an attempt to connect uonorturbative string theory aud locyp quantum eravity 10,"," Recently, Lee Smolin has been developing an attempt to connect nonperturbative string theory and loop quantum gravity \cite{leemerge}."
+ However. iu spite of this continuous cross fertilization. tlje three main lines of development have kept their esseutial separation.," However, in spite of this continuous cross fertilization, the three main lines of development have kept their essential separation."
+ As pointed out. the three direction of investigation where already clearly identified by Charlκ Misner in 1959 [15|..," As pointed out, the three direction of investigation where already clearly identified by Charles Misner in 1959 \cite{misner}."
+ In. the coicludiuse remark of the qgravitation. ni 1963. Peter Berean roted [117] ," In the concluding remark of the , in 1963, Peter Bergmann noted \cite{pologne}
+ "
+a long-period. sinusoidal TTV can be fitted to the O—C points. with only one significant outlier.,"a long-period, sinusoidal TTV can be fitted to the $O-C$ points, with only one significant outlier."
+ All the observations reported here were made as part of the TASTE (The Asiago Search for Transit timing variations. of Exoplanets) project (?).., All the observations reported here were made as part of the TASTE (The Asiago Search for Transit timing variations of Exoplanets) project \citep{nascimbeni2010}.
+ TASTE is collecting high-precision. short-cadence light curves for a selected sample of transiting exoplanets. to discover low-mass planetary companions or exomoons with the TTV/TDV method (transit time/duration variation).," TASTE is collecting high-precision, short-cadence light curves for a selected sample of transiting exoplanets, to discover low-mass planetary companions or exomoons with the TTV/TDV method (transit time/duration variation)."
+ We refer to that paper for a detailed description of our instrumental setup. observing strategy. and data reduction/analysis.," We refer to that paper for a detailed description of our instrumental setup, observing strategy, and data reduction/analysis."
+ HAT-P-13b is among the sample we are following., HAT-P-13b is among the sample we are following.
+ We collected five transit light curves of HAT-P-13b using the AFOSC imager with its new E2V 42-20 CCD detector mounted at the Asiago 1.82mtelescope!.. An observation log is shown in Table l.., We collected five transit light curves of HAT-P-13b using the AFOSC imager with its new E2V 42-20 CCD detector mounted at the Asiago 1.82m. An observation log is shown in Table \ref{observ}.
+ All the observations were made using a standard Cousins Α filter and 4x binning., All the observations were made using a standard Cousins $R$ filter and $4\times4$ binning.
+" We employed binning and windowing to speed up the readout and decrease as much as possible the technical 7dead"" times between the exposures.", We employed binning and windowing to speed up the readout and decrease as much as possible the technical “dead” times between the exposures.
+ We achieved an average >70% duty-cycle and a « 10s net cadence for all our photometric series., We achieved an average $>70\%$ duty-cycle and a $<10$ s net cadence for all our photometric series.
+ We acquired both sky- and dome flat-field frames during each night: bias and dark frames were taken at both the beginning and the end of a light curve to constrain possible instrumental drifts., We acquired both sky- and dome flat-field frames during each night; bias and dark frames were taken at both the beginning and the end of a light curve to constrain possible instrumental drifts.
+" Stellar profiles were defocused to ~4—6"" FWHM (that 1s. over ~1300 physical pixels) in order to minimize systematic. errors arising from imperfect flat-field correction. guiding drifts. and pixel-to-pixel inhomogeneity."," Stellar profiles were defocused to $\sim4-6''$ FWHM (that is, over $\sim1300$ physical pixels) in order to minimize systematic errors arising from imperfect flat-field correction, guiding drifts, and pixel-to-pixel inhomogeneity."
+" The 9’x2'.6 CCD window that we read included HAT-P-13 às well as the main reference star TYC 3416-1608-1. a star with a magnitude and colour similar to HAT-P-13 (Wy,=10.80 versus (vs.) 10.50 and By—Vy=0.81 vs. 0.52)."," The $9'\times 2'.6$ CCD window that we read included HAT-P-13 as well as the main reference star TYC 3416-1608-1, a star with a magnitude and colour similar to HAT-P-13 $V_T=10.80$ versus (vs.) 10.50 and $B_T-V_T=0.81$ vs. 0.52)."
+ We performed differential aperture photometry on HAT-P-13 using STARSKY (?).. an independent pipeline that we specifically developed for the TASTE project.," We performed differential aperture photometry on HAT-P-13 using STARSKY \citep{nascimbeni2010}, an independent pipeline that we specifically developed for the TASTE project."
+ This code is designed to keep under control any possible source of systematic errors. and implements a fully empirical. iterative approach to identify and correct them.," This code is designed to keep under control any possible source of systematic errors, and implements a fully empirical, iterative approach to identify and correct them."
+ The output light curve is the one with the smallest effective RMS., The output light curve is the one with the smallest effective RMS.
+ Specific diagnostics are evaluated at each iteration to constrain the amount of correlated noise., Specific diagnostics are evaluated at each iteration to constrain the amount of correlated noise.
+ The final. detrended light curves are shown in Fig. 1..," The final, detrended light curves are shown in Fig. \ref{lcurves},"
+ both unbinned and binned over 120 s intervals., both unbinned and binned over 120 s intervals.
+" The photometric RMS scatter is in the range c,=1.7—2.9 mmag for the unbinned points and the range c,»o=0.6—1.1] mmag for the 120 s bins."," The photometric RMS scatter is in the range $\sigma_u
+= 1.7-2.9$ mmag for the unbinned points and the range $\sigma_{120} =
+0.6-1.1$ mmag for the 120 s bins."
+ Three of the light curves in Fig., Three of the light curves in Fig.
+ | represent the most accurate light curves of HAT-P-I3b published so far., 1 represent the most accurate light curves of HAT-P-13b published so far.
+ We ran the JKTEBOP code version 25 (?) to fit a transit model over our light curves., We ran the JKTEBOP code version 25 \citep{southworth2004} to fit a transit model over our light curves.
+ We used a quadratic law for limb darkening. fixing both the linear and the quadratic term i. wo to the theoretical values interpolated from the ? tables. for the stellar parameters of HAT-P-13 derived by ?..," We used a quadratic law for limb darkening, fixing both the linear and the quadratic term $u_1$, $u_2$ to the theoretical values interpolated from the \citet{claret2000}
+ tables, for the stellar parameters of HAT-P-13 derived by \citet{bakos2009}."
+" Three of the remaining parameters of the transit (inclination i. ratio. and sum of the fractional radit ΑιRz. Rk,+ Rp) were estimated by fitting the two highest quality light curves (2011 Jan 3 and Feb 6)."," Three of the remaining parameters of the transit (inclination $i$, ratio, and sum of the fractional radii $R_a/R_b$, $R_a+R_b$ ) were estimated by fitting the two highest quality light curves (2011 Jan 3 and Feb 6)."
+" We then fixed i. R,/R,. R,+Rp to these respective values. and fitted each individual transit only for the central instant 75."," We then fixed $i$, $R_a/R_b$, $R_a+R_b$ to these respective values, and fitted each individual transit only for the central instant $T_0$."
+" Since the formal errors derived by the least squares routine are known to be far too optimistic. we took advantage of two techniques implemented in JKTEBOP to estimate realistic errors: à Monte Carlo test (MC) and a bootstrapping method based on the cyclic permutations of the residuals (RP or ""prayer bead"" algorithm. ?))."," Since the formal errors derived by the least squares routine are known to be far too optimistic, we took advantage of two techniques implemented in JKTEBOP to estimate realistic errors: a Monte Carlo test (MC) and a bootstrapping method based on the cyclic permutations of the residuals (RP or “prayer bead” algorithm, \citealt{southworth2008}) )."
+ The errors from the RP algorithm are significantly larger. suggesting a non-negligible amount of red noise in our light curves.," The errors from the RP algorithm are significantly larger, suggesting a non-negligible amount of red noise in our light curves."
+ We therefore adopted conservatively the RP l-c errors in our analysis., We therefore adopted conservatively the RP $\sigma$ errors in our analysis.
+ The best-fit Zo for each transit. converted from UT to barycentric Julian date (BJD). are shown in Table 2 along with their estimated uncertainties.," The best-fit $T_0$ for each transit, converted from UT to barycentric Julian date (BJD), are shown in Table \ref{timings} along with their estimated uncertainties."
+We performed. a Monte-Carlo simulation similar to those done in Lluetal.(2008). for estimating the measurement errors.,We performed a Monte-Carlo simulation similar to those done in \citet{hu} for estimating the measurement errors.
+ The detail of the procedure is as follows. (, The detail of the procedure is as follows. (
+1)s,i).
+"pectrin, Following VandenBerketal. (2001).. we generated. a composite spectrum using all of our samples."," Following \citet{van}, we generated a composite spectrum using all of our samples."
+ This composite spectrum represents a tvpical quasar spectrum for our samples. (, This composite spectrum represents a typical quasar spectrum for our samples. (
+ii)profiles.,ii).
+ Ne applied the measurement methods written in 3.2 to the composite spectrum. and obtained. typical emission line profiles forV).. aandMeL. (," We applied the measurement methods written in \ref{sec:fit} to the composite spectrum, and obtained typical emission line profiles for, and. ("
+ui)s,iii).
+"pectra, We combined these line profiles with the power-Iow continuum and the Jabmer continuum.", We combined these line profiles with the power-low continuum and the Balmer continuum.
+ Thus the simulated spectrum is written as follows., Thus the simulated spectrum is written as follows.
+ Note that. for simplicity. we ignored the broadening of the ccontinuum.," Note that, for simplicity, we ignored the broadening of the continuum."
+ Values of input. parameters. which are given in the parentheses in equation (15)). are randomly sampled from probability distributions that are mace to reflect the observations.," Values of input parameters, which are given in the parentheses in equation \ref{eq:simflux}) ), are randomly sampled from probability distributions that are made to reflect the observations."
+ Thus. we generated 1.000 simulated: spectra. (," Thus, we generated 1,000 simulated spectra. ("
+iv)template.,iv).
+ Using all the noise spectra produced. by the SDSS pipeline for our samples. we generated a composite spectrum following VandenBerketal.(2001). and named it a noise template.," Using all the noise spectra produced by the SDSS pipeline for our samples, we generated a composite spectrum following \citet{van} and named it a noise template."
+" ""This noise template is sealed so that the resulting median. S/N to be 10 per pixel at the continuum level for cach simulated spectrum. and is treated as its noise."," This noise template is scaled so that the resulting median S/N to be 10 per pixel at the continuum level for each simulated spectrum, and is treated as its noise."
+ Now we have the 1.000 simulated spectra with their noise.," Now we have the 1,000 simulated spectra with their noise."
+ Phe measurement methods written in $3.2 are applied to these simulated spectra., The measurement methods written in \ref{sec:fit} are applied to these simulated spectra.
+" We calculate the value dsj,=(BawsDu)£Py lor cach simulated. spectrum where £75, represents the input parameters. (i.e. the values given in the parentheses in equation (15))) and. {δω represents the corresponding measured. values for the simulated spectra."," We calculate the value $\delta_{sim}=(P_{out} -
+P_{in})/P_{in}$ for each simulated spectrum where $P_{in}$ represents the input parameters (i.e., the values given in the parentheses in equation \ref{eq:simflux}) )) and $P_{out}$ represents the corresponding measured values for the simulated spectra."
+" We regard (0,;,. à standard deviation of 9,5,"" as la error of the measurement."," We regard $\sigma_{sim}$, a standard deviation of $\delta_{sim}$, as $1\sigma$ error of the measurement."
+ Thus we evaluate the measurement errors to be for EW ofUV). for EW of aand for IEM(MgD).," Thus we evaluate the measurement errors to be for EW of, for EW of and for )."
+ “Phe simulation implies the meastwement error to be for the luminosity. which is less than estimated in the power-law fitting.," The simulation implies the measurement error to be for the luminosity, which is less than estimated in the power-law fitting."
+ Therefore we decided to evaluate the measurement error to be for ΑΙ ancl ALzi00-, Therefore we decided to evaluate the measurement error to be for $\lambda L_{3000}$ and $\lambda L_{5100}$.
+ Figure 7 shows the observed celistribution., Figure \ref{fig:fefe_dist} shows the observed distribution.
+ As can be seen from the comparison V)between this and Figure 2.. our photoionization models uncderprecict the bbv a factor of 10. failing to account for the observations.," As can be seen from the comparison between this and Figure \ref{fig:fefe_phot}, our photoionization models underpredict the by a factor of 10, failing to account for the observations."
+ This result. is consistent with the preceding. study by Baldwinetal.(2004)., This result is consistent with the preceding study by \citet{bal04}.
+.. AXdditional microturbulence to the photoionized clouds gets the situation even. worse., Additional microturbulence to the photoionized clouds gets the situation even worse.
+ Thus. Figure 7. seems to challenge classical photoionized pictures of Felr-emitting clouds.," Thus, Figure \ref{fig:fefe_dist} seems to challenge classical photoionized pictures of -emitting clouds."
+ On the other hand. in the case of the collisionally ionized mocdels shown in Figure 3.. the observed {Hux ratios are well. reproduced: with 1077107!em7 and with 6.000«7;10.000 Kk. ‘These results give two remarks: (1) the Felr-emitting louds in quasars are heated το 6.000.<7;«10.000 Ix: (2) the UV and the X-ray photons. which are the heating source in our photoionization models. fail to heat the gas to such temperatures (probably heat the gas too πο).," On the other hand, in the case of the collisionally ionized models shown in Figure \ref{fig:fefe_col}, the observed flux ratios are well reproduced with $10^{22} < N_H
+< 10^{24}\ \mathrm{cm^{-2}}$ and with $6,000 < T_e < 10,000$ K. These results give two remarks: (1) the -emitting clouds in quasars are heated to $6,000 < T_e < 10,000$ K; (2) the UV and the X-ray photons, which are the heating source in our photoionization models, fail to heat the gas to such temperatures (probably heat the gas too hot!)."
+ One possible interpretation is that the t--cmitting clouds are heated by an alternative mechanism such as through shocks., One possible interpretation is that the -emitting clouds are heated by an alternative mechanism such as through shocks.
+ Lere we note that there is a reverberation mapping stucly implying shock heating for comission., Here we note that there is a reverberation mapping study implying shock heating for emission.
+ Ixuehinctal.(2008). analyzed optical comission bands in the Ark 120. finding that they do not respond to the continuum. variation.," \citet{kue} analyzed optical emission bands in the Ark 120, finding that they do not respond to the continuum variation."
+ Thus the optical lelr-emitting region may be heated by other mechanism than photoionization., Thus the optical -emitting region may be heated by other mechanism than photoionization.
+ These results favor the shock heating [or the optical Felr-emitting region. but there are also difficulties.," These results favor the shock heating for the optical -emitting region, but there are also difficulties."
+ First. the amount of shock-processed matter would probably be too laree.," First, the amount of shock-processed matter would probably be too large."
+ Second. as Ixuchnctal.(2008) showed. collisionally ionized. mocdels failed to match the shape of the optical rT-emission band.," Second, as \citet{kue} showed, collisionally ionized models failed to match the shape of the optical -emission band."
+ Fhird. the fact that there is no response to the continuum variation for optical eemission bands can also be interpreted as that the emitting region is too large to vary optical eemission in observable timescales.," Third, the fact that there is no response to the continuum variation for optical emission bands can also be interpreted as that the emitting region is too large to vary optical emission in observable timescales."
+ Unless. shocks are a viable solution. the failure of the photoionization moclel simply indicates that it is not predicting the correct heating rate. or that the radiative transport. calculations are not COLPLDOCl.," Unless shocks are a viable solution, the failure of the photoionization model simply indicates that it is not predicting the correct heating rate, or that the radiative transport calculations are not correct."
+ One possible cause disturbing classical photoionization models to. reproduce. the observations may be the assumption that the eemission is isotropic., One possible cause disturbing classical photoionization models to reproduce the observations may be the assumption that the emission is isotropic.
+ Ferlandοἳal.(2009)— recently suggested that UV. lines are beamecl toward a central source while optical lines are emitted isotropically., \citet{fer09} recently suggested that UV lines are beamed toward a central source while optical lines are emitted isotropically.
+ Phen photoionization models can reproduce the observed. UV to optical flux ratio if the LFeir-emitting clouds arei clistributes asvinmetrically so that we mainly observe their. shieldec aces., Then photoionization models can reproduce the observed UV to optical flux ratio if the -emitting clouds are distributed asymmetrically so that we mainly observe their shielded faces.
+ Llowever. this needs special gcometrical distributions ike Type LL AXGNs: a thick FeIr-emitting gas surrounding he central source with a substantial covering factor. anc intervening between the central region and our eves.," However, this needs special geometrical distributions like Type II AGNs; a thick -emitting gas surrounding the central source with a substantial covering factor, and intervening between the central region and our eyes."
+ At the present time. it is not clear whether or not. photoionization models can reproduce the emission line strengths other than uunder such the situation.," At the present time, it is not clear whether or not photoionization models can reproduce the emission line strengths other than under such the situation."
+ Much. broader. exploration of photoionization model calculations is certainly. needed., Much broader exploration of photoionization model calculations is certainly needed.
+ Now we can roughlv estimate column densities from the observed, Now we can roughly estimate column densities from the observed
+of the nebula decreases as the energy mereases. au effect also reported by??..,"of the nebula decreases as the energy increases, an effect also reported by\cite*{tsk00}."
+" The high spatial and spectral resolution of the lustreation aboard the satellite alows us for the first time to perform spatially resolved svectral analysis of the N-rav nebula associated o 3Cha,", The high spatial and spectral resolution of the instrumentation aboard the satellite allows us for the first time to perform spatially resolved spectral analysis of the X-ray nebula associated to 3C58.
+ To this eud. we extracted spectra from ὃ concentric aunuli with Ar=δ centered at the same xositiou of the N-vay profiles of Fig. 1..," To this end, we extracted spectra from 8 concentric annuli with $\Delta r
+= 8\arcsec$ centered at the same position of the X-ray profiles of Fig. \ref{mos_prof},"
+ covering the core ehuissiou up to roni the center., covering the core emission up to from the center.
+ (ναι the uncertainties i present calibration of the MOS aud PN cameras. and the euerev dependence of the vignetting correction above 5 keV. we ave restricted our spectral analysis to photons iu the 0.55 keV ewrev band.," Given the uncertainties in present calibration of the MOS and PN cameras, and the energy dependence of the vignetting correction above 5 keV, we have restricted our spectral analysis to photons in the 0.5--5 keV energy band."
+ The spectra have been backerouncd subtracted using the same backeround region introduced. in the previous subsection. and have been rebiined to ensure that a mininmuu of 30 counts are present iu each οποιον channel.," The spectra have been background subtracted using the same background region introduced in the previous subsection, and have been rebinned to ensure that a minimum of 30 counts are present in each energy channel."
+ As for response matrices aud effective area files. we have used the latest version of staudiucd MOS awl PN inatrices provided by the calibration teal nuuedinuaallqqetrzarmtttelb.115 aud nunediumi).," As for response matrices and effective area files, we have used the latest version of standard MOS and PN matrices provided by the calibration team 15 and medium)."
+ We have sununued the spectra extracted with MOSL and MOS2. and we have also rescaled the response matrix to reflect flisoperation!.," We have summed the spectra extracted with MOS1 and MOS2, and we have also rescaled the response matrix to reflect this."
+. Besides the 5 aunuli. we have also defined a region (the “edge” region hereafter). represented by the muion of two ellipses⋅ watching⋅ the outer X-ray- edge of. thenebula. nunms a circle with same radius as the ati annulus.," Besides the 8 annuli, we have also defined a region (the ""edge"" region hereafter), represented by the union of two ellipses matching the outer X-ray edge of the, minus a circle with same radius as the 8th annulus."
+ This region is particularly suited for the stuvo of the N-arav Cluission coniüng from the outermost friuges of the X-ray aud radio nebula., This region is particularly suited for the study of the X-ray emission coming from the outermost fringes of the X-ray and radio nebula.
+ We have used three different ciissiou niodels to fif the 3€58 data. nanelv a power-law model. a power-law imnodoel plus the optically thin plasiia model of ?? with Fe L caleulation of ??.. hereafter. aud a power-law model plus a black-body spectre. X1body hereafter.," We have used three different emission models to fit the 3C58 data, namely a power-law model, a power-law model plus the optically thin plasma model of \cite*{mgo85} with Fe L calculation of \cite*{log95}, hereafter, and a power-law model plus a black-body spectrum, pl+bbody hereafter."
+ Thesc +eun=Lardο modes enceonpass what we could possibly expec roni an N-ray nebula. the list two represeutius eventual contributions from a thermal shell (asin the case of known pleriou-comyosite SNR). aud from a conipact source πι he center as poiuted out by ??..," These three models encompass what we could possibly expect from an X-ray nebula, the last two representing eventual contributions from a thermal shell (as in the case of known plerion-composite SNR), and from a compact source in the center as pointed out by \cite*{hbw95}."
+ The eniperature of the niodoel aud of the pl|body nodels lave been constrained iu fιο 0.1LO keV and in he 0.12.0 keV. respectivelv.," The temperature of the model and of the pl+bbody models have been constrained in the 0.1–10 keV and in the 0.1–2.0 keV, respectively."
+ Aliudances are those of TT., Abundances are those of \cite*{ag89}.
+ All the models have been inodified by the interstellar absorption according to cross-sections of 77.. where we rave let the equivalent hydrogen «ο deusity Vy vary.," All the models have been modified by the interstellar absorption according to cross-sections of \cite*{mm83}, where we have let the equivalent hydrogen column density $N_H$ vary."
+" Since we have noted that the nouthermal conipone-ε provides most of the flux iu the bandwidth. aud that the residual thermal couiponeuts of the model ancl plibbocvoare correlated with the value of Nyy. we have also performed a set of fittings ∏⊼∐↓∶↴⋁↑↕∐∖⊀∖∐↖⇁⋜↧↕⋯∖↑∪↓∖↓∩−↓↸⊳⋯−∙↖↖⇁↕∐↸⊳∐↕↴∖↴↸⊳∪∐∏≻⋜↧↑↕↴⋝↕↸∖⋅⋅ - Vo 2 ⋅⋅ ⋅ with previous estimates (?7., 7TT))a xd it is also consistent with the resultwe obtained letting it vary."," Since we have noted that the non-thermal component provides most of the flux in the bandwidth, and that the residual thermal components of the model and pl+bbody are correlated with the value of $N_H$, we have also performed a set of fittings fixing the $N_H$ value to $4\times 10^{21}$ $^{-2}$, which is compatible with previous estimates \cite{hbw95}, \cite{tsk00}) ) and it is also consistent with the resultwe obtained letting it vary."
+ Fie., Fig.
+ 3 shows the best-fit value of the absorption aud of the power-law photon iudex (5) as a function of the distance from reunant center. obtained with a fit to a power-law enission Mode oulv.," \ref{rgamma} shows the best-fit value of the absorption and of the power-law photon index $\gamma$ ) as a function of the distance from remnant center, obtained with a fit to a power-law emission model only."
+ The data clearly show the effect of svuchroron burn-off of lagh energy clectrous as tle radius increases., The data clearly show the effect of synchrotron burn-off of high energy electrons as the radius increases.
+ This effect has also been observed in €21.5-0.9 oth. wih Chaτα C273) aud (?7)) aud it is related to inhomogeneity in the particle cüstribuion inside the pleriou nebula., This effect has also been observed in G21.5-0.9 both with Chandra \cite{scs00}) ) and \cite{wbb00}) ) and it is related to inhomogeneity in the particle distribution inside the plerion nebula.
+ The straight line iu he lower panel of Fie., The straight line in the lower panel of Fig.
+ 3. represents the linear best-fit to he ~ rrolation. =A|Br with A=1.97£0.03 and B=s8.0£0.6«107 aud kis iu avesecouds.," \ref{rgamma}
+ represents the linear best-fit to the $\gamma - r$ relation, $\gamma=A+Br$ with $A=1.97\pm 0.03$ and $B=8.0\pm 0.6\times 10^{-3}$ and $r$ is in arcseconds."
+ The MOS fits o the single power-law reported in Fig., The MOS fits to the single power-law reported in Fig.
+" 39. are statistically from ring Ll to ring 6. while not acceptable in rings 7-8 and in ""the edge” region."," \ref{rgamma} are statistically from ring 1 to ring 6, while not acceptable in rings 7-8 and in “the edge"" region."
+" This is true for thi Ny, free and Nyy fixed fits. aud Table 1. reports the \?/dof values of the fits."," This is true for both $N_H$ free and $N_H$ fixed fits, and Table \ref{gof} reports the $\chi^2/dof$ values of the fits."
+ It is interesting to note that the Nyy valuc. when left free to vary. is significantly lower than he average value of Lx1075 ?. for the outer nebula reelous (Fie. 3)).," It is interesting to note that the $N_H$ value, when left free to vary, is significantly lower than the average value of $4\times 10^{21}$ $^{-2}$ for the outer nebula regions (Fig. \ref{rgamma}) )."
+ Moreover. the best-fit 5 of the shell is off he treud dictated by fit to the spectra of the rugs when Nyy is left free ο vary. while it shows lower deviation iu he ft with Vy; fined.," Moreover, the best-fit $\gamma$ of the shell is off the trend dictated by fit to the spectra of the rings when $N_H$ is left free to vary, while it shows lower deviation in the fit with $N_H$ fixed."
+ If the N-rav. cussion of the outer dines ds donnaed by the non-thermal component of the enron. we do not expect sieUficative variation of the absorption. aud the data seems ο confirm that a fixed Nyy nav be more ao»propriate.," If the X-ray emission of the outer rings is dominated by the non-thermal component of the plerion, we do not expect significative variation of the absorption, and the data seem to confirm that a fixed $N_H$ may be more appropriate."
+" However. the fact that fits of he outer rings and ""edee data are less acceptable than fits to inner rines data strongvo supgeests a contribution youn other com)0neuts."," However, the fact that fits of the outer rings and “edge"" data are less acceptable than fits to inner rings data strongly suggests a contribution from other components."
+ Iu order to assess the preseuce of auy additional ciission from a central compact source (as seeosted by II95) iux from anv thermal shell. we now consider the results of fits with and pl|bbody models.," In order to assess the presence of any additional emission from a central compact source (as suggested by H95) and from any thermal shell, we now consider the results of fits with and pl+bbody models."
+ The inchision of the additional componcut or bbodv) docs no increase the uul] hypothesis probabilitv above54%... excep for ring &.," The inclusion of the additional component or bbody) does not increase the null hypothesis probability above, except for ring 8."
+ This is also reported iu Table 1.. in which the reader fuds tιο data needed to evauate the gooducss of," This is also reported in Table \ref{gof}, , in which the reader finds the data needed to evaluate the goodness of"
+underlying Il-year cycle and several incidences of phase jumps are evident. meaning that periods are not easy to determine.,"underlying 11-year cycle and several incidences of phase jumps are evident, meaning that periods are not easy to determine."
+ The three-year periodicity is not visible in a periodogram of the activity proxy residuals., The three-year periodicity is not visible in a periodogram of the activity proxy residuals.
+ The solar-cycle shifts that are observed in. p-mode frequencies are usually well correlated with proxies of the Sun's activity. such as the 10.7 em radio flux.," The solar-cycle shifts that are observed in $p$ -mode frequencies are usually well correlated with proxies of the Sun's activity, such as the 10.7 cm radio flux."
+ However. in the declining phase of cycle 23 and the current solar minimum we find unusually large differences between the frequencies observed in BiSON data and the activity levels.," However, in the declining phase of cycle 23 and the current solar minimum we find unusually large differences between the frequencies observed in BiSON data and the activity levels."
+ The current cycle minimum indicated by the helioseismic data is significantly deeper than the minima observed by the activity proxies and the structure that is clearly evident in the frequency shifts is not replicated in the proxy data., The current cycle minimum indicated by the helioseismic data is significantly deeper than the minima observed by the activity proxies and the structure that is clearly evident in the frequency shifts is not replicated in the proxy data.
+ We also observe a quasi-biennial signal in the p-mode frequencies at activity levels., We also observe a quasi-biennial signal in the $p$ -mode frequencies at activity levels.
+ Interestingly. this signal is visible at //gi--activity levels in the proxy data.," Interestingly, this signal is visible at -activity levels in the proxy data."
+ As the frequency shifts respond to conditions beneath the surface of the Sun whereas the proxies respond to changes at or above the surface. we suggest that these differences may be caused by changes in the magnetic flux that have yet to manifest at the surface.," As the frequency shifts respond to conditions beneath the surface of the Sun whereas the proxies respond to changes at or above the surface, we suggest that these differences may be caused by changes in the magnetic flux that have yet to manifest at the surface."
+ It is also possible that the magnetic flux responsible for the discrepancies between the frequency shifts and the activity proxies will never reach the Sun's surface., It is also possible that the magnetic flux responsible for the discrepancies between the frequency shifts and the activity proxies will never reach the Sun's surface.
+ The analysis presented here was based on averages made over groups of modes (0€(3. 14€n< 23).," The analysis presented here was based on averages made over groups of modes $0\le\ell\le3$, $14\le n\le 23$ )."
+ Further work on individual modes may allow one to isolate the location of the variability because each mode shows a different sensitivity to the latitudinal distribution of the surface activity., Further work on individual modes may allow one to isolate the location of the variability because each mode shows a different sensitivity to the latitudinal distribution of the surface activity.
+ Such an analysis is currently in progress., Such an analysis is currently in progress.
+ In the seismic data. the previous solar minimum exhibited à double minimum and it appears that the current solar minimum is showing a similar structure.," In the seismic data, the previous solar minimum exhibited a double minimum and it appears that the current solar minimum is showing a similar structure."
+ Furthermore. the most recent p-mode frequencies indicate that the current minimum is still declining.," Furthermore, the most recent $p$ -mode frequencies indicate that the current minimum is still declining."
+ Therefore. it may still be some time before we observe the rising phase of cycle 24.," Therefore, it may still be some time before we observe the rising phase of cycle 24."
+ There have been suggestions that this recent strange behavior of the Sun is indicative that the current grand maximum ts about to end (Abreuetal.2008)., There have been suggestions that this recent strange behavior of the Sun is indicative that the current grand maximum is about to end \citep{Abreu2008}.
+. That would indeed be an occurrence of great significance., That would indeed be an occurrence of great significance.
+ Although our results cannot predict whether this is true they do indicate that the next solar cycle should be observed with great interest., Although our results cannot predict whether this is true they do indicate that the next solar cycle should be observed with great interest.
+ This Letter utilizes data collected by the Birmingham Solar-Oscillations Network (BISON)., This Letter utilizes data collected by the Birmingham Solar-Oscillations Network (BiSON).
+ We thank the members of the BISON team. both past and present. for their technical and analytical support.," We thank the members of the BiSON team, both past and present, for their technical and analytical support."
+ We also thank P. Whitelock and P. Fourie at the South African Astronomical Observatory (SAAO). the Carnegie Institution of Washington. the Australia Telescope National Facility (Australian Commonwealth Scientific and Research Organization. CSIRO). E.J. Rhodes (Mt. Wilson. Californà) and members (past and present) of the Instituto de Astrofisica de Canarias (LAC). Tenerife.," We also thank P. Whitelock and P. Fourie at the South African Astronomical Observatory (SAAO), the Carnegie Institution of Washington, the Australia Telescope National Facility (Australian Commonwealth Scientific and Research Organization, CSIRO), E.J. Rhodes (Mt. Wilson, Californa) and members (past and present) of the Instituto de Astrofisica de Canarias (IAC), Tenerife."
+ BISON 15 funded by the Science and Technology Facilities Council (STFC)., BiSON is funded by the Science and Technology Facilities Council (STFC).
+ The authors also acknowledge the financial support of STFC., The authors also acknowledge the financial support of STFC.
+We recorded all spectra in (wo sections to improve the signal-to-noise ratio (see Figure 3)) and the experimental conditions for each region are listed in Table 2..,We recorded all spectra in two sections to improve the signal-to-noise ratio (see Figure \ref{fig3}) ) and the experimental conditions for each region are listed in Table \ref{tab2}.
+ The two selected regions were a consequence of (he available fillers and constituent materials of the svstem., The two selected regions were a consequence of the available filters and constituent materials of the system.
+ The first region covered the range 740 1690 | and was recorded with thallium bromoiodide (XRS-5) windows. a potassium bromide (IXDr) beamsplitter. a zinc selenide (ZnSe) lens and a liquid nitrogen cooled mercurv-cadimium telluride (MCT) detector.," The first region covered the range 740 – 1690 $^{-1}$ and was recorded with thallium bromoiodide (KRS-5) windows, a potassium bromide (KBr) beamsplitter, a zinc selenide (ZnSe) lens and a liquid nitrogen cooled mercury-cadmium telluride (MCT) detector."
+" The lower wavenumber cut off was limited due to the baud gap of the MCT detector and a small part of the NIL, νο unbrella mode could not be observed.", The lower wavenumber cut off was limited due to the band gap of the MCT detector and a small part of the $_{3}$ $\nu_{2}$ umbrella mode could not be observed.
+are no longer required. (he black hole densities and IGM ionization [fractions are likely to be considerably less.,"are no longer required, the black hole densities and IGM ionization fractions are likely to be considerably less."
+ All these constraints depend heavily on uncertain parameterizations of the efficiency. of star formation and ionizing photon production., All these constraints depend heavily on uncertain parameterizations of the efficiency of star formation and ionizing photon production.
+ However. with more precise measurements of CMB optical depth from future missions. there is hope (hat more stringent constraints on hieh-z star formation and black-hole accretion will be possible.," However, with more precise measurements of CMB optical depth from future missions, there is hope that more stringent constraints on $z$ star formation and black-hole accretion will be possible."
+ We are grateful to David Spergel. Licia Verde. Rachel Bean. and Nick Gnedin for useful discussions regarding the interpretation of WMAP data and numerical simulations.," We are grateful to David Spergel, Licia Verde, Rachel Bean, and Nick Gnedin for useful discussions regarding the interpretation of WMAP data and numerical simulations."
+ We thank Douglas Scott and. Wan Yan Wong for providing their calculations of recombination history., We thank Douglas Scott and Wan Yan Wong for providing their calculations of recombination history.
+ This research at the University of Colorado was supported by astrophysical theory. grants from NASA (NAC5-7262) and NSF (ASTO2-06042)., This research at the University of Colorado was supported by astrophysical theory grants from NASA (NAG5-7262) and NSF (AST02-06042).
+lind. good candidates in the ground-based data of TrES Lyvrl as these classes are characterized by low amplitudes.,find good candidates in the ground-based data of TrES Lyr1 as these classes are characterized by low amplitudes.
+ The classification algorithm was able to detect many &ood candidate class members., The classification algorithm was able to detect many good candidate class members.
+ By candidate we mean a target belonging to the class with the highest class probability above 2 dilferent. cutoll values pri: 0.5 and 0.75 and with a generalized Mahbalanobis distance d«3 to that class., By candidate we mean a target belonging to the class with the highest class probability above 2 different cutoff values $p_{min}$: 0.5 and 0.75 and with a generalized Mahalanobis distance $d<3$ to that class.
+ A quick visual check of the light curves ancl phase plots of the targets with a distance above 3 showed that a large fraction of light. curves sullers from instrumental elfects., A quick visual check of the light curves and phase plots of the targets with a distance above 3 showed that a large fraction of light curves suffers from instrumental effects.
+ The results of the classification are listed in Table 5.., The results of the classification are listed in Table \ref{candidates}.
+ As with ColtoT. the main objective for ος was the search for planets.," As with CoRoT, the main objective for TrES was the search for planets."
+ We do not find many Long Period Variables (LPV). Cepheids and RR Lyr among its targets.," We do not find many Long Period Variables (LPV), Cepheids and RR Lyr among its targets."
+ The total time span of the light curves is also too short to be able to detect. Mira type variables., The total time span of the light curves is also too short to be able to detect Mira type variables.
+ lrrespective of the observed. field. on the sky. we should always find a number of eclipsing binaries and. ellipsoidal variables.," Irrespective of the observed field on the sky, we should always find a number of eclipsing binaries and ellipsoidal variables."
+ Light curves of eclipsing binaries are very cilferent from those of pulsating stars and therefore generally well separated. using the phase dillerences between the first 3 harmonics of the first frequency., Light curves of eclipsing binaries are very different from those of pulsating stars and therefore generally well separated using the phase differences between the first 3 harmonics of the first frequency.
+ Most. detected. candidate binaries have therefore a very high. probability (77 90%) of belonging to the ECL class., Most detected candidate binaries have therefore a very high probability $>90\%$ ) of belonging to the ECL class.
+ We found about 158 reliable eclipsing binaries., We found about 158 reliable eclipsing binaries.
+ Some good examples of eclipsing binary light. curves are shown in Fie. 4., Some good examples of eclipsing binary light curves are shown in Fig. \ref{lc_ecl}.
+ Ht is remarkable that. although eclipses are not always casily seen in the light curve. they clearly show up in the phase plot. ancl are detected by the classification algorithm.," It is remarkable that, although eclipses are not always easily seen in the light curve, they clearly show up in the phase plot and are detected by the classification algorithm."
+ Despite the facet. that Cepheids and IU Lyr are easy. to distinguish from other classes due to their large amplitudes. almost no good candidates were found.," Despite the fact that Cepheids and RR Lyr are easy to distinguish from other classes due to their large amplitudes, almost no good candidates were found."
+ Examples of the few candidates found. are shown in Fig. 5..," Examples of the few candidates found, are shown in Fig. \ref{lc_rp}."
+ As no colour information was available. confusion between 3 Cop and 8 Set stars occurs. because of overlapping frequency ranges.," As no colour information was available, confusion between $\beta$ Cep and $\delta$ Sct stars occurs, because of overlapping frequency ranges."
+ For this reason we merged these 2 classes into a single class., For this reason we merged these 2 classes into a single class.
+ Ht is possible that. for the same target. these classes have similar probabilities below 0.5. but add up to a value well above 0.5.," It is possible that, for the same target, these classes have similar probabilities below 0.5, but add up to a value well above 0.5."
+ Similarly. we could often not make a clear distinction between > Dor and SPB stars. because they show similar eravitv-mocde spectra.," Similarly, we could often not make a clear distinction between $\gamma$ Dor and SPB stars, because they show similar gravity-mode spectra."
+ This problem may be solved bv adding supplementary information like temperature. colours or a spectrum. not only for the targets but also for the training sets.," This problem may be solved by adding supplementary information like temperature, colours or a spectrum, not only for the targets but also for the training sets."
+ Although frequencies around multiples of leefd have been set unreliable. especially the 5. Dor and SPB classes suller from the combination of daily aliasing and instrumental effects.," Although frequencies around multiples of c/d have been set unreliable, especially the $\gamma$ Dor and SPB classes suffer from the combination of daily aliasing and instrumental effects."
+ For this class. a visual inspection," For this class, a visual inspection"
+"There is an interesting history about this effect (e.g., Balbus 2003).","There is an interesting history about this effect (e.g., Balbus 2003)."
+ The instability can be illustrated through an analogy to two masses in a differentially rotating disk that are slightly offset from each other along a radius., The instability can be illustrated through an analogy to two masses in a differentially rotating disk that are slightly offset from each other along a radius.
+" These two masses are connected by a spring, to represent the effect of an axial magnetic field."," These two masses are connected by a spring, to represent the effect of an axial magnetic field."
+ The end result is that the two masses on different orbits seek to increase their displacement from one another., The end result is that the two masses on different orbits seek to increase their displacement from one another.
+ Coupling by the spring leads to a runaway situation ensues., Coupling by the spring leads to a runaway situation ensues.
+ Key for the context of magnetic diagnostics is how this instability impacts the field topology and strength throughout the disk., Key for the context of magnetic diagnostics is how this instability impacts the field topology and strength throughout the disk.
+ The mass-spring analogy above has built into it the assumption of flux freezing., The mass-spring analogy above has built into it the assumption of flux freezing.
+ The separation of the masses along with the differential rotation would appear to evolve the axial field through the disk into a toroidal one., The separation of the masses along with the differential rotation would appear to evolve the axial field through the disk into a toroidal one.
+" Different researchers have studied the operation of the MRI in accretion disks through both semi-analytic work and numerical simulations (e.g., Balbus Hawley 1991, 1992; Hawley Balbus 1991, 1992; Hawley, Gammie, Balbus 1995; Hawley 2000; Fromang Stone 2009; Lesur Longaretti 2009; Maheswaran Cassinelli 2009)."," Different researchers have studied the operation of the MRI in accretion disks through both semi-analytic work and numerical simulations (e.g., Balbus Hawley 1991, 1992; Hawley Balbus 1991, 1992; Hawley, Gammie, Balbus 1995; Hawley 2000; Fromang Stone 2009; Lesur Longaretti 2009; Maheswaran Cassinelli 2009)."
+ The goals of the MRI simulations are to understand better the physics of angular momentum transport through disks and disk structure., The goals of the MRI simulations are to understand better the physics of angular momentum transport through disks and disk structure.
+" However, this paper seeks better insight into whether and how disk magnetism might be directly detected, with a focus on the Zeeman and Hanle effects for spectral lines."," However, this paper seeks better insight into whether and how disk magnetism might be directly detected, with a focus on the Zeeman and Hanle effects for spectral lines."
+ Ignace Gayley (2008) reported on a simplistic calculation of the Zeeman effect and the Hanle effect for a Keplerian disk with a purely toroidal field., Ignace Gayley (2008) reported on a simplistic calculation of the Zeeman effect and the Hanle effect for a Keplerian disk with a purely toroidal field.
+" Here, application of the Hanle effect to Keplerian disks has been developed in a more complete way."," Here, application of the Hanle effect to Keplerian disks has been developed in a more complete way."
+" But before discussing its application to the MRI scenario, it is worth commenting on the conclusions of Ignace Gayley for the use of the Zeeman effect."," But before discussing its application to the MRI scenario, it is worth commenting on the conclusions of Ignace Gayley for the use of the Zeeman effect."
+ 'The MRI leads to a field topology that consists of a toroidal component and a randomized component., The MRI leads to a field topology that consists of a toroidal component and a randomized component.
+" 'The toroidal components exists in mixed polarity, by which it is meant that some sectors run cw and others run ccw."," The toroidal components exists in mixed polarity, by which it is meant that some sectors run cw and others run ccw."
+" For the weak longitudinal Zeeman effect, the Stokes-V flux scales with the net magnetic flux associated with a spatial resolution element (e.g., Mathys 2002)."," For the weak longitudinal Zeeman effect, the Stokes-V flux scales with the net magnetic flux associated with a spatial resolution element (e.g., Mathys 2002)."
+ With bulk motions frequency (or wavelength or velocity) resolution ultimately maps to geometrical zones at the source (an example of this involving the Sobolev approximation was detailed by Gayley Ignace 2010 for spherical winds)., With bulk motions frequency (or wavelength or velocity) resolution ultimately maps to geometrical zones at the source (an example of this involving the Sobolev approximation was detailed by Gayley Ignace 2010 for spherical winds).
+" The key point is that if both the randomized field and the polarity of the toroidal component changes on small scales, then the flux of circularly polarized light arising from the Zeeman effect will be strongly suppressed owing to little net magnetic flux."," The key point is that if both the randomized field and the polarity of the toroidal component changes on small scales, then the flux of circularly polarized light arising from the Zeeman effect will be strongly suppressed owing to little net magnetic flux."
+ The issue of variations of the field on “small scales” must be evaluated with care., The issue of variations of the field on “small scales” must be evaluated with care.
+" In the Sobolev approximation for purely Keplerian rotation, we have seen that the isovelocity zones are “loop” structures."," In the Sobolev approximation for purely Keplerian rotation, we have seen that the isovelocity zones are “loop” structures."
+" These loops extend out to large radius for low velocity shifts, but they can be quite small at high velocity shifts."," These loops extend out to large radius for low velocity shifts, but they can be quite small at high velocity shifts."
+ The highest velocity shifts degenerate to a pair of points for either limb of the star at the projected stellar equator!, The highest velocity shifts degenerate to a pair of points for either limb of the star at the projected stellar equator!
+" More importantly, every isovelocity zone intersects with the innermost radius of the disk. ("," More importantly, every isovelocity zone intersects with the innermost radius of the disk. ("
+"If the disk extended down to the star, this would be the photospheric radius.)","If the disk extended down to the star, this would be the photospheric radius.)"
+ A steep power-law density ensures that the bulk of emission or scattered light comes from inner disk radii., A steep power-law density ensures that the bulk of emission or scattered light comes from inner disk radii.
+" As a result, the Zeeman effect would be most sensitive to a magnetic field at these locations, and thus “small scales"" refers to turnovers in the field direction that are small compared to the inner radius of the disk."," As a result, the Zeeman effect would be most sensitive to a magnetic field at these locations, and thus “small scales” refers to turnovers in the field direction that are small compared to the inner radius of the disk."
+" Consequently, the detection of the Zeeman effect from a disk with a given spectral resolution and density distribution constrains the characteristic spatial wavelengths at which the field turns over with radius from the star, azimuth around the star, and vertical height through the"," Consequently, the detection of the Zeeman effect from a disk with a given spectral resolution and density distribution constrains the characteristic spatial wavelengths at which the field turns over with radius from the star, azimuth around the star, and vertical height through the"
+Ever since powerful computers and digital spectra have become available. there have been efforts to develop algorvitlins for automatic spectral classification (forareviewonthecarlyworkssee ?)..,"Ever since powerful computers and digital spectra have become available, there have been efforts to develop algorithms for automatic spectral classification \citep[for a review on the early works see][]{Kurtz:1984}."
+ The advantages of automated procedures απ compared to manual classification are obvious., The advantages of automated procedures as compared to manual classification are obvious.
+ First of all. ouly a few experts are able to perform accurate manual classifications. aud it was therefore sought to “freeze” this expert knowledge iuto computer programs.," First of all, only a few experts are able to perform accurate manual classifications, and it was therefore sought to “freeze” this expert knowledge into computer programs."
+ Such programs would allow to obtain classifications by criteria. and much larger data sets could be processed thiu by manual classification.," Such programs would allow to obtain classifications by criteria, and much larger data sets could be processed than by manual classification."
+ The latter issue has become ever more demanding. with wpcomime survey missions likeDIVAL..NGST?.. or GAIA?.," The latter issue has become ever more demanding, with upcoming survey missions like, or ."
+. With all these satellites. it is planned to detect millions of objects. or even one billion objects in the case of GAIA.," With all these satellites, it is planned to detect millions of objects, or even one billion objects in the case of GAIA."
+ Tn the last decade. much progress was made in the field of automatic spectral classification. and it was demonstrated that computers are actually capable of verformine this task (forarecent.comprehensivereviewsee 7).," In the last decade, much progress was made in the field of automatic spectral classification, and it was demonstrated that computers are actually capable of performing this task \citep[for a recent, comprehensive review
+see][]{Bailer-Jones:2001}."
+ Using Rurtz οσο distance approach (?7).. ? automatically classified digitized objective prisii spectra roni TWouk’s plates. with good results (0=1.14 ATsvpesh," Using Kurtz' metric distance approach \citep{Kurtz:1984}, \cite{LaSala:1994} automatically classified digitized objective prism spectra from Houk's plates, with good results $\sigma=1.14$ MK-types)."
+ 2 used a simular approach. and applied it to slit spectra with similar spectral resolution iud a slightly arecr wavceleusth coverage (see Tab.," \cite{Penprase:1994} used a similar approach, and applied it to slit spectra with similar spectral resolution and a slightly larger wavelength coverage (see Tab."
+ 1. for a comparison of the data used. aud results obtained).," \ref{performance_comparison} for a comparison of the data used, and results obtained)."
+ The spectral type accuracy he reached for BOFS stars was a bit worse than hat of LaSala: ic... 0—1.5 MIS-tvpes.," The spectral type accuracy he reached for B0–F5 stars was a bit worse than that of LaSala; i.e., $\sigma=1.5$ MK-types."
+" However. as we will see below, it is very difficult to compare the performance of classification algorithms based on the results published iu the literature. because (a) rarely ever is the signal-to-nolse ratio CS7N ) of the data documented. aud the achievable classification accuracy depends critically on S/N: (b) different waveleneth rauges aud spectral resolutions were used: and (ο) the algorithius were applied to stars in differing ranges of spectral type."," However, as we will see below, it is very difficult to compare the performance of classification algorithms based on the results published in the literature, because (a) rarely ever is the signal-to-noise ratio $S/N$ ) of the data documented, and the achievable classification accuracy depends critically on $S/N$; (b) different wavelength ranges and spectral resolutions were used; and (c) the algorithms were applied to stars in differing ranges of spectral type."
+ The influence of the latter on the achievable classification accuracy ds nicely demonstrated by coluparing the results of ? with those of ?.., The influence of the latter on the achievable classification accuracy is nicely demonstrated by comparing the results of \cite{Weaver/Torres-Dodgen:1995} with those of \cite{Weaver/Torres-Dodgen:1997}.
+ Iu the orner paper. the authors report on supervised automatic classification of stars of spectral type AQΑθ with a nulti-laver artificial neural network (ANN) with one iddenu laver. trained with a back-propagation algoritlun.," In the former paper, the authors report on supervised automatic classification of stars of spectral type A0–A9 with a multi-layer artificial neural network (ANN) with one hidden layer, trained with a back-propagation algorithm."
+ They reached a ean absolute deviation of 0.12 spectral vpes and 0.15 Iuuinosity classes., They reached a mean absolute deviation of 0.42 spectral types and 0.15 luminosity classes.
+ In the second paper. the ANN was applied to stars in the range OlM6. and the uean absolute deviations were oulv 1.26 spectral types aud 0:58 LDIunuuinositv classes.," In the second paper, the ANN was applied to stars in the range O4–M6, and the mean absolute deviations were only 1.26 spectral types and 0.38 luminosity classes."
+ The results of Weaver Torres-Dodgen have also shown that spectral classification in the near infrared can be doue with the same accuracy as du the “classical” MIN. spectral range. with spectra of much lower resolution.," The results of Weaver Torres-Dodgen have also shown that spectral classification in the near infrared can be done with the same accuracy as in the “classical” MK spectral range, with spectra of much lower resolution."
+ The resolution used by Weaver Torres-Dodecu was only per pixel. and their spectra ranges950A.," The resolution used by Weaver Torres-Dodgen was only per pixel, and their spectral range."
+. Their results are comparable to that achieved bv others at three times lieher spectra resolution iu the optical or UV., Their results are comparable to that achieved by others at three times higher spectral resolution in the optical or UV.
+ To continue with our brief review. in receut vears. ANNs have been successfully used for supervise automatic spectral classification by a couple of groups.," To continue with our brief review, in recent years, ANNs have been successfully used for supervised automatic spectral classification by a couple of groups."
+ All of them used inultilaver back-propagation networks (AIBPNs)., All of them used multilayer back-propagation networks (MBPNs).
+ ? automatically classified spectra of O3C5 stars obtained with the International Ultraviolet Explorer (UE: dispersion per pixel) with an AIBPN., \cite{Viera/Ponz:1995} automatically classified spectra of O3–G5 stars obtained with the International Ultraviolet Explorer (IUE; dispersion per pixel) with an MBPN.
+ The lo error was 1.11 spectral types., The $\sigma$ error was 1.11 spectral types.
+ They found their ANN classification to be superior to a classification with a metric distance method (0=1.38 types)., They found their ANN classification to be superior to a classification with a metric distance method $\sigma=1.38$ types).
+ The data used by ? were optical, The data used by \cite{Singhetal:1998} were optical
+The Ixuiper belt is a population of panetesimals with cliameters as large as a few thiousaud kilometers (?) auda uass of approximately a few 0.0144 (???)..,"The Kuiper belt is a population of planetesimals with diameters as large as a few thousand kilometers \citep{Brown2006} and a mass of approximately a few $0.01 M_\oplus$ \citep{Gladman2001,Bernstein2004,Fuentes2008}."
+ The stae oltl e belt is enigmatic. as such laree objecds are uot likely to [oru iu sucl a low mass Ixuiper belt. «over the age of the solar system (?)..," The state of the belt is enigmatic, as such large objects are not likely to form in such a low mass Kuiper belt, over the age of the solar system \citep{Kenyon2001}."
+ Raher. it is likely that a touch nore massive initial yell underwent accreion for a time belore some large-scale nass-deyletion event occurred. such as a stellar passage (??).. a sweep through of the nealnr-Diotiou resonances of ej»Xune (?)..?). or the scateriug of ΝΕΟΣ by rogue planets (?)..," Rather, it is likely that a much more massive initial belt underwent accretion for a time before some large-scale mass-depletion event occurred, such as a stellar passage \citep{Ida2000,Levison2004}, a sweep through of the mean-motion resonances of Neptune \citep{Levison2008}, or the scattering of KBOs by rogue planets \citep{Gladman2006}."
+ Whateve ‘the event. it is aso respousibe. at least in part. fo“tl e vnamically excited orbits of NBO populatious such as the 2Ηπο» or the scattering populatlou (?)..," Whatever the event, it is also responsible, at least in part, for the dynamically excited orbits of KBO populations such as the plutinos or the scattering population \citep{Gladman2008}."
+ The size distribution of the Ixuiper belt contaius a fossil record of the end-state of the accretion processes that occurred in that region., The size distribution of the Kuiper belt contains a fossil record of the end-state of the accretion processes that occurred in that region.
+ Ixnowledge «‘the size distriutiou ca1 COSiu disruplon strengths of the bodies. formation time-scales iu tie outer solar system. aud the early conditious of the proto-plauetary disk.," Knowledge of the size distribution can constrain disruption strengths of the bodies, formation time-scales in the outer solar system, and the early conditions of the proto-planetary disk."
+ This inakes the determination of the size clistribiion a primary constraint on Ixuiper belt formation scenarios., This makes the determination of the size distribution a primary constraint on Kuiper belt formation scenarios.
+ Because of the large cistaice to the Ixuiper belt. the size clistribiion s uot determined directly. but rather. is inferred from the shape of the observed luminosity fu1Ct1OL (LF).," Because of the large distance to the Kuiper belt, the size distribution is not determined directly, but rather, is inferred from the shape of the observed luminosity function (LF)."
+ Ealv observatious determined that the LF for bright ojects GnGI)<o26) was well represeuted by a power-law with a logarithinic slope oc0.7 (222222 ..," Early observations determined that the LF for bright objects $(m(R)\lesssim 26)$ was well represented by a power-law with a logarithmic slope $\alpha\sim 0.7$ \citep{Jewitt1998,Gladman1998,Gladman2001,Allen2002,Petit2006,Fraser2008}. ."
+ This suggests that the size clistribition of IKRBOs is a with logarithmic slope q~1.5., This suggests that the size distribution of KBOs is a power-law with logarithmic slope $q\sim4.5$.
+ The deepest survey of ΡΩΣ ou the ecliptic. presented by 2.. Iniid a derthi of faint objects demostrating ihat the LF must “roll-over™ or break to a shallower slope at iC)S27 (?).," The deepest survey of KBOs on the ecliptic, presented by \citet{Bernstein2004}, found a derth of faint objects demonstrating that the LF must “roll-over” or break to a shallower slope at $m(R)\lesssim27$ \citep{Bernstein2004}."
+" Interpretiig the roll-over in terms of the biper belt size clistribution üunplies that the large-object. power-aw size clistribution breaks at a WBO diameter D,~100 kin asstunine albecos.", Interpreting the roll-over in terms of the Kuiper belt size distribution implies that the large-object power-law size distribution breaks at a KBO diameter $D_b\sim 100$ km assuming albedos.
+ Mocels of ace'etion intie outer solar system precicted such a breaς αἱ least au order of magnuitcle sinaller (??)..," Models of accretion in the outer solar system predicted such a break at least an order of magnitude smaller \citep{Kenyon2001,Kenyon2002}."
+ The goal of or work was to coulirm tle results of 2.. aud accurately measure the shape ol the LF at and beyoud the roll-over maguitude.," The goal of our work was to confirm the results of \citet{Bernstein2004}, and accurately measure the shape of the LF at and beyond the roll-over magnitude."
+ Froin this. the shape of the under-Iyiug IKBO size distribution ea be inferred.," From this, the shape of the under-lying KBO size distribution can be inferred."
+ We present |ere a survey of the ecliptic with limiting magnitude m(R)~27. auc use this alone with previous observations to cletermite the shape of the LF.," We present here a survey of the ecliptic with limiting magnitude $m(R)\sim 27$, and use this along with previous observations to determine the shape of the LF."
+ Iu section 2 we presel our observations and in sectious 3 we present our survey results., In section 2 we present our observations and in sections 3 we present our survey results.
+ La sectious ft and 5 we couside| past ecliptic Ixuiper belt surveys aud cousider various LF fuctional formes., In sections 4 and 5 we consider past ecliptic Kuiper belt surveys and consider various LF functional forms.
+ In sectious 6 aud 7 we present our analysis aud a discussion of the results., In sections 6 and 7 we present our analysis and a discussion of the results.
+ New observations were mace in service mode with Suprime-C'aim on the Subaru &.2 in telescope (?).., New observations were made in service mode with Suprime-Cam on the Subaru 8.2 m telescope \citep{Miyazaki2002}.
+" Suprime-cam is a 10 chip mosaic CCD camera with a 30x27"" field-of-view (FOV) and a 07.202 pixel scale.", Suprime-cam is a 10 chip mosaic CCD camera with a $\times$ 27' field-of-view (FOV) and a $0''.202$ pixel scale.
+ Observations were made in the r-band on the nights of April 22nd (night 1). audMay," Observations were made in the r-band on the nights of April 22nd (night 1), andMay"
+Type LEP supernovae (SNe) are core-collapse (CC?) induced. explosions that eject massive hyerogen-rich envelope.,Type II-P supernovae (SNe) are core-collapse (CC) induced explosions that eject massive hydrogen-rich envelope.
+ These. SNe show a nearly constant-Iuminosity plateau in their light curve., These SNe show a nearly constant-luminosity plateau in their light curve.
+ Phe plateau phase lasts for ~ 100 days. ancl finally ends up in a rapid decline of the Luminosity. as the expanding ejecta becomes fully transparent.," The plateau phase lasts for $\sim$ 100 days, and finally ends up in a rapid decline of the luminosity as the expanding ejecta becomes fully transparent."
+ CC SNe and their progenitors play important role in understanding the evolution. of the most. massive stars., CC SNe and their progenitors play important role in understanding the evolution of the most massive stars.
+ Extensive searches. for SN. progenitors showed. that the progenitors of Type - SNe are stars with initial masses between  S and 17 M. (Smarttctal.2009).., Extensive searches for SN progenitors showed that the progenitors of Type II-P SNe are stars with initial masses between $\sim$ 8 and 17 $_\odot$ \citep{smarttprog}.
+. Phe lower value is roughly in agreement with the prediction of current theoretical evolutionary tracks. but. the upper limit. is somewhat lower than expected.," The lower value is roughly in agreement with the prediction of current theoretical evolutionary tracks, but the upper limit is somewhat lower than expected."
+ Phe fate of the more massive (M & 20 AL.) stars is controversial., The fate of the more massive (M $\geq$ 20 $_\odot$ ) stars is controversial.
+ Phey are thought to be he input channel for Tvpe Ib/c SNe (Gaskelletal1986).. oit this has not been confirmed directly by observations vet (Smartt2009).," They are thought to be the input channel for Type Ib/c SNe \citep{gaskell86}, but this has not been confirmed directly by observations yet \citep{smartt_rev}."
+. Any observational studs of the physies of SNe and heir progenitors (and. obviously. all astronomical objects) depends heavily on the knowledge of their distances.," Any observational study of the physics of SNe and their progenitors (and, obviously, all astronomical objects) depends heavily on the knowledge of their distances."
+ Because distance is such a crucial parameter. strong elforts have en devoted to the development. of distance measurement echniques that are applicable for Evpe L-P SNe.," Because distance is such a crucial parameter, strong efforts have been devoted to the development of distance measurement techniques that are applicable for Type II-P SNe."
+ The Expanding Photosphere Method (IPM) (IxirshnerIxwan 1974).. a variant of the Daade-Messelink method. uses the idea of comparing the angular and the physical cliameter of the expanding SN ejecta.," The Expanding Photosphere Method (EPM) \citep{kirshner_epm}, a variant of the Baade-Wesselink method, uses the idea of comparing the angular and the physical diameter of the expanding SN ejecta."
+ The input quantities are Huxes. temperatures ancl velocities observed at several epochs during the photospherie phase.," The input quantities are fluxes, temperatures and velocities observed at several epochs during the photospheric phase."
+ Thus. IPM requires both shotometric ancl spectroscopic monitoring of SNe throughout the plateau phase.," Thus, EPM requires both photometric and spectroscopic monitoring of SNe throughout the plateau phase."
+ A clear advantage of LPAI is that it does not require external calibration via SNe with known clistances., A clear advantage of EPM is that it does not require external calibration via SNe with known distances.
+ Another interesting property is that EPAL is much less sensitive to uncertainties in the interstellar reddening and absorption toward SNe., Another interesting property is that EPM is much less sensitive to uncertainties in the interstellar reddening and absorption toward SNe.
+ As, As
+Eclipsing binaries are systems in. which the masses. radi. temperatures. and absolute luminosities can be measured for two stars with the same age and metallicity (Andersen.J..1991:Youngetal.. 2001).,"Eclipsing binaries are systems in which the masses, radii, temperatures, and absolute luminosities can be measured for two stars with the same age and metallicity \citep{andersen91,youngarnett01}."
+. In addition. apsidal motion measurements. if available. can further constrain interior. stellar. physics (Young&Arnett.2005:Schwarzschild.M..1957).," In addition, apsidal motion measurements, if available, can further constrain interior stellar physics \citep{youngarnett05,schwarz57}."
+. Therefore. these objects are effective tools used extensively for calibrating and testing theoretical stellar evolution models.," Therefore, these objects are effective tools used extensively for calibrating and testing theoretical stellar evolution models."
+ Pre-main sequence eclipsing binaries are. particularly important., Pre-main sequence eclipsing binaries are particularly important.
+ Young stars are changing rapidly as they contract onto the main sequence and heat up. so an individual solar-type star covers a large range of temperatures and radii in its first 50 Myr.," Young stars are changing rapidly as they contract onto the main sequence and heat up, so an individual solar-type star covers a large range of temperatures and radii in its first 50 Myr."
+ Furthermore. young stars often have star spots. coronal activity. and/or circumstellar dust which complicates the physics of their formation and early evolution.," Furthermore, young stars often have star spots, coronal activity, and/or circumstellar dust which complicates the physics of their formation and early evolution."
+ Thus. empirical measurements of fundamental properties of young stars with a range of masses and ages are needed to fully constrain the pre-main sequence evolution and understand how the properties of young stars are affected by other parameters (e.g. magnetic fields. dust. metallicity. planets. multiplicity).," Thus, empirical measurements of fundamental properties of young stars with a range of masses and ages are needed to fully constrain the pre-main sequence evolution and understand how the properties of young stars are affected by other parameters (e.g. magnetic fields, dust, metallicity, planets, multiplicity)."
+ Pre-main sequence eclipsing binaries are the only objects which can provide precise measurements of the most fundamental physical properties of young stars. however these objects are extremely rare.," Pre-main sequence eclipsing binaries are the only objects which can provide precise measurements of the most fundamental physical properties of young stars, however these objects are extremely rare."
+ There are only 6 known low-mass pre-main sequence eclipsing binaries with M.-1.5 M: RXJ 0529.4+0041A (Covinoetal..2000.2004).. VI174 Ori (Stassunetal. 2004). 2MASS J05352184-20130546085 (Stassunetal..2006. 2007).. JW 380 (Irwinetal..2007). Par 1802 (Cargileetal...2008:Stassunal..2008).. and ASAS J052821+0338.5 (Stempelsetal..2008).," There are only 6 known low-mass pre-main sequence eclipsing binaries with $M<1.5$ $_{\odot}$: RXJ 0529.4+0041A \citep{covino00,covino04}, V1174 Ori \citep{stassun04}, 2MASS J05352184-20130546085 \citep{stassun06,stassun07}, , JW 380 \citep{irwin07}, Par 1802 \citep{cargile,stassun08}, and ASAS J052821+0338.5 \citep{stempels08}."
+. In addition. EK Cep (Popper.D.M..1987).. TY CrA(Caseyetal..1995).. and RS Cha (Alecianetal..2005.2007) are known to be higher mass eclipsing systems with at least one PMS component.," In addition, EK Cep \citep{popper87}, TY \citep{casey98}, and RS Cha \citep{rscha1,rscha2} are known to be higher mass eclipsing systems with at least one PMS component."
+ In this paper. we announce the discovery of à new low-mass PMS eclipsing binary. and the first such object discovered outside of the Orion star forming region.," In this paper, we announce the discovery of a new low-mass PMS eclipsing binary, and the first such object discovered outside of the Orion star forming region."
+ We describe the observations that were used to identify the object as a PMS eclipsing binary refsec:observations))., We describe the observations that were used to identify the object as a PMS eclipsing binary \\ref{sec:observations}) ).
+ Using the discovery data. we determine initial orbital parameters for the system and approximate physical properties of its component stars refsec:results)).," Using the discovery data, we determine initial orbital parameters for the system and approximate physical properties of its component stars \\ref{sec:results}) )."
+ Finally. we discuss future observations needed to fully analyse this important system refsec:summary)).," Finally, we discuss future observations needed to fully analyse this important system \\ref{sec:summary}) )."
+ TYC 7310-503-] ία =14:58:37.7. 6 2-35:40:30.4). MML 53. is a late-type star spatially located in the region of the Scorpius Centarus OB association complex.," TYC 7310-503-1 $\alpha=$ 14:58:37.7, $\delta=$ -35:40:30.4), MML 53, is a late-type star spatially located in the region of the Scorpius Centarus OB association complex."
+ It was identifed as an X-ray source by ROSAT (RXJI458.6-3541) which lead to further spectroscopic study by several groups with the aim of identifying young stars and mapping the overall star forming region (Wichmannetal..1997.b:Mamajek2002:Torresetal..2006;White 2007).," It was identifed as an X-ray source by ROSAT (RXJ1458.6-3541) which lead to further spectroscopic study by several groups with the aim of identifying young stars and mapping the overall star forming region \citep{Wichmann,Wichmannb,Mamajek,Torres,White}."
+. The object was first identified as a T Tauri star (K3-type) through the detection of Li 1.16708 absorption by Wichmannetal., The object was first identified as a T Tauri star (K3-type) through the detection of Li $\lambda6708$ absorption by \citet{Wichmannb}.
+(1997b).. Ha emission and significant lithium absorption were detected in all subsequent spectroscopic observations. therefore MML 53 is confirmed to be a young.pre-mainsequence (PMS) object.," $\alpha$ emission and significant lithium absorption were detected in all subsequent spectroscopic observations, therefore MML 53 is confirmed to be a young,pre-mainsequence (PMS) object."
+ Mamajeketal.(2002) also kinematically and spatially defined MML 53 as a member of the 15-22 Myr old, \citet{Mamajek} also kinematically and spatially defined MML 53 as a member of the 15-22 Myr old
+"ealaxies and they have color dispersions larger than those in carly-type ealaxies,",galaxies and they have color dispersions larger than those in early-type galaxies.
+ The pCMD shape of NCC 5191 is approximately consisteut with the late-type galaxy pCAIDs described by Lauvou-Fosteretal.(2007).. although their details arecly depend on the individual characteristics of ate-tvpoe ealaxies.," The pCMD shape of NGC 5194 is approximately consistent with the late-type galaxy pCMDs described by \citet{lan07}, although their details largely depend on the individual characteristics of late-type galaxies."
+ Since the red and blue pixel sequences are individually defined iu the BV aud VI colors. he pixels iu those sequences have slightly differeut ueanines according to the criterion color.," Since the red and blue pixel sequences are individually defined in the $B-V$ and $V-I$ colors, the pixels in those sequences have slightly different meanings according to the criterion color."
+ The bV red sequence pixels cousist of mostly P1 jxels aud some P2. P3 aud P5 pixels. while the V Pred sequence pixels consist of 1nostly P2 aud Pil pixels and a sinall fraction of P3 aud P5 pixels.," The $B-V$ red sequence pixels consist of mostly P4 pixels and some P2, P3 and P5 pixels, while the $V-I$ red sequence pixels consist of mostly P2 and P4 pixels and a small fraction of P3 and P5 pixels."
+" However. since the wmmber of P2 and P5 pixels is much smaller than that of Pl pixels (fp,>NU at Ro <90"" and fp,>SOW at Ro <1307: sce Fie. 11))."," However, since the number of P2 and P5 pixels is much smaller than that of P4 pixels $f_{P4}>80\%$ at R $<90''$ and $f_{P4}>50\%$ at R $<130''$; see Fig. \ref{radfrac}) ),"
+ the dominant pixel population iu the red pixel sequence is the P1. for both 5.V aud V1 colors.," the dominant pixel population in the red pixel sequence is the P4, for both $B-V$ and $V-I$ colors."
+ Thus. if we suppose solar metallicity (Z = 0.02) and moderate dust extinction τι= 1). the huninositv-welehltec mean stellar ages of the red sequence pixels are iiostlv 2L Gyr.," Thus, if we suppose solar metallicity (Z = 0.02) and moderate dust extinction $\tau_V=1$ ), the luminosity-weighted mean stellar ages of the red sequence pixels are mostly $\gtrsim1$ Gyr."
+ However. sole pixels iu the pluie features (particularly in the WF pCMDJ) iiiv © Very voune (<10 Myr) and dustv populations. as discussecl iu refpopvartent..," However, some pixels in the plume features (particularly in the $V-I$ pCMD) may be very young $<10$ Myr) and dusty populations, as discussed in \\ref{popvartext}."
+ There are two major factors affecting the pCMD red pixel sequences., There are two major factors affecting the pCMD red pixel sequences.
+ First. since pixels closer to the center of a galaxy tend to be brighter. a red pixel sequence reflects the color eradicut along the radius of a galaxy (or a galaxy bulec).," First, since pixels closer to the center of a galaxy tend to be brighter, a red pixel sequence reflects the color gradient along the radius of a galaxy (or a galaxy bulge)."
+ It is well mown that the main origin of the color eradicnut of bulees or elliptical galaxies is reir (c.g.Iwo&Th2005: 2010).. which is an maportaut factor determining I shape of the red pixel sequence.," It is well known that the main origin of the color gradient of bulges or elliptical galaxies is their \citep[e.g.][]{ko05,lab09,tor10}, which is an important factor determining the shape of the red pixel sequence."
+ When it is supposed that the red pixel sequence slope is surely determined by metallicity variation. the estimated metallicity difference between V220 nag 7 aud V—17 mae 7 is as aree as ΑΙΓΟΠΠ ~2 (from Z z0.0002 to 7 zz0.02: with mean stellar age = 3 Gyr aud τι.=ld oassunpptions).," When it is supposed that the red pixel sequence slope is purely determined by metallicity variation, the estimated metallicity difference between $V=20$ mag $^{-2}$ and $V=17$ mag $^{-2}$ is as large as $\Delta$ [Fe/H] $\sim2$ (from Z $\approx0.0002$ to Z $\approx0.02$; with mean stellar age = 3 Gyr and $\tau_V=1$ assumptions)."
+ However. since the typical uetallicity eradicuts of early-type galaxies (of which properties are VOYV simular to those of lees) are sinaller than 0.6 (Spolaoretal.2010).. his estimated metallicity difference is too large.," However, since the typical metallicity gradients of early-type galaxies (of which properties are very similar to those of bulges) are smaller than 0.6 \citep{spo10}, this estimated metallicity difference is too large."
+ The second factor is internalattenuation., The second factor is internal.
+ As reported x Lauvou-Fosteretal.(2007).. ποιο elliptical galaxies show hook features iu their red pixel sequences. which scems to be affected by dust.," As reported by \citet{lan07}, some elliptical galaxies show hook features in their red pixel sequences, which seems to be affected by dust."
+ Since dust makes optical colors redder. the shape of a red pixel sequeuce is distorted according to the distribution of dust in a galaxwv.," Since dust makes optical colors redder, the shape of a red pixel sequence is distorted according to the distribution of dust in a galaxy."
+ When it is supposed that the red pixel sequence slope js purely determined by dust attenuation. the estimated optical depth differeuce by dust between Έ520 nae 5 and V...—17 mae 2 is as huge as Any~L corresponding to AEQΤο0.2.," When it is supposed that the red pixel sequence slope is purely determined by dust attenuation, the estimated optical depth difference by dust between $V=20$ mag $^{-2}$ and $V=17$ mag $^{-2}$ is as large as $\Delta\tau_V\sim4$, corresponding to $\Delta E(B-V)\sim0.2$ ."
+ However. the actual red pixel sequence iu the pCCD does not exactly aeree with ατα: the sequence of the metallicity effect or that of the dust effect.," However, the actual red pixel sequence in the pCCD does not exactly agree with either the sequence of the metallicity effect or that of the dust effect."
+ Thus. the red pixel sequence is thought to be affected by both metallicity and dust variations. and in addition. the effect of age variation may not be zero.," Thus, the red pixel sequence is thought to be affected by both metallicity and dust variations, and in addition, the effect of age variation may not be zero."
+ As shown iu Lauvon-redhook. features). the effect of dust on the rec pixel sequence slope sccius to be different from that of metallicity.," As shown in \citet[][see their Section 3.3.1 on the \emph{red hook} , the effect of dust on the red pixel sequence slope seems to be different from that of metallicity."
+ Thus. if pCAID data for a lavee enough nuuber Guore than several teus) of ealaxies with incdependcutly estimated metallicity auc dust information are secured. we can establish a statistical method to distinguish the effects of motallicity and dust on the red pixel sequence slope and color offset.," Thus, if pCMD data for a large enough number (more than several tens) of galaxies with independently estimated metallicity and dust information are secured, we can establish a statistical method to distinguish the effects of metallicity and dust on the red pixel sequence slope and color offset."
+ After that. iuverselv. i will be possible to estimate the approximate motallicity aud dust content of ealaxies ouly using their pCAIDs.," After that, inversely, it will be possible to estimate the approximate metallicity and dust content of galaxies only using their pCMDs."
+ As displaved in Fie., As displayed in Fig.
+ Saa. the blue pixel sequence reflects the properties of the disk and spiral arms: the star-forming area with voung stars.," \ref{cmap}a a, the blue pixel sequence reflects the properties of the disk and spiral arms: the star-forming area with young stars."
+ Like the red pixel sequence. the blue pixel sequence has slightly different meanines according to the criterion color.," Like the red pixel sequence, the blue pixel sequence has slightly different meanings according to the criterion color."
+ The pixels in the BoV. blue pixel sequence cousist of P1. P2 and P3 pixels. while the pixels in the WO£ blue pixel sequence cousist of P1. P3 and P5 pixels.," The pixels in the $B-V$ blue pixel sequence consist of P1, P2 and P3 pixels, while the pixels in the $V-I$ blue pixel sequence consist of P1, P3 and P5 pixels."
+ For those populations. if solar inetallicitv (Z = 0.02) aud moderate dust extinction (7= 1) are supposed. the Iuninositv- mean stellar ages of the blue sequence pixels are mostly =1 Cyr.," For those populations, if solar metallicity (Z = 0.02) and moderate dust extinction $\tau_V=1$ ) are supposed, the luminosity-weighted mean stellar ages of the blue sequence pixels are mostly $\lesssim1$ Gyr."
+ When it is supposed that the color variation in the blue pixel sequence is purely due to age variation. the estimated," When it is supposed that the color variation in the blue pixel sequence is purely due to age variation, the estimated"
+l.,\ref{table:eos}.
+ The maximum accreted mass for which the iterative solver converges reliably (grid-averaged: residual x:5 per cent) depends on theEOS. with Muuscm1o107.3107.2.*.3.PALS for models AD. respectively. (," The maximum accreted mass for which the iterative solver converges reliably (grid-averaged residual $\leq 5$ per cent) depends on theEOS, with $M_{\mathrm{a},\mathrm{max}} \approx 1\times10^{-3}, \
+3\times10^{-8}, \ 2\times10^{-7}, \ 3\times10^{-6} \mathrm{M}_{\sun}$ for models A–D, respectively. ("
+"Asà corollary. the gradient did, in the vicinity of the rightmost data point for each model in Fig.","Asa corollary, the gradient $\mathrm{d}\mu/\mathrm{d}M_{\mathrm{a}}$ in the vicinity of the rightmost data point for each model in Fig."
+ 2 is unphysically steep.), \ref{fig:dipole_mass} is unphysically steep.)
+" The method we use to caleulate the dipole moment cdilfers sliehtly from that in PMÓO4: we integrate c directly rather than /3,.. according to [or the (th multipole moment. circumventing one set of numerical derivatives and improving the accuracy of the results."," The method we use to calculate the dipole moment differs slightly from that in PM04; we integrate $\psi$ directly rather than $B_{r}$, according to for the th multipole moment, circumventing one set of numerical derivatives and improving the accuracy of the results."
+ Equation (14)) is 10 per cent more accurate than equation (34) in PPMOA for a 64 erid., Equation \ref{multipole_moments}) ) is $\sim 10$ per cent more accurate than equation (34) in PM04 for a $64 \times 64$ grid.
+ The discrepancy shrinks to «l per cent for a 256.256 grid., The discrepancy shrinks to $ < 1$ per cent for a $256 \times 256$ grid.
+ lt dis clear from Fig., It is clear from Fig.
+ 2 int the characteristic mass AM required to significantly distort the initial configuration varies with the EOS., \ref{fig:dipole_mass} that the characteristic mass $M_{\mathrm{c}}$ required to significantly distort the initial configuration varies with the EOS.
+ Lowe define Al. to be the accreted mass that halves ji from its initial value fn. to be consistent with the empirical scaling introduced by 2.. viz.," If we define $M_{\mathrm{c}}$ to be the accreted mass that halves $\mu$ from its initial value $\mu_{\mathrm{i}}$ , to be consistent with the empirical scaling introduced by \citet{shibazaki1989}, viz."
+ then Fig., then Fig.
+" 2. vields Alo,&5)10UALS. Mogm3 “Me. Moez1210""M and Mop22510PALS for models AD in Table 1.."," \ref{fig:dipole_mass} yields $M_{\mathrm{c},\mathrm{A}} \approx 5 \times
+10^{-4} \mathrm{M}_{\sun}$, $M_{\mathrm{c},\mathrm{B}} \approx 3 \times 10^{-8} \mathrm{M}_{\sun}$ $M_{\mathrm{c},\mathrm{C}} \approx 1 \times 10^{-7} \mathrm{M}_{\sun}$ and $M_{\mathrm{c},\mathrm{D}}
+\approx 2 \times 10^{-6} \mathrm{M}_{\sun}$ for models A–D in Table \ref{table:eos}."
+ Plainly. varving the EOS makes a big dilference.," Plainly, varying the EOS makes a big difference."
+ AZ. is reduced by a factor of between 310? (model D) and 2.107 (mocel B) relative to an isothermal mountain., $M_{\mathrm{c}}$ is reduced by a factor of between $3 \times 10^{2}$ (model D) and $2 \times 10^{4}$ (model B) relative to an isothermal mountain.
+" “Phis is because adiabatie mountains are up to ~Q0? times taller than isothermal ones for AZ,=Ale (see Fie."," This is because adiabatic mountains are up to $\sim
+10^{2}$ times taller than isothermal ones for $M_{\mathrm{a}} = M_{\mathrm{c}}$ (see Fig."
+e 3 below and Section 4.5))., \ref{fig:dipole_radius} below and Section \ref{section_4:hydromagnetic_structure}) ).
+ At higher altitudes. the magnetic stress Cxrr 7) ds weaker and hence the pressure eracient pushes the magnetic field sideways more than in an isothermal mountain.," At higher altitudes, the magnetic stress $\propto r^{-6}$ ) is weaker and hence the pressure gradient pushes the magnetic field sideways more than in an isothermal mountain."
+ In the limit of small Λι one can show (see Appendix D) that the scaling. of the characteristic mass AM. [or adiabatic mountains is proportional to the square of he magnetic field. strength. as lor isothermal magnetic mountains (see Section 3.2 in PMOOA).," In the limit of small $M_{\mathrm{a}}$, one can show (see Appendix \ref{appendix:gs_analytic}) ) that the scaling of the characteristic mass $M_{\mathrm{c}}$ for adiabatic mountains is proportional to the square of the magnetic field strength, as for isothermal magnetic mountains (see Section 3.2 in PM04)."
+" Aclitionally. AJ, is also inversely proportional to an extra factor Z6Xo.D) (evaluated as a contour plot on the Xo E plane in Fig. DI1))."," Additionally, $M_{\mathrm{c}}$ is also inversely proportional to an extra factor $I(\Lambda_{0}, \Gamma)$ (evaluated as a contour plot on the $\Lambda_{0}$ $\Gamma$ plane in Fig. \ref{fig:lambda_gamma}) ),"
+ which depends only on the EOS parameters and the accreted mass through equations (DBI18)) and (D22))., which depends only on the EOS parameters and the accreted mass through equations \ref{integral_plotted}) ) and \ref{lambda_general}) ).
+ In this limit. one inds the following scalings of the magnetic dipole moment: Hox(1lAVALB7) (model A). Hox(dkppap2) (models 12 and. D) and jx(1byAL“2B7) (model C). where Aypeop are constants.," In this limit, one finds the following scalings of the magnetic dipole moment: $\mu \propto (1-k_{\mathrm{A}}M_{\mathrm{a}} B^{-2})$ (model A), $\mu \propto (1-k_{\mathrm{B,D}}M_{\mathrm{a}}^{9/5} B^{-2})$ (models B and D) and $\mu \propto (1-k_{\mathrm{C}}M_{\mathrm{a}}^{3/2} B^{-2})$ (model C), where $k_{\mathrm{A,B,C,D}}$ are constants."
+ We confirm in Section 5. that the realistic EOS (model I2) is wellapproximated by nioclel C and hence follows the same scaling., We confirm in Section \ref{section_5} that the realistic EOS (model E) is wellapproximated by model C and hence follows the same scaling.
+" Lt is important. to note that these pCAZ.) sealings are only valid in the AL, limit (i.c. A,xAM. where Ado is EOS-depencent)."," It is important to note that these $\mu(M_{\mathrm{a}})$ scalings are only valid in the $M_{\mathrm{a}}$ limit (i.e. $M_{\mathrm{a}} \leq M_{\mathrm{c}}$, where $M_{\mathrm{c}}$ is EOS-dependent)."
+" For AL,cAlo. the analytical solution no longer applies aud numerical results have to be used."," For $M_{\mathrm{a}} > M_{\mathrm{c}}$, the analytical solution no longer applies and numerical results have to be used."
+ In Fig. 3.," In Fig. \ref{fig:dipole_radius},"
+" we plot μμ as a function of altitude above the surface for models AD by replacing Ru, with rin equation (14)).", we plot $\mu/\mu_{\mathrm{i}}$ as a function of altitude above the surface for models A–D by replacing $R_{\mathrm{m}}$ with $r$ in equation \ref{multipole_moments}) ).
+ The purpose is to illustrate how the screening currents are distributed racially for dillerent EOS., The purpose is to illustrate how the screening currents are distributed radially for different EOS.
+ The aceretec masses are chosen to be the characteristic masses Af. of cach model in Table 1.., The accreted masses are chosen to be the characteristic masses $M_{\mathrm{c}}$ of each model in Table \ref{table:eos}.
+" The dipole moment urns up by z5 per cent at reLy, because the Neumann voundary condition 0uc/080=0. which holds the field lines »rpendicular to the outer ericl boundary. does not apply strictly to a dipole field."," The dipole moment turns up by $\approx 5$ per cent at $r \approx
+R_{\mathrm{m}}$ because the Neumann boundary condition $\partial \psi/\partial
+\theta = 0$, which holds the field lines perpendicular to the outer grid boundary, does not apply strictly to a dipole field."
+ For the isothermal mountain (mocoel A). the screening.ὃν currents are located. 10n107Do. times closer o the neutron star surface than in models D. anc the isocdensity contours contract towards the surface by the same actor (see Section 4.5)).," For the isothermal mountain (model A), the screening currents are located $10^{1} - 10^{2}$ times closer to the neutron star surface than in models B–D, and the isodensity contours contract towards the surface by the same factor (see Section \ref{section_4:hydromagnetic_structure}) )."
+ Fie., Fig.
+ 4 displavs the mass. quacrupole moment. of. the mountain. expressed in terms of the mass ellipticity €. as a function of AZ.," \ref{fig:ellipticity_mass} displays the mass quadrupole moment of the mountain, expressed in terms of the mass ellipticity $\epsilon$, as a function of $M_{\mathrm{a}}$."
+" The ellipticity is given by ο=|.{μμ fo. where £7; denotes the moment-of-inertia tensor. the z-axis lies along the magnetic axis of svmmetry. and we deline dy=(2/5)AL,Reie"," The ellipticity is given by $\epsilon = |I_{zz} - I_{yy}|/I_{0}$ , where $I_{ij}$ denotes the moment-of-inertia tensor, the $z$ -axis lies along the magnetic axis of symmetry, and we define $I_{0} =
+(2/5)M_{\ast}R_{\mathrm{in}}^{2}$."
+ do zeroth order. both Ad. and « are proportional to the surface density pour.," To zeroth order, both $M_{\mathrm{a}}$ and $\epsilon$ are proportional to the surface density $\rho_{\mathrm{surf}}$."
+" Hence. the ellipticity is proportional to accreted mass for AL,< "," Hence, the ellipticity is proportional to accreted mass for $M_{\mathrm{a}} < M_{\mathrm{c}}$ ."
+"Ado. M AL,7Ade. the hyelrostatic pressure overwhelms the Lorentz force and the mountain spreads laterally. distributing the extra accreted mass evenly over a larger area (the enlarged magnetic polar cap) ancl moderating the growth of the ellipticity such that cleαλ« L/Mi."," At $M_{\mathrm{a}}
+\approx M_{\mathrm{c}}$, the hydrostatic pressure overwhelms the Lorentz force and the mountain spreads laterally, distributing the extra accreted mass evenly over a larger area (the enlarged magnetic polar cap) and moderating the growth of the ellipticity such that $\mathrm{d}\epsilon/\mathrm{d}M_{\mathrm{a}} < 1/\mathrm{M}_{\sun}$ ."
+ The apparent turnover in ο after it peaks in Vig., The apparent turnover in $\epsilon$ after it peaks in Fig.
+ 4 is anumerical artefact.which sets in as the convergence of the numerical algorithm worsens (see Section 4.1)).," \ref{fig:ellipticity_mass} is anumerical artefact,which sets in as the convergence of the numerical algorithm worsens (see Section \ref{section_4:dipole_moment}) )."
+" In realitw. for Al,2 Alo. the ellipticity saturates at the value where deαλα=0 in a hard-surface model."," In reality, for $M_{\mathrm{a}} > M_{\mathrm{c}}$ , the ellipticity saturates at the value where $\mathrm{d}\epsilon/\mathrm{d}M_{\mathrm{a}} = 0$ in a hard-surface model."
+" 2. examined accretion on to a non-rigid neutron starcrust. thereby allowing the accreted matter to sink. and showed that the ellipticity does not saturate (ic. de/dAZ,2 0) up to AM 0.12M.;."," \citet{wette2010} examined accretion on to a non-rigid neutron starcrust, thereby allowing the accreted matter to sink, and showed that the ellipticity does not saturate (i.e. $\mathrm{d}\epsilon/\mathrm{d}M_{\mathrm{a}} > 0$ ) up to $M_{\mathrm{a}} \lesssim 0.12
+\mathrm{M}_{\sun}$ ."
+ Despite this monotonic increase. the ellipticity of," Despite this monotonic increase, the ellipticity of"
+In order to measure the elfective band-power windows in the present case. we ran four additional models with the ficlucial power spectrum of Croft et shorteiteCWBOL reduced by in four different wave-bands. corresponding to four & values.,"In order to measure the effective band-power windows in the present case, we ran four additional models with the fiducial power spectrum of Croft et \\shortcite{CWB01} reduced by in four different wave-bands, corresponding to four $k$ values."
+ Linear power spectra for these models are shown in Figure lOaa. and their respective Dux power spectra are shown in retlieCLbb. Ehe dillerence between the dillerent models is comparable to the residual uncertainty in the mean [lux power spectrum reffisltL)). and thus we are only able to measure. the dillerence. between the various models with the reduced power to an accuracy of about30%.," Linear power spectra for these models are shown in Figure \ref{figCL}a a, and their respective flux power spectra are shown in \\ref{figCL}b b. The difference between the different models is comparable to the residual uncertainty in the mean flux power spectrum \\ref{figRL}) ), and thus we are only able to measure the difference between the various models with the reduced power to an accuracy of about."
+. A more accurate determination of the banc-power windows would require an implausibly [arge number of simulations., A more accurate determination of the band-power windows would require an implausibly large number of simulations.
+ As can be seen. the abrupt reductions in the linear power spectrum leac not to abrupt reductions in the [lux »ower spectrum. but rather to broad depressions.," As can be seen, the abrupt reductions in the linear power spectrum lead not to abrupt reductions in the flux power spectrum, but rather to broad depressions."
+ Figure 11 shows the inferred. band-power windows hemselves. the derivative of the lux power spectrum with respect to the input lincar power spectrum at a given wavenumber. both for the case when the amplitudes: of he input power spectra are kept fixed. and when the amplitudes. are adjusted to fit the observed. [lux power spectrum.," Figure \ref{figCD} shows the inferred band-power windows themselves, the derivative of the flux power spectrum with respect to the input linear power spectrum at a given wavenumber, both for the case when the amplitudes of the input power spectra are kept fixed, and when the amplitudes are adjusted to fit the observed flux power spectrum."
+ Evicenth the effective bancd-power windows are quite broad., Evidently the effective band-power windows are quite broad.
+ reffigCD.— indicates. that the. banc-power windows are broader at smaller scales., \\ref{figCD} indicates that the band-power windows are broader at smaller scales.
+ Figure 120 attempts to quantify this feature by showing the same derivatives as reffigC D.. except as a function of r=2x/K& rather than &.," Figure \ref{figCR} attempts to quantify this feature by showing the same derivatives as \\ref{figCD}, except as a function of $r=2\pi/k$ rather than $k$."
+ ‘To bring out the similarity of shapes. the four band-power windows from reffigCD. have been shifted horizontally so that their center points rj=2x/hy coincide (the centers of the bancd-power windows Ay do not necessarily fall on A’. but are always within one bin in & space. from which we conclude that this discrepancy is most likely due to the finite. A-space sampling). and vertically to the same maximum.," To bring out the similarity of shapes, the four band-power windows from \\ref{figCD} have been shifted horizontally so that their center points $r_0 = 2\pi/k_0$ coincide (the centers of the band-power windows $k_0$ do not necessarily fall on $k^\prime$, but are always within one bin in $k$ space, from which we conclude that this discrepancy is most likely due to the finite $k$ -space sampling), and vertically to the same maximum."
+ The four band-power windows do not have exactly the same width in real space., The four band-power windows do not have exactly the same width in real space.
+ Fitting to à Gaussian of width o. we find that σ as a function⋅. of ⋅∕A can be approximated. as In order to illustrate that itis indeed peculiar velocities that are responsible for the smoothing of the recovered linear power spectrum. we show in reffigCIt. the derivative of the [lux power spectrum with respect to the prior linear power spectrum ad jo-0.0096(Km/s)! (an equivalent of the long-dashed line) with all peculiar velocities set to zero.," Fitting to a Gaussian of width $\sigma$, we find that $\sigma$ as a function of $k^\prime$ can be approximated as In order to illustrate that it is indeed peculiar velocities that are responsible for the smoothing of the recovered linear power spectrum, we show in \\ref{figCR} the derivative of the flux power spectrum with respect to the prior linear power spectrum at $k^\prime=0.0096\dim{(km/s)}^{-1}$ (an equivalent of the long-dashed line) with all peculiar velocities set to zero."
+ In this case the peak is much narrower. and its width is explained by finite sampling of the k-space alone.," In this case the peak is much narrower, and its width is explained by finite sampling of the $k$ -space alone."
+ Lt is important to underscore here that it is the power spectrum. itself which is smoothed. and not the Εαν.," It is important to underscore here that it is the power spectrum itself which is smoothed, and not the flux."
+ Croft et did consider the ellect. of smoothing the Lyman-alpha absorption spectrum 110 in their paper). but the effect that we discuss here is. dillerent it is a," Croft et \\shortcite{CWB01} did consider the effect of smoothing the Lyman-alpha absorption spectrum 10 in their paper), but the effect that we discuss here is different – it is a"
+"In order to attenuate possible calibration problems. we thus calculated both the effective wavelengths ;l,jj and count- conversion factors CF for the UVOT filters by folding the BL Lacertae spectrum with their effective areas2008).","In order to attenuate possible calibration problems, we thus calculated both the effective wavelengths $\lambda_{\rm eff}$ and count-rate-to-flux conversion factors $\rm CF$ for the UVOT filters by folding the BL Lacertae spectrum with their effective areas."
+. We first built a composite observed spectrum of BL Lacertae by combining a mean OM spectrum (obtained from thethree XMM-Newton pointings of 2007—2008. 2009)) with an average UVOT spectrum (resulting from 16 UVOT observing epochs analysed in. this paper)," We first built a composite observed spectrum of BL Lacertae by combining a mean OM spectrum (obtained from thethree XMM-Newton pointings of 2007--2008, ) with an average UVOT spectrum (resulting from 16 UVOT observing epochs analysed in this paper)."
+ To compensate for the different brightness state. we increased the OM flux densities by ~6% so that the two spectra match in the V band.," To compensate for the different brightness state, we increased the OM flux densities by $\sim 6$ so that the two spectra match in the $V$ band."
+ The composite spectrum ts shown in refcalib. (top panel). together with its log-parabolic fit that we used in the folding procedure.," The composite spectrum is shown in \\ref{calib} (top panel), together with its log-parabolic fit that we used in the folding procedure."
+ The resulting effective wavelengths (see 88 in 2008)) are: 5439. 4381. 3500. 2776. 2295. and 2225 ffor the v. b. t. ivl. ΙΟ. and uvw2 filters. respectively. showing a clear shift towards longer wavelengths in the ultraviolet.," The resulting effective wavelengths (see 8 in ) are: 5439, 4381, 3500, 2776, 2295, and 2225 for the $v$, $b$, $u$, $uvw1$, $uvm2$, and $uvw2$ filters, respectively, showing a clear shift towards longer wavelengths in the ultraviolet."
+ As forthe count-rate-to-flux conversion factors. we obtained 2.60. 1.47. 1.65. 4.31. 8.54. and 6.72x107ereem™s!Av! from. the v. to. the. inw2.— band. respectively.," As forthe count-rate-to-flux conversion factors, we obtained 2.60, 1.47, 1.65, 4.31, 8.54, and $6.72 \times 10^{-16} \rm \, erg \, cm^{-2} \, s^{-1} \, \AA^{-1}$ from the $v$ to the $uvw2$ band, respectively."
+ These new CF differ from those given by for the GRB models by € [ο with the only exceptions of CFGnw1) and CF(nw2). which are nowl," These new $\rm CF$ differ from those given by for the GRB models by $\la 1$ , with the only exceptions of ${\rm CF}(uvw1)$ and ${\rm CF}(uvw2)$, which are now."
+arger?.. The new 2g would produce a decrease of extinction in the avi land ava? bands with respect to those given by(2008).. and an increase in the 42 band.," The new $\lambda_{\rm eff}$ would produce a decrease of extinction in the $uvw1$ and $uvm2$ bands with respect to those given by, and an increase in the $uvw2$ band."
+ Indeed. the Galactic mean extinction. curve shows a dramatic bump peaking at À~275A owing to absorption by graphite dust.," Indeed, the Galactic mean extinction curve shows a dramatic bump peaking at $\lambda \sim 2175 \, \AA$ owing to absorption by graphite dust."
+ Actually. an accurate evaluation of extinctioni1 this critical frequency range requires that the law ts folded with the filter's effective area and BL Lacertae spectrum. similarly to what was done above for the ο and CF: where Aq Is the extinction in the A band. Εκ is the effective area of that band. and Fy 1s the source flux density.," Actually, an accurate evaluation of extinctionin this critical frequency range requires that the law is folded with the filter's effective area and BL Lacertae spectrum, similarly to what was done above for the $\lambda_{\rm eff}$ and $\rm CF$: where $A_{\Lambda}$ is the extinction in the $\Lambda$ band, $E_\Lambda$ is the effective area of that band, and $F_{\lambda}$ is the source flux density."
+ The result is a Galactic extinction of 1.10. 1.44. 1.74. 2.40. 3.04. and 2.92 mag from the v to the ww? band. respectively.," The result is a Galactic extinction of 1.10, 1.44, 1.74, 2.40, 3.04, and 2.92 mag from the $v$ to the $uvw2$ band, respectively."
+ We verified the stability of our results by iterating the procedure with the recalibrated UVOT flux densities and lar., We verified the stability of our results by iterating the procedure with the recalibrated UVOT flux densities and $\lambda_{\rm eff}$.
+ The bottom panel of refealib. shows the mean SED obtained after recalibration of the UVOT data according to our procedure., The bottom panel of \\ref{calib} shows the mean SED obtained after recalibration of the UVOT data according to our procedure.
+ For comparison. we also show the OM and UVOT SEDs derived from the average spectra shown in the top panel. for which the amount of Galactic extinction was calculated from the law at the standard ;ty.," For comparison, we also show the OM and UVOT SEDs derived from the average spectra shown in the top panel, for which the amount of Galactic extinction was calculated from the law at the standard $\lambda_{\rm eff}$."
+ Notice that the recalibration process has shifted ἐν5) redward so much that it overlaps with the standard εἰ 9)., Notice that the recalibration process has shifted $\lambda_{\rm eff}(uvw2)$ redward so much that it overlaps with the standard $\lambda_{\rm eff}(uvm2)$ .
+ This is because of the noticeable red tail of the (2 filter as well as to the red BL Lacertae spectrum., This is because of the noticeable red tail of the $uvw2$ filter as well as to the red BL Lacertae spectrum.
+ Recalibration has solved the iywl-dip problem. which ts a common feature of UVOT SEDs for a number of blazars at different redshifts2009).," Recalibration has solved the $uvw1$ -dip problem, which is a common feature of UVOT SEDs for a number of blazars at different redshifts."
+. Moreover. it seems to confirm the UV excess claimed by that was ascribed to thermal emission from the accretion disc. even if this excess may be less pronounced than indicated by the OM data.," Moreover, it seems to confirm the UV excess claimed by that was ascribed to thermal emission from the accretion disc, even if this excess may be less pronounced than indicated by the OM data."
+ Our analysis highlights the importance of calculating the amount of extinction. in. the critical UV bands. close to the 2175 bbump. by folding the Galactic mean extinction law through the effective area curves and source spectrum.," Our analysis highlights the importance of calculating the amount of extinction in the critical UV bands, close to the 2175 bump, by folding the Galactic mean extinction law through the effective area curves and source spectrum."
+ In any case. as pointed out by(2007).. one has to keep in mind that the use of an average dereddening curve implies a significant error owing to the scatter of Galactic extinction curves.," In any case, as pointed out by, one has to keep in mind that the use of an average dereddening curve implies a significant error owing to the scatter of Galactic extinction curves."
+ The X-ray Telescope (XRT: 2005)) data were processed with version 0.12.3 of the task contained in the FTOOLS package. applying standard screening criteria.," The X-ray Telescope (XRT; ) data were processed with version 0.12.3 of the task contained in the FTOOLS package, applying standard screening criteria."
+ Inspection of the light curves revealed that the count rate was low. from 0.15 to 0.24 counts s!. so that observations were performed in photon counting mode. and no correction for pile-up was necessary.," Inspection of the light curves revealed that the count rate was low, from 0.15 to 0.24 counts $\rm s^{-1}$, so that observations were performed in photon counting mode, and no correction for pile-up was necessary."
+ Source and background spectra were extracted with from a circular region of 20 pixel (47 aresec) radius centred on the source and from a surrounding annulus of 30 and 50 pixel radii. respectively.," Source and background spectra were extracted with from a circular region of 20 pixel (47 arcsec) radius centred on the source and from a surrounding annulus of 30 and 50 pixel radii, respectively."
+ We used version O11 of the response matrix available in the HEASARC calibration database (CALDB). and calculated the ancillary response file withxrtmkarf. using the exposure map created by xrtpipeline.," We used version 011 of the response matrix available in the HEASARC calibration database (CALDB), and calculated the ancillary response file with, using the exposure map created by ."
+ The source spectra were binned with to have a minimum of 20 counts in each bin. and they were finally analysed with version 12.5.1 of the task. using the energy channels greater than 0.3 keV. Spectral analysis was performed for each observation following(2009): we first fitted à single power law with freeabsorption. and then fixed the Galactic absorption to Ny=3.4x107!em-7. which takes into account both atomic and molecular column density.," The source spectra were binned with to have a minimum of 20 counts in each bin, and they were finally analysed with version 12.5.1 of the task, using the energy channels greater than 0.3 keV. Spectral analysis was performed for each observation following: we first fitted a single power law with free, and then fixed the Galactic absorption to $N_{\rm H} = 3.4 \times 10^{21} \, \rm cm^{-2}$, which takes into account both atomic and molecular column density."
+ Statisties is not good enough to evaluate if a double power law model can improve the fit., Statistics is not good enough to evaluate if a double power law model can improve the fit.
+ The results of spectral fitting on XRT data are reported in Table 2. for all observations with an exposure longer than 3 ksee: 11 gives the date and start time of the observation; 22 its duration: 33 the hydrogen column: 44 the power law photon index: 55 the | keV flux density: 66 the y7/v (and degrees of freedom)., The results of spectral fitting on XRT data are reported in Table \ref{xrt_pow} for all observations with an exposure longer than 3 ksec; 1 gives the date and start time of the observation; 2 its duration; 3 the hydrogen column; 4 the power law photon index; 5 the 1 keV flux density; 6 the $\chi^2/\nu$ (and degrees of freedom).
+ One spectrum (August 29) is shown in refxrt.., One spectrum (August 29) is shown in \\ref{xrt}. .
+ Fits with free absorption resulted ina very variable Ny. which is unlikely to correspond to a real change of absorption.," Fits with free absorption resulted ina very variable $N_{\rm H}$ , which is unlikely to correspond to a real change of absorption."
+ The average and median values are 3.46 and 3.44x107! em7. respectively. confirming that the value assumed for the Galactic," The average and median values are 3.46 and $3.44 \times 10^{21} \, \rm cm^{-2}$ , respectively, confirming that the value assumed for the Galactic"
+"For the inner and outer bars of 3357 we only have blue spectra, from 3990 to 4440A.","For the inner and outer bars of 357 we only have blue spectra, from 3990 to 4440."
+". There is only one Iron index in the Lick system within this range, at 4383Á."," There is only one Iron index in the Lick system within this range, at 4383."
+". Unfortunately, the red pseudo-continuum of this Lick Fe4383 index falls partially out of our blue spectra, so we have redefined it in order to introduce a new Fe4383°"" index suitable for our spectral range requirements."," Unfortunately, the red pseudo-continuum of this Lick Fe4383 index falls partially out of our blue spectra, so we have redefined it in order to introduce a new $^{SR}$ index suitable for our spectral range requirements."
+" Fe4383?7 is named after (due to its motivation) and keeps the blue pseudo-continuum of the Lick Fe4383 index, but modifies the range definitions for the main bandpass and the red pseudo-continuum (see Table 2)). B2 ??,, Vazdekisetal.(2010) C1,,"," $^{SR}$ is named after (due to its motivation) and keeps the blue pseudo-continuum of the Lick Fe4383 index, but modifies the range definitions for the main bandpass and the red pseudo-continuum (see Table \ref{tab:fe}) \ref{fig:spFe4383} \ref{sec:cc}, \citet{2010MNRAS.404.1639V}
+ \ref{fig:cctest},"
+ emission are includel.,$\beta$ emission are included.
+ This limits the sample to the brighter stars at each spectral type that mect the SNR criterion and does not niclude all of the AL dafs that are SCLSIive to Πα cussion., This limits the sample to the brighter stars at each spectral type that meet the SNR criterion and does not include all of the M dwarfs that are sensitive to $\alpha$ emission.
+ Spectral types earlier than AP and later than MT were nof inclided due to an iusufficieit sample size of active stars., Spectral types earlier than M2 and later than M7 were not included due to an insufficient sample size of active stars.
+ As seen in previous studies. the magnetic activity fractions decrease as a function of Galactic heielt. confirming an age-activity relationsip: stars closer to Plane are statistically vounger and more likely to IC activo.," As seen in previous studies, the magnetic activity fractions decrease as a function of Galactic height, confirming an age-activity relationship; stars closer to the Plane are statistically younger and more likely to be active."
+ Figue 55. also cdemonstrates that IIo au (ID trace each other extremely well for he entire. DRT saluple. inplvius that the I> activity ifotinies are essentialv the same as those for Πα.," Figure \ref{betafrac} also demonstrates that $\alpha$ and $\beta$ trace each other extremely well for the entire DR7 sample, implying that the $\beta$ activity lifetimes are essentially the same as those for $\alpha$."
+ The other activity tracOrs aso appear to correlate with Ta eudssion., The other activity tracers also appear to correlate with $\alpha$ emission.
+ Figure6 shows jc activity fractious for all five enussion lines in ALL s as a functiou of vertical distance from the Calacti UM‘plane., Figure \ref{m4frac} shows the activity fractions for all five emission lines in M4 dwarfs as a function of vertical distance from the Galactic plane.
+ Each panel plots the activity fractions for stars where the data are good enough to measure (1.0. are sensitive to) I> (top left). II* (top right). IIó (bottom oft) and Call I (bottom right) activity as well as t1ο ehussion lines redder than that specific tracer.," Each panel plots the activity fractions for stars where the data are good enough to measure (i.e. are sensitive to) $\beta$ (top left), $\gamma$ (top right), $\delta$ (bottom left) and CaII K (bottom right) activity as well as the emission lines redder than that specific tracer."
+ The activdtv fractious for all of the lines are in excellent aerecme wewith each other aud indicate that all five ciissiou dites are produced for the same aout of time curing tie lifetime of au AL dwarf., The activity fractions for all of the lines are in excellent agreement with each other and indicate that all five emission lines are produced for the same amount of time during the lifetime of an M dwarf.
+ This new result is complemieitary to previous studies that found strong correla10119 αλλος different cussion line streugths in M dwarts(Rauscher&Marcy2006:Walkow-icz&Dawley 2009)..," This new result is complementary to previous studies that found strong correlations among different emission line strengths in M dwarfs \citep{rm06,walkowicz09}. ."
+ Altrough the relative emission line streneths for iucividial stars may fluctuate over time2001).. the mean values of the emission lines are wel correlated in the large DR? sample.," Although the relative emission line strengths for individual stars may fluctuate over time, the mean values of the emission lines are well correlated in the large DR7 sample."
+ The unprecedented size and Galactic distribution of our spectroscopic siuuple confirms that not only the strength. but the duratio1 of activity appears to be similar for all of the Baliner aud Call activity indicators.," The unprecedented size and Galactic distribution of our spectroscopic sample confirms that not only the strength, but the duration of activity appears to be similar for all of the Balmer and CaII activity indicators."
+ Figure 6 also demonstrates a possible (and inrportaut) selection effect xeseut in iaegnitude-Imited or activitv-selected saies., Figure \ref{m4frac} also demonstrates a possible (and important) selection effect present in magnitude-limited or activity-selected samples.
+ A close examination of Figure 6 (aud Figure 5 )) reveals that the activity fractious are larger for the stars flat are sensitive to the bluer cussion lines., A close examination of Figure \ref{m4frac} (and Figure \ref{betafrac}) ) reveals that the activity fractions are larger for the stars that are sensitive to the bluer emission lines.
+ This troil was not expected since the stars used to compute the activity fraction were drawn from the same volunes iu cach of the four paucls., This trend was not expected since the stars used to compute the activity fraction were drawn from the same volumes in each of the four panels.
+ The only difference was that the stars selected. for the bluer emissou Line analyses were required to have higher SNR., The only difference was that the stars selected for the bluer emission line analyses were required to have higher SNR.
+ With this in wind. we can explain the different activity fractions as ollows.," With this in mind, we can explain the different activity fractions as follows."
+" The actual spread in absolute magnitude (οι ""uimositv) at a eiven spectral type cau be quite Some of this spread uv be due to differences iu he physicalbL. properties of the stars.", The actual spread in absolute magnitude (or luminosity) at a given spectral type can be quite Some of this spread may be due to differences in the physical properties of the stars.
+ Tawleyetal.(1996) showed that magjeticallv active stars are brighter in AM han their inactive counterparts., \citet{hawley96} showed that magnetically active stars are brighter in $M_V$ than their inactive counterparts.
+" Receutly. Dochauskietal.(2011.hereafterPaperIT) fouud that both active and higher metalicity M. chwarts (nav stars are both) appear to be brighter in M, at a given color or spectral type."," Recently, \citet[hereafter Paper II]{boostat}
+ found that both active and higher metallicity M dwarfs (many stars are both) appear to be brighter in $M_r$ at a given color or spectral type."
+ Because the bluer Cluission lines that we are exanuning in Figure 6 require higher SNR spectra to accurately miecasure activity. the ALL dwarfs at a given distance that mect the SNR criterion are preferetially nxe hnüuous than other stars in the same distance biu.," Because the bluer emission lines that we are examining in Figure \ref{m4frac} require higher SNR spectra to accurately measure activity, the M4 dwarfs at a given distance that meet the SNR criterion are preferentially more luminous than other stars in the same distance bin."
+ If activity is correlated with luuinositvy. our SNR selection will bias our sample toward coutaimine a higher fraction of active stars.," If activity is correlated with luminosity, our SNR selection will bias our sample toward containing a higher fraction of active stars."
+ Iu addition to activity fractions. we also 1ivestieated the mielan Πορ m each cussion liue as a fuiction of spectral typo," In addition to activity fractions, we also investigated the median luminosity in each emission line as a function of spectral type."
+ Previous studies have investigated how, Previous studies have investigated how
+the CIAO scriptpsertract. and fit the spectra in NSPEC (Arnaud1996)5.,"the CIAO script, and fit the spectra in XSPEC (Arnaud."
+. We correct the effective area functions for the time-dependent low-energv quantum effieiencv We exclude bins with most photons below 0.3 keV or above 10 keV. We attempt to fit three models to hese spectra. all with photoclectric absorption as a free xuaneter forced to be equal to or greater than the cluster value (9.1:«1079 cm 72).," We correct the effective area functions for the time-dependent low-energy quantum efficiency We exclude bins with most photons below 0.3 keV or above 10 keV. We attempt to fit three models to these spectra, all with photoelectric absorption as a free parameter forced to be equal to or greater than the cluster value $9.4\times10^{20}$ $^{-2}$ )."
+ For all aualvsis in this paper. we use photoclectric absorption N-ray cross-sections of Balucinska-Chiurch AIcCanunon (1992) in the NSPEC ghabs model.," For all analysis in this paper, we use photoelectric absorption X-ray cross-sections of Balucinska-Church McCammon (1992) in the XSPEC phabs model."
+ Our models ave: a thermal bremsstraliluug spectruni as associated with CVs: a power-law model: aud a lyvdrogen atinosphere model (Lloyd 2003) as appropriate or GQLAINBs containing thermal ucutrou stars with D«1029 C. with the radius fixed at 10 kin.," Our models are: a thermal bremsstrahlung spectrum as associated with CVs; a power-law model; and a hydrogen atmosphere model (Lloyd 2003) as appropriate for qLMXBs containing thermal neutron stars with $B<10^{10}$ G, with the radius fixed at 10 km."
+ The dichotomy )etween harder and softer sources apparent in the X-rav CMDSs is also clear in the spectral fitting. with CX2 and CX6 showing good fits to the hydrogen atmosphere spectral imodols while CNT. CN3. CNL and CNS do rot.," The dichotomy between harder and softer sources apparent in the X-ray CMDs is also clear in the spectral fitting, with CX2 and CX6 showing good fits to the hydrogen atmosphere spectral models while CX1, CX3, CX4, and CX5 do not."
+ CN2 and €X6 require large values for a powerlaw jiofon index (= 5) aud very small bronisstrahluug cluperatures (<<0.6 keV). which are not consistent models or any known phlivsical sources at these Iuuinosities.," CX2 and CX6 require large values for a powerlaw photon index $>5$ ) and very small bremsstrahlung temperatures $<0.6$ keV), which are not consistent models for any known physical sources at these luminosities."
+ CXL. CN3. CXL and CNS eive. breimisstralilung temperatures consistent with ~7 keV or more. as appropriate for ininous CVs. particularly magnetic CVs (Eracleous. Halpern Patterson 1991: Mul 2001).," CX1, CX3, CX4, and CX5 give bremsstrahlung temperatures consistent with $\sim7$ keV or more, as appropriate for luminous CVs, particularly magnetic CVs (Eracleous, Halpern Patterson 1991; Mukai 2001)."
+ Mekal models (Liedalil et. al., Mekal models (Liedahl et al.
+ 1995) eive iudistinguishable results. οἼναι he low metallicity (|Fe/II]|2.1.75) aud high temperatures.," 1995) give indistinguishable results, given the low metallicity ([Fe/H]=–1.75) and high temperatures."
+ This result confinis our tentative classification of these sources as cataclysuuc variables in 82.1., This result confirms our tentative classification of these sources as cataclysmic variables in 2.1.
+ We note that CNG requires a hieher Nyy than the cluster value for auv of our models. while the other sources are consistent with the cluster value.," We note that CX6 requires a higher $N_H$ than the cluster value for any of our models, while the other sources are consistent with the cluster value."
+ ITeimke et al. (, Heinke et al. (
+2003a) note euliauced Vy; towards NS and X7 in £7 Tuc. prestunably from gas inside or surroundius the system.,"2003a) note enhanced $N_H$ towards X5 and X7 in 47 Tuc, presumably from gas inside or surrounding the system."
+ Our preferred spectral fits to these six sources are shown in Figure Ll. aud results for all three models are listed in Table 3.," Our preferred spectral fits to these six sources are shown in Figure 4, and results for all three models are listed in Table 3."
+ For the remaining sources within the halfiuass radius (except CX15. which has a very unusual spectrum: see &2.1). we extract a combined spectrün aud fit this to derive the mean spectral shape and dhunünositv/couut ratio.," For the remaining sources within the half-mass radius (except CX15, which has a very unusual spectrum; see 2.1), we extract a combined spectrum and fit this to derive the mean spectral shape and luminosity/count ratio."
+ We extract a total of 235 counts. and fit them with a thermal bremsstrahhme model of kT=2.3!p keV (for fixed Ny—9.1.«10°? 7). (AZ=LS for 8 dof).," We extract a total of 235 counts, and fit them with a thermal bremsstrahlung model of $=2.3^{+.91}_{-.64}$ keV (for fixed $N_H=9.4\times10^{20}$ $^{-2}$ ), $\chi^2_{\nu}=1.45$ for 8 dof)."
+ A mela fit gives very similar results. while fits with a power law or blackbody require very different coluun deusitics:.," A mekal fit gives very similar results, while fits with a power law or blackbody require very different column densities;."
+" The powerlaw requires Ny=2649<107?"" cmi2 with a photon iudex of 2.14 (42=L8 for 8 dof). while the blackbody fit requires Ny=0!4«10? 7. much less than the cluster value."," The powerlaw requires $N_H=26\pm9\times10^{20}$ $^{-2}$ with a photon index of $^{.4}_{-.3}$ $\chi^2_{\nu}=1.8$ for 8 dof), while the blackbody fit requires $N_H=0^{+3}_{-0}\times10^{20}$ $^{-2}$, much less than the cluster value."
+ These results indicate that the fainter sources have lower teiiperatures than the bright CVs. as expected for a mux of active biuaries and (perhaps nounmagnetic) CVs. as seen in £7 Tuc (Ediiouds et al.," These results indicate that the fainter sources have lower temperatures than the bright CVs, as expected for a mix of active binaries and (perhaps nonmagnetic) CVs, as seen in 47 Tuc (Edmonds et al."
+ 2003a. b).," 2003a, b)."
+ We use the broiisstraliluug spectral fits to derive fluxes., We use the bremsstrahlung spectral fits to derive fluxes.
+" To calculate the Iuninosities of each of the fainter sources, we multiplied their iutegrated πιοαν by the ratio of each sources counts to the combined source counts (Table 1)."," To calculate the luminosities of each of the fainter sources, we multiplied their integrated luminosity by the ratio of each source's counts to the combined source counts (Table 1)."
+ We do this for both the 0.5-2.5 keV baud aud the 0.5-6 keV. baud., We do this for both the 0.5-2.5 keV band and the 0.5-6 keV band.
+" Derived huuinosityv errors are sniplv Poisson or Celrels errors from the detected counts, without including spectral uncertainties. aud are thus uudoerestinates."," Derived luminosity errors are simply Poisson or Gehrels errors from the detected counts, without including spectral uncertainties, and are thus underestimates."
+ We extracted event files frou: cach detected source within the halfimass radius aud tested them using the IRAFvorfst to attempt to disprove the bypothesis that the source flux is coustant., We extracted event files from each detected source within the half-mass radius and tested them using the IRAF to attempt to disprove the hypothesis that the source flux is constant.
+ Two sources (CN1 and CX8) showed variability at the coufidence level according to both the Cramervou Mises test and LNolmoeorovSuuiruov tests (Daniel 19903., Two sources (CX1 and CX8) showed variability at the confidence level according to both the Cramer-von Mises test and Kolmogorov-Smirnov tests (Daniel 1990).
+ CXE showed variability at the confidence level in both tests. while no other source showed evidence of variability.," CX4 showed variability at the confidence level in both tests, while no other source showed evidence of variability."
+ We present the lighteurves roni these three sources. plus the (uouvariable) Biehteurve your CN2 (a probable qLMXND) in Figure 5.," We present the lightcurves from these three sources, plus the (nonvariable) lightcurve from CX2 (a probable qLMXB) in Figure 5."
+ Clear flares are present in all three of the variable sources., Clear flares are present in all three of the variable sources.
+ X-ray faring may be present iu either CVs or ABs. but is not expected from MSPs.," X-ray flaring may be present in either CVs or ABs, but is not expected from MSPs."
+ The laree flare visible from CNS is renuiscent of a flare from an AB. but we cannot make any firm statements about these sources from their variability alone.," The large flare visible from CX8 is reminiscent of a flare from an AB, but we cannot make any firm statements about these sources from their variability alone."
+ The Cramer-vou Mises aud EKolhuogorov-Siiirnov ests are naturally far more sensitive to variability from xieht sources than faint sources. so the lack of identified variability from faint sources does not indicate that they did not vary during the observation.," The Cramer-von Mises and Kolmogorov-Smirnov tests are naturally far more sensitive to variability from bright sources than faint sources, so the lack of identified variability from faint sources does not indicate that they did not vary during the observation."
+ The racial distribution of N-rav sources m a dvnanucally relaxed cluster allows au estinate of the average ass of the ταν sources., The radial distribution of X-ray sources in a dynamically relaxed cluster allows an estimate of the average mass of the X-ray sources.
+ Ποιοetal.(2003€) describe a procedure for estimating the typical qLAINB ass from the spatial distribution of a sample of 20 qLAINBs in seven clusters., \citet{heinke03c} describe a procedure for estimating the typical qLMXB mass from the spatial distribution of a sample of 20 qLMXBs in seven clusters.
+ This procedure is based. on inuixinimuu-likeliood fitting of a parameterized form to the radial profile of the source distribution., This procedure is based on maximum-likelihood fitting of a parameterized form to the radial profile of the source distribution.
+ The key parameter is the ratio q=Aly/AL. of the source mass to the mass of the typical star that defines the optical core radius., The key parameter is the ratio $q = M_X/M_\ast$ of the source mass to the mass of the typical star that defines the optical core radius.
+ The approach asstuues that the spatial distribution of these typical stars is well described by a classic Kine(1966) model. which is the case for ALSO (Ferraro.Paltrinicri.Rood.&Dor-man 1999).," The approach assumes that the spatial distribution of these typical stars is well described by a classic \citet{king66} model, which is the case for M80 \citep{ferraro99}."
+. The radial profile of the source surface density takes the form. where Sy is an overall normalization aud ως is the optical core radius determined for turuoffauass stars.," The radial profile of the source surface density takes the form, where $S_0$ is an overall normalization and $r_{c\ast}$ is the optical core radius determined for turnoff-mass stars."
+ For ALSO. Ferraro.Paltrimieri.Rood.&Dorman(1999) have obtained ες=675.," For M80, \citet{ferraro99} have obtained $r_{c\ast}=6\farcs5$."
+ Tn fitting the radial profile of the source distribution in ALSO. if is necessary to correct the source sample for backeround contaiunation. aud cusure a uuiforni conrpleteness limit.," In fitting the radial profile of the source distribution in M80, it is necessary to correct the source sample for background contamination, and ensure a uniform completeness limit."
+ We address the latter by using oulv sources with more than LO counts. as we are complete to this flux limit from the cluster core out to four hilf- radii.," We address the latter by using only sources with more than 10 counts, as we are complete to this flux limit from the cluster core out to four half-mass radii."
+ The expected umuber of background sources above 10 counts is 0.7 sources within the halfinass, The expected number of background sources above 10 counts is 0.7 sources within the half-mass
+Henriques et al. 2008..,"Henriques et al. \cite{Henriq08},"
+ Henriques Thomas 2010.. Yang et al. 2009))," Henriques Thomas \cite{Henriq10}, Yang et al. \cite{Yang09}) )"
+ and their contribution to the intracluster light., and their contribution to the intracluster light.
+ However. those studies were not specifically tailored to address the tidal stripping scenario for UCD formation.," However, those studies were not specifically tailored to address the tidal stripping scenario for UCD formation."
+ Yang et al. (2009)), Yang et al. \cite{Yang09}) )
+ argue that tidal disruption of satellites will be more efficient for ratios between satellite and host halo mass., argue that tidal disruption of satellites will be more efficient for ratios between satellite and host halo mass.
+ At the same time. they show that the stellar mass contributed to the intracluster medium by dissolving satellites is higher for massive host halos ~107 M. than for lower mass halos ~10 M.," At the same time, they show that the stellar mass contributed to the intracluster medium by dissolving satellites is higher for massive host halos $\sim 10^{14}$ $_{\odot}$ than for lower mass halos $\sim 10^{12}$ $_{\odot}$."
+ Henriques et al. (2008)), Henriques et al. \cite{Henriq08}) )
+ estimate that about one half of satellite galaxies get disrupted and/or accreted to their host halos and that of the overall cluster light is found in the intracluster medium. originating from tidally stripped stars.," estimate that about one half of satellite galaxies get disrupted and/or accreted to their host halos and that of the overall cluster light is found in the intracluster medium, originating from tidally stripped stars."
+ In their most recent paper. Henriques Thomas (2010)) improve their treatment of tidal disruption.," In their most recent paper, Henriques Thomas \cite{Henriq10}) ) improve their treatment of tidal disruption."
+ In turn they revise the intra-cluster light fraction provided by disrupted galaxies upwards to30%., In turn they revise the intra-cluster light fraction provided by disrupted galaxies upwards to.
+. Their models predict of the satellites with masses between 10” and 10!° M. to be tidally disrupted., Their models predict of the satellites with masses between $10 ^9$ and $10 ^{10}$ $M_{\odot}$ to be tidally disrupted.
+ This fraction increases to when considering satellites that have already lost their dark matter halo (“orphan satellites”) in. previous interactions (Henriques 2011. private communication).," This fraction increases to when considering satellites that have already lost their dark matter halo (“orphan satellites”) in previous interactions (Henriques 2011, private communication)."
+ This ~50% global disruption fraction is lower than the fraction of ~90% needed for the limiting case that half of today’s UCDs originate from tidal processes (see above)., This $\sim$ global disruption fraction is lower than the fraction of $\sim$ needed for the limiting case that half of today's UCDs originate from tidal processes (see above).
+ Radial trends of disruption efficiency would still be expected (Bekki et al. 20091).," Radial trends of disruption efficiency would still be expected (Bekki et al. \cite{Bekki03}) ),"
+ such that this disruption fraction can be considered a lower limit to the value applicable to the clustercentric areas studied in this paper., such that this disruption fraction can be considered a lower limit to the value applicable to the clustercentric areas studied in this paper.
+ In this paper we have proposed a definition for the specific frequency of UCDs. which we denote as S4(cop.," In this paper we have proposed a definition for the specific frequency of UCDs, which we denote as $S_{N,UCD}$."
+" We adopt the following functional Sxecp=Nucpl0QAUM77agUUM. This definition normalizes the number of UCDs in a given environment to a unit host galaxy luminosity M5,. analogous to the definition of the specific frequency Succ for GCs."," We adopt the following functional $S_{N,UCD} = N_{UCD} 10^{0.4 (M_{V,host}-M_{V,0})} c_{w}$ This definition normalizes the number of UCDs in a given environment to a unit host galaxy luminosity $M_{V,host}$, analogous to the definition of the specific frequency $S_{N,GC}$ for GCs."
+ The premise of our definition is that if UCDs follow the extrapolation of the GCLF to bright magnitudes. then S4S vGc.," The premise of our definition is that if UCDs follow the extrapolation of the GCLF to bright magnitudes, then $S_{N,UCD} = S_{N,GC}$ ."
+ This premise defines the value of the zero point Myo.," This premise defines the value of the zero point $M_{V,0}$."
+ Considering UCDs as compact stellar systems with My<—]0.25 mag. we find that the value of Myo=—20 mag fulfills the above premise.," Considering UCDs as compact stellar systems with $M_V<-10.25$ mag, we find that the value of $M_{V,0}=-20$ mag fulfills the above premise."
+ For the case of extragalactic surveys with a brighter completeness limit My<-11.0 mag. we need to adopt Myo=—22 mag.," For the case of extragalactic surveys with a brighter completeness limit $M_V<-11.0$ mag, we need to adopt $M_{V,0}=-22$ mag."
+" The term c, is introduced to correct the specific frequency for the well known systematic variation of the GCLF width c with host galaxy magnitude.", The term $c_{w}$ is introduced to correct the specific frequency for the well known systematic variation of the GCLF width $\sigma$ with host galaxy magnitude.
+ Details of this correction are outlined in Sect. 2.., Details of this correction are outlined in Sect. \ref{sndefsec}.
+ In Sect., In Sect.
+ 3 we apply our proposed definition of Ss(05 to results of spectroscopic UCD searches performed by our group in the Fornax. Hydra and Centaurus galaxy clusters. and two Hickson Compact Groups.," \ref{envsn} we apply our proposed definition of $S_{N,UCD}$ to results of spectroscopic UCD searches performed by our group in the Fornax, Hydra and Centaurus galaxy clusters, and two Hickson Compact Groups."
+ We also include the Local Group., We also include the Local Group.
+ Our main finding is that the specific frequencies derived for UCDs match those of GCs very well., Our main finding is that the specific frequencies derived for UCDs match those of GCs very well.
+ For four of the SIX Investigated environments. there are GC specific frequency measurements available. allowing a direct comparison.," For four of the six investigated environments, there are GC specific frequency measurements available, allowing a direct comparison."
+" For those four environments we find a mean ΑΟ,=6.74 2.0. vs. Φλας=6.7+2.2."," For those four environments we find a mean $S_{N,UCD}=6.7 \pm$ 2.0, vs. $S_{N,GC}=6.7 \pm 2.2$."
+ The ratio of UCD-to-GC specific frequency is therefore +=1.00x0.44.," The ratio of UCD-to-GC specific frequency is therefore $\frac{6.7 \pm 2.0}{6.7 \pm
+ 2.2}=1.00 \pm 0.44$."
+ The mean ορ of all six investigated environments is 5.3 4 1.7. in agreement with the average GC specific frequency of 5.0 + 0.7 of the available literature data for the corresponding host magnitude range My<—2] mag.," The mean $S_{N,UCD}$ of all six investigated environments is 5.3 $\pm$ 1.7, in agreement with the average GC specific frequency of 5.0 $\pm$ 0.7 of the available literature data for the corresponding host magnitude range $M_V<-21$ mag."
+" Our findings are consistentwith the hypothesis that most UCDs are formed by the same process as the overall GC population. and with a similar formation efficiency to GCs,"," Our findings are consistent with the hypothesis that most UCDs are formed by the same process as the overall GC population, and with a similar formation efficiency to GCs."
+ We also present an extension of our analysis in the Fornax cluster by using the large available data set of spectroscopically confirmed UCDs and GCs., We also present an extension of our analysis in the Fornax cluster by using the large available data set of spectroscopically confirmed UCDs and GCs.
+ This literature data set comprises about 180 confirmed UCDs with My<—10.25 mag., This literature data set comprises about 180 confirmed UCDs with $M_V<-10.25$ mag.
+ We find that the specific frequencies of UCDs around giant galaxies and in the intracluster space are consistent with their being drawn from the bright tail of the GCLF., We find that the specific frequencies of UCDs around giant galaxies and in the intracluster space are consistent with their being drawn from the bright tail of the GCLF.
+ There is still room for a possible UCD overabundance in the intracluster space by a factor of ~2. given the incompleteness in spectropscopic coverage in those regions.," There is still room for a possible UCD overabundance in the intracluster space by a factor of $\sim 2$, given the incompleteness in spectropscopic coverage in those regions."
+ We do not find significant evidence of a different spatial distribution between UCDs and GCs., We do not find significant evidence of a different spatial distribution between UCDs and GCs.
+ It has been proposed that the present-day population of UCDs is indeed of sources formed via tidal stripping of dwarf galaxies and sources formed in the same process as the main GC population (e.g. Haseegan et al. 2005..," It has been proposed that the present-day population of UCDs is indeed of sources formed via tidal stripping of dwarf galaxies and sources formed in the same process as the main GC population (e.g. Haşeegan et al. \cite{Hasega05},"
+ Mieske et al. 2006..," Mieske et al. \cite{Mieske06},"
+ Chilingarian et al. 2011..," Chilingarian et al. \cite{Chilin11},"
+ Norris et al. 201I..," Norris et al. \cite{Norris11},"
+ da Rocha et al. 2011)., da Rocha et al. \cite{Daroch11}) ).
+ The error bars of the specific frequencies derived for UCDs suggest ~50 as an upper limit for the importance of the dwarf galaxy channel., The error bars of the specific frequencies derived for UCDs suggest $\sim$ 50 as an upper limit for the importance of the dwarf galaxy channel.
+ We show in Sect., We show in Sect.
+ 4. that this would require at least > of primordial dwarf galaxies in the central - 50-70 kpe of the considered galaxy clusters to have already been disrupted., \ref{progen} that this would require at least $\gtrsim$ of primordial dwarf galaxies in the central $\sim$ 50-70 kpc of the considered galaxy clusters to have already been disrupted.
+ If indeed a significant amount of UCDs were formed from tidal processes. then the tidal stripping of stars from low-mass dark matter halos in the central Fornax cluster has been extremely We conclude that the number counts of UCDs are fully consistent with them being the bright tail of the GC population.," If indeed a significant amount of UCDs were formed from tidal processes, then the tidal stripping of stars from low-mass dark matter halos in the central Fornax cluster has been extremely We conclude that the number counts of UCDs are fully consistent with them being the bright tail of the GC population."
+ From a statistical point of view there Is no need to invoke an additional formation. channel., From a statistical point of view there is no need to invoke an additional formation channel.
+ The statistical. error. bars constrain the fraction of tidally stripped dwarfs to not more than of UCDs., The statistical error bars constrain the fraction of tidally stripped dwarfs to not more than of UCDs.
+ We thank the anonymous referee for very useful comments and suggestions that helped to improve the paper., We thank the anonymous referee for very useful comments and suggestions that helped to improve the paper.
+ We thank Bruno Henriques for providing us with details on the number fraction of disrupted satellite galaxies in their models., We thank Bruno Henriques for providing us with details on the number fraction of disrupted satellite galaxies in their models.
+ IM acknowledges support through DFG grant BEI091/13-1., IM acknowledges support through DFG grant BE1091/13-1.
+The PNe whose medium-resolution spectra are presented here are selected form the bright end of the planetary nebula luminosity function (PNLF) of the M81 disk.,The PNe whose medium-resolution spectra are presented here are selected form the bright end of the planetary nebula luminosity function (PNLF) of the M81 disk.
+" They have two peculiarities, as a group: first, we do not detect the 44686 emission lines in any of them; this is a signature of moderately low central star temperature (less than ~ 100,000 K)."," They have two peculiarities, as a group: first, we do not detect the $\lambda$ 4686 emission lines in any of them; this is a signature of moderately low central star temperature (less than $\sim$ 100,000 K)."
+" Second, none of the PNe whose abundance analysis was performed (i.e., whose electron temperature has been measured from the emission lines) is of Type L thus they must be the progeny of stars with turnoff masses (Μιο) in the-1-2 Mo range."," Second, none of the PNe whose abundance analysis was performed (i.e., whose electron temperature has been measured from the emission lines) is of Type I, thus they must be the progeny of stars with turnoff masses $_{\rm to}$ ) in $\sim$ 1-2 $_{\odot}$ range."
+ The properties of the M81 PN sample are important to define their nature when we compare the population of M81 PNe to those of other galaxies., The properties of the M81 PN sample are important to define their nature when we compare the population of M81 PNe to those of other galaxies.
+ M81 has similar metallicity to that of the Milky Way (Davidge 2009)., M81 has similar metallicity to that of the Milky Way (Davidge 2009).
+" In M81, as in most external galaxies, the observed PN population belong to the brightest few magnitude bins."," In M81, as in most external galaxies, the observed PN population belong to the brightest few magnitude bins."
+ The nature of the brightest PNe varies notably with the metallicity of the host galaxy., The nature of the brightest PNe varies notably with the metallicity of the host galaxy.
+" In their runaway evolution from the AGB, the PN central stars (CS) reach hot temperature while keeping very bright; later, both brightness and temperature decline."," In their runaway evolution from the AGB, the PN central stars (CS) reach hot temperature while keeping very bright; later, both brightness and temperature decline."
+" A PN is visible only when the CS temperature is high enough to ionize hydrogen, and the nebula is optically thin to the rradiation (e.g., Stanghellini Renzini 2000)."," A PN is visible only when the CS temperature is high enough to ionize hydrogen, and the nebula is optically thin to the radiation (e.g., Stanghellini Renzini 2000)."
+" At relatively high metallicities such as those of M81 or the Galaxy, the brightest PNe are not those with the hottest CSs (e.g., Villaver et al."," At relatively high metallicities such as those of M81 or the Galaxy, the brightest PNe are not those with the hottest CSs (e.g., Villaver et al."
+ 2002)., 2002).
+" In fact, the post-AGB shells with very hot CSs are statistically still enshrouded in dust at early phases or their evolution, thus still thick to the rradiation."," In fact, the post-AGB shells with very hot CSs are statistically still enshrouded in dust at early phases or their evolution, thus still thick to the radiation."
+ A lower metallicity implies a shorter thinning time (Kàuufl et al., A lower metallicity implies a shorter thinning time (Käuufl et al.
+ 1993)., 1993).
+" Plenty of Type I PNe have been observed within the bright end of the PNLF of low metallicity galaxies such as M33, the LMC, and the SMC; this is not the case in M81, where the high metallicity prevents massive progenitor PNe (i.e., the Type I) to populate the bright end of the PNLF."," Plenty of Type I PNe have been observed within the bright end of the PNLF of low metallicity galaxies such as M33, the LMC, and the SMC; this is not the case in M81, where the high metallicity prevents massive progenitor PNe (i.e., the Type I) to populate the bright end of the PNLF."
+" To illustrate this point, in Figure 4 we show the luminosity function (PNLF) of Galactic PNe in the Hf--equivalent absolute magnitude."," To illustrate this point, in Figure 4 we show the luminosity function (PNLF) of Galactic PNe in the -equivalent absolute magnitude."
+ The ffluxes are from Cahn et al. (, The fluxes are from Cahn et al. (
+"1992), and the PN distances from Stanghellini et al.","1992), and the PN distances from Stanghellini et al."
+ 2008 (see also Stanghellini Haywood 2010)., 2008 (see also Stanghellini Haywood 2010).
+" The solid histogram represents the homogeneous Galactic sample of 331 PNe where both distance and fflux are known, and where spectroscopy has been performed and a lower limit or a detection of the 44686 line has been acquired."," The solid histogram represents the homogeneous Galactic sample of 331 PNe where both distance and flux are known, and where spectroscopy has been performed and a lower limit or a detection of the $\lambda$ 4686 line has been acquired."
+ The shaded histogram is the subsample of Galactic PNe with Ligsg/Ig< 5%., The shaded histogram is the subsample of Galactic PNe with $_{4686}$ $_{\beta} <$ $\%$.
+" It is clear that at least 80-90% of the bright Galactic PNe do not show the llines, in qualitative agreement with not finding this emission line in the M81 PN sample."," It is clear that at least $\%$ of the bright Galactic PNe do not show the lines, in qualitative agreement with not finding this emission line in the M81 PN sample."
+ In order to quantify the expectation of detecting the eemission line in a given PN population of a certain metallicity we use Káuufl et al, In order to quantify the expectation of detecting the emission line in a given PN population of a certain metallicity we use Käuufl et al.
+"s (1993) models to derive PN opacity, where the non-evolving term of the opacity,r;(0), scales directly with the dust-to-gas ratio (see their Eq.","'s (1993) models to derive PN opacity, where the non-evolving term of the $\tau_{\lambda}(0)$, scales directly with the dust-to-gas ratio (see their Eq."
+ 10)., 10).
+" We also express this ratio as expected at different metallicities, Ma/M,~0.01(0/H)/(O/H)uw (Draine 2009), where the right hand side term refers to the ratio of the oxygen vs. hydrogen abundance in a given galaxy compared to that of the Milky Way."," We also express this ratio as expected at different metallicities, $_{\rm d}$ $_{\rm g}\sim0.01$ $_{\rm MW}$ (Draine 2009), where the right hand side term refers to the ratio of the oxygen vs. hydrogen abundance in a given galaxy compared to that of the Milky Way."
+" To obtain the ratio for M81 scaled to that of the Galaxy we use the rregion oxygen abundances in M81 (4.99x10-*) and in the Galaxy (5.25x10-, Deharveng et al."," To obtain the ratio for M81 scaled to that of the Galaxy we use the region oxygen abundances in M81 $\times$ $^{-4}$ ) and in the Galaxy $\times10^{-4}$, Deharveng et al."
+" 2000) to find Μα/Με~0.01, thus for a given type of M81 PN the thinning time would be similar to that of the Galactic homologous."," 2000) to find ${\rm M_d/M_g \sim 0.01}$, thus for a given type of M81 PN the thinning time would be similar to that of the Galactic homologous."
+" We conclude that the lack of u--emitting PNe in M81 is totally expected, given the Galactic PNLF for both m-—emitting and non-emitting PNe (Figure 4)."," We conclude that the lack of -emitting PNe in M81 is totally expected, given the Galactic PNLF for both –emitting and non-emitting PNe (Figure 4)."
+ It is interesting to compare the PN population of M81 with that of M33 as well., It is interesting to compare the PN population of M81 with that of M33 as well.
+ Magrini et al. (, Magrini et al. (
+"2010) obtained an average oxygen abundance of <O/H>=2.04x10~*for rregions in the M33 disk, corresponding to a dust-to-mass ratio of~4x1073, or, less than half that of M81.","2010) obtained an average oxygen abundance of $<$ $>$ $\times 10^{-4}$ for regions in the M33 disk, corresponding to a dust-to-mass ratio $\sim$ $\times$ $^{-3}$, or, less than half that of M81."
+ By running the Kauufl et al., By running the Käuufl et al.
+"'s models for M33 we obtain (Balmer line) thinning times~1.5 times shorter than in M81 or the Galaxy, which can account for the relatively higher frequency of n--emitting, bright PNe that where observed in M33 by MO9.","'s models for M33 we obtain (Balmer line) thinning $\sim$ 1.5 times shorter than in M81 or the Galaxy, which can account for the relatively higher frequency of -emitting, bright PNe that where observed in M33 by M09."
+" We know that the brightest PNe in a given population excludes the progeny of the massive AGB stars, i.e., those with Mio>2 Mo (Stanghellini Renzini 2000)."," We know that the brightest PNe in a given population excludes the progeny of the massive AGB stars, i.e., those with $_{\rm to}>2$ $_{\odot}$ (Stanghellini Renzini 2000)."
+ The sample of M81 PNe under study here thus does not include the progeny of the more massive AGB stars., The sample of M81 PNe under study here thus does not include the progeny of the more massive AGB stars.
+ This agrees very well with the observational fact that we do not observe any Type I PNe in our M81 sample., This agrees very well with the observational fact that we do not observe any Type I PNe in our M81 sample.
+"Substituting Fo) by στὰ, Eq.","Substituting $F_{\rm{bol}}$ by $ \sigma T_{\rm{eff}}^{4}$ , Eq."
+" 4 can be rewritten as where Ms,=—-2.5log|ZL—FS—x(A)da|+m$. is. computed using∙ the SED of the star at itsvA surface.", \ref{bolo_corre2} can be rewritten as where $\textsf{M}_{S_X}=-2.5 \log \left(\frac{\int_{\lambda_{1}}^{\lambda_{2}}F_{\lambda} S_X(\lambda) d\lambda}{\int_{\lambda_{1}}^{\lambda_{2}} f^{0}_{\lambda} S_X(\lambda) d\lambda}\right) + m^{0}_{S_X}$ is computed using the SED of the star at its surface.
+ Equation. 5 is. similar∙⋅ to the one derived in ? and will be used to compute the bolometric correction in the photometric bands., Equation \ref{bolo_corre3} is similar to the one derived in \cite{Bessel98} and will be used to compute the bolometric correction in the photometric bands.
+" Once we compute the bolometric correction BCs,, the bolometric . ∂↻⊱⊙↕∐↥∁∐↕∂≝∐↥↥∐↺⊜∐↻∘↕∁∂∐↻⊜⊂↥⊜∐⋁⊜⊂⊓↕⊙∐↕⊟∁↥⋅∍⋅⋅ ."," Once we compute the bolometric correction $BC_{S_X}$, the bolometric absolute magnitude $M_{\rm{bol}}$ can be derived from Eq. \ref{bolo_corre1}."
+" Figure 15 and 12 display the bolometric correction for the G, and bbands for Gap»different metallicitiesEN and surface gravities."," Figure \ref{fig:Bolom-plot} and 12 display the bolometric correction for the $G$, and bands for different metallicities and surface gravities."
+"2. Panels (a), (c) and (d) of Fig."," Panels (a), (c) and (d) of Fig."
+ 15 display the bolometric correction and its. dependence with- effective. temperature and metallicity., \ref{fig:Bolom-plot} display the bolometric correction and its dependence with effective temperature and metallicity.
+e. Panel (b) shows the variation of the bolometric correction in G with respect to surface gravity and for solar metallicity., Panel (b) shows the variation of the bolometric correction in $G$ with respect to surface gravity and for solar metallicity.
+ The bolometric corrections in G and aare near zero for F-type stars and for the entire. metallicitye. range., The bolometric corrections in $G$ and are near zero for F-type stars and for the entire metallicity range.
+" For the maximum around BCgp is found =0.75to be related to Gpp,,stars with T,g around 4500—5000 K. For cool stars (logTeg€3.6 dex) and for each temperature, there is a large dispersion in bolometric correction values with respect to surface gravity and metallicity."," For , the maximum around $BC_{RP}=$ 0.75 is found to be related to stars with $T_{\rm{eff}}$ around $4500-5000$ K. For cool stars $\log T_{\rm{eff}}\leq3.6$ dex) and for each temperature, there is a large dispersion in bolometric correction values with respect to surface gravity and metallicity."
+ The extinction curve used in the previous section was taken from ? assuming an average galactic value of Ry=3.1., The extinction curve used in the previous section was taken from \cite{1989ApJ...345..245C} assuming an average galactic value of $R_V=3.1$.
+" This curve agrees with ? and ? in the wavelength range of Gaia’s passbands, 330-1000 nm."," This curve agrees with \cite{1999PASP..111...63F} and \cite{2007ApJ...663..320F} in the wavelength range of 's passbands, 330–1000 nm."
+" In ?,, which contains the most updated discussion on the absorption law, extinction curves with Ry values in the range 2.4—3.6 are considered for a sample 243 stars in sight lines with diffuse interstellar medium."," In \cite{2007ApJ...663..320F}, which contains the most updated discussion on the absorption law, extinction curves with $R_V$ values in the range 2.4–3.6 are considered for a sample 243 stars in sight lines with diffuse interstellar medium."
+" magnitudes have been recomputed forall spectra of the BaSeL3.1 library and Ajzss0=0, 1, 3and5magwithRy and Ry3.6."," magnitudes have been recomputed forall spectra of the BaSeL3.1 library and $=0$, 1, 3 and 5 mag with $R_V=2.4$ and $R_V=3.6$."
+ The left panel of Fig., The left panel of Fig.
+" 16 shows the G—V vs V—I, polynomial relationships for each value of Ry.", \ref{fig:RV_var} shows the $G-V$ vs $V-I_c$ polynomial relationships for each value of $R_V$.
+ No differences are noticeable: the polynomials overlap in the three cases., No differences are noticeable: the polynomials overlap in the three cases.
+" In the right panel of Fig. 16,,"," In the right panel of Fig. \ref{fig:RV_var},"
+ we display the effect of the variation of Ry on the colour-colour diagram., we display the effect of the variation of $R_V$ on the colour-colour diagram.
+" We have computed G—Ggp and G—Grp with respect to Gpgp-Ggp ccolour for an absorption value A;255921 mag, andusingthethreedif ferentvaluesof Ry."," We have computed $G-G_{BP}$ and $G-G_{RP}$ with respect to $-$ colour for an absorption value $=1$ mag, and usingthe three different values of $R_V$ ."
+ The effect of modifying Ry is also negligible inthis case., The effect of modifying $R_V$ is also negligible inthis case.
+ Figure 17 and 13 show several ratios of total-to-selective absorption including absorption, Figure \ref{fig:absorcio1} and 13 show several ratios of total-to-selective absorption including absorption
+found anv set of parameters that left (he simulations just like the Newtonian one. which. il is worth stressing again. is consistent with observations.,"found any set of parameters that left the simulations just like the Newtonian one, which, it is worth stressing again, is consistent with observations."
+ It was just the results of model ID that leacl us to make a pure Molffatian disk: our moclel III whose simulations were made wilh (he same parameters used in the models I ancl lI., It was just the results of model II that lead us to make a pure Moffatian disk; our model III whose simulations were made with the same parameters used in the models I and II.
+ The results of model IHE can be seen in Figures 7- 9.., The results of model III can be seen in Figures \ref{fig7}- \ref{fig9}.
+ In the top right panel of Figure 7.. one finds the phase-space points.," In the top right panel of Figure \ref{fig7}, one finds the phase-space points."
+ These points are distributed somewhat differently. as compared to previous models. because (he rotation curve has now a different shape.," These points are distributed somewhat differently as compared to previous models, because the rotation curve has now a different shape."
+ Also. after 1 Gyr of simulated time. we note that the scaltering in the phase space plane is more prominent than in the model II and the initial configuration 1s completelv lost.," Also, after 1 Gyr of simulated time, we note that the scattering in the phase space plane is more prominent than in the model II and the initial configuration is completely lost."
+ This effect is neither due to simulation inaccuracies nor to the resolution of the simulation. as can be seen that the energy. violation (see the top right panel of Figure 7)) is small. namely. logAPδη —2.0.," This effect is neither due to simulation inaccuracies nor to the resolution of the simulation, as can be seen that the energy violation (see the top right panel of Figure \ref{fig7}) ) is small, namely, $\log{\Delta E/ E_0} \lesssim -2.0$."
+ Our results show in (he end that the exponential profile is not consistent with: the Moffatian potential., Our results show in the end that the exponential profile is not consistent with the Moffatian potential.
+ As occurred with model IL. the rotation curve of the model HI is below (he initial one.," As occurred with model II, the rotation curve of the model III is below the initial one."
+ For E.«10 kpe. the rotation curve is linear ancl associated with the central core movement. similar to a rigid body rotation.," For $\rm{R < 10}$ kpc, the rotation curve is linear and associated with the central core movement, similar to a rigid body rotation."
+ For R.2»10 kpe. the rotation curve is similar (o a (vpical rotation curve of late-tvpe svstems. but. in analogy wilh model II. this similarity is only apparent. as we will see below.," For $\rm{R > 10}$ kpc, the rotation curve is similar to a typical rotation curve of late-type systems, but, in analogy with model II, this similarity is only apparent, as we will see below."
+ Figures 8 and 9 cdisplav (he particles positions at Cr.y) plane. where the simulated lime is indicated in (he respective boxes.," Figures \ref{fig8} and \ref{fig9} display the particle's positions at $(x,y)$ plane, where the simulated time is indicated in the respective boxes."
+ Note that. al initial (mes. the same atypical central void appears wilh its surrounding ring. that. collapses aud fragments in (wo small clumps. that evolves separately eventually merging al /e0.9 Cyr.," Note that, at initial times, the same atypical central void appears with its surrounding ring, that collapses and fragments in two small clumps, that evolves separately eventually merging at $t \sim 0.9$ Gyr."
+ This simulation shows that for a typical initial equilibrium configuration. DMs disk," This simulation shows that for a typical initial equilibrium configuration, BM's disk"
+The ssurface brightness yields an ionizing flux of 6(H)—8.1 x10’ photons cm? s! tthat then leads to a derived ionizing luminosity of Q(H)=2.5x1015 photons !.,The surface brightness yields an ionizing flux of $\phi$ $8.1\times 10^{7}$ photons $^{-2}$ $^{-1}$ that then leads to a derived ionizing luminosity of $2.5\times 10^{49}$ photons $^{-1}$.
+ Photoionization of the large scale Barnard's Loop may arise from multiple stars., Photoionization of the large scale Barnard's Loop may arise from multiple stars.
+" It is unlikely that two candidate stars that are located within the Orion Nebula aand A)) are important contributors to ionizing Barnard's Loop because it is well known (O'Dell2001) that the nebula is optically thick to LyC radiation in all directions except possibly to the southwest, where the foreground Veil thins and X-ray emission from million degree gas is found (Güdeletal. 2008)."," It is unlikely that two candidate stars that are located within the Orion Nebula and ) are important contributors to ionizing Barnard's Loop because it is well known \citep{od01} that the nebula is optically thick to LyC radiation in all directions except possibly to the southwest, where the foreground Veil thins and X-ray emission from million degree gas is found \citep{gud08}."
+".The four remaining candidates for causing photoionization would have a total luminosity of Q(H)=1.9x1015 photons!,, in reasonable agreement with the LyC luminosity derived from the ssurface brightness."," .The four remaining candidates for causing photoionization would have a total luminosity of $1.9\times 10^{49}$ photons, in reasonable agreement with the LyC luminosity derived from the surface brightness."
+ The effective temperatures of these stars range from 30000 K to 32000 K. We have determined the properties of the radiation field illuminating Barnard's Loop by comparing the predictions for emission line ratios from stellar atmosphere models of various effective temperatures (Tstar)., The effective temperatures of these stars range from 30000 K to 32000 K. We have determined the properties of the radiation field illuminating Barnard's Loop by comparing the predictions for emission line ratios from stellar atmosphere models of various effective temperatures ).
+" We use the development version of the spectral simulation code Cloudy, last reviewed by Ferlandetal.(1998) and noted in appendix A. We present here several large grids of model calculations that, in some cases, varied both aand the flux of hydrogen-ionizing photons, and in others, onlyTsar."," We use the development version of the spectral simulation code Cloudy, last reviewed by \citet{fer98} and noted in appendix A. We present here several large grids of model calculations that, in some cases, varied both and the flux of hydrogen-ionizing photons, and in others, only."
+ Later we also calculate and present models including differences in the abundance ratio (Z/H) of heavy elements to hydrogen., Later we also calculate and present models including differences in the abundance ratio (Z/H) of heavy elements to hydrogen.
+" In many ways our calculations are similar to those of but we use a version of Cloudy with up-to-date atomic coefficients, more recent stellar"," In many ways our calculations are similar to those of \citet{sem00} but we use a version of Cloudy with up-to-date atomic coefficients, more recent stellar"
+(his profile of dpwag/di in TLUSTY by mapping it directly onto the column density grid used in the code.,this profile of $dp_{\rm mag}/dm$ in TLUSTY by mapping it directly onto the column density grid used in the code.
+ In order to avoid Che upper vertical boundary effects on (the simulation data at low column densities (see the kinks in Figure 12 at mi=0.1 gem 7). we extrapolate ΠΠdin as à constant for in«0.1 g 37.," In order to avoid the upper vertical boundary effects on the simulation data at low column densities (see the kinks in Figure \ref{figdpdmfrac} at $m=0.1$ g $^{-2}$ ), we extrapolate $dp_{\rm mag}/dm$ as a constant for $m<0.1$ g $^{-2}$."
+ Figure 13. compares (he emergent spectrum from (he annulus when magnetic pressure support is or is not included in the vertical hydrostatic balance., Figure \ref{figpmagspec} compares the emergent spectrum from the annulus when magnetic pressure support is or is not included in the vertical hydrostatic balance.
+ In contrast to ehanging the dissipation profile (Figure 10)). magnetic pressure support produces a much more noticeable change. aud causes the spectrum to be significantly. harder.," In contrast to changing the dissipation profile (Figure \ref{figdissspec}) ), magnetic pressure support produces a much more noticeable change, and causes the spectrum to be significantly harder."
+ The reason for this can be seen in Figure 14.. which shows the vertical clensity and temperature profiles in (he (wo structures. as well as (he locations of the effective photospheres al lrequencies corresponding to (he (wo spectral peaks.," The reason for this can be seen in Figure \ref{figpmagvertstruc}, which shows the vertical density and temperature profiles in the two structures, as well as the locations of the effective photospheres at frequencies corresponding to the two spectral peaks."
+ Without magnetic pressure support. large gas pressure gradients are required (o support the disk against the high gravitw al hieh alüitude.," Without magnetic pressure support, large gas pressure gradients are required to support the disk against the high gravity at high altitude."
+ These necessarily require steep densitv gradients because the temperature profile (largely set by radiative equilibrium) is too flat., These necessarily require steep density gradients because the temperature profile (largely set by radiative equilibrium) is too flat.
+ Magnetic pressure support relaxes this constraint. permitting a much more extended densitv profile.," Magnetic pressure support relaxes this constraint, permitting a much more extended density profile."
+ Because the location οἱ the effective photosphere is set by an approximately fixed column density. the larger density scale height in a magnetically supported annulus implies that the density at the effective photosphere will be lower than would be the case if magnetic pressure was neglected.," Because the location of the effective photosphere is set by an approximately fixed column density, the larger density scale height in a magnetically supported annulus implies that the density at the effective photosphere will be lower than would be the case if magnetic pressure was neglected."
+ The lower densitv reduces (he ratio of absorption to scattering opacitv. which alone would tend to harden the spectrum by producing a modified blackbody.," The lower density reduces the ratio of absorption to scattering opacity, which alone would tend to harden the spectrum by producing a modified blackbody."
+ In addition. the lower density implies that metals are more ionized than they would otherwise be. and this reduces the strength of absorption edges.," In addition, the lower density implies that metals are more ionized than they would otherwise be, and this reduces the strength of absorption edges."
+ This is illustrated in detail in Figure 15.. which shows the departure coellicient lor (he CVI eround state. which is responsible for the absorption edee feature al 0.5 keV in the unmagnetized spectrum.," This is illustrated in detail in Figure \ref{figdepcoeff}, which shows the non-LTE departure coefficient for the CVI ground state, which is responsible for the absorption edge feature at $~0.5$ keV in the unmagnetized spectrum."
+ In the magnetized atmosphere. non-LTE elfects reduce the ground state population of this ion even further. ancl drive the edge into emission.," In the magnetized atmosphere, non-LTE effects reduce the ground state population of this ion even further, and drive the edge into emission."
+ For all these reasons. we expect alinosphere caleulations which include magnetic pressure support will generically result in harder spectra.," For all these reasons, we expect atmosphere calculations which include magnetic pressure support will generically result in harder spectra."
+ This applies to electrically conducting accretion disks with MBRI turbulence in all astroplvsical contexts. from cataclvsmic variables to active galactic nuclei.," This applies to electrically conducting accretion disks with MRI turbulence in all astrophysical contexts, from cataclysmic variables to active galactic nuclei."
+ Figure 14. also compares the time ancl horizontally averaged density ancl temperature structures from the Iirose.IXrolik.&Stone(2005) simulation with those computed by TLUSTY.," Figure \ref{figpmagvertstruc} also compares the time and horizontally averaged density and temperature structures from the \citet{hir05}
+ simulation with those computed by TLUSTY."
+ The density profiles are in excellent agreement. and even the temperatures compare favorably beneath the effective photosphere.," The density profiles are in excellent agreement, and even the temperatures compare favorably beneath the effective photosphere."
+It can be seen in relfig:vep.. that the field D. has an approximately Gaussian distribution even well above the photospheric level (2=0).,"It can be seen in \\ref{fig:vgp}, that the field $B_z$ has an approximately Gaussian distribution even well above the photospheric level $z=0$ )."
+ At the bottom. the Gaussian component has a standard deviation. oi~11.3 Gauss. comparable to the intrinsic noise in a [ive-minute average of MDI magnetograms (LinandNorton2001): it decreases with height as the width of the Gaussian narrows.," At the bottom, the Gaussian component has a standard deviation, $\sigma_z^{(g)}\simeq11.3$ Gauss, comparable to the intrinsic noise in a five-minute average of MDI magnetograms \citep{Liu2001}; it decreases with height as the width of the Gaussian narrows."
+ The actual distribution departs from the Gaussian fit principally by an excess (Lail)⋅⋡ mostly outside⋅ zo.(ο), The actual distribution departs from the Gaussian fit principally by an excess (tail) mostly outside $\pm\sigma_z^{(g)}$.
+ Histograms⋅ made from. the horizontal. components of. the extrapolated fields at varving heights look very similar to those of the vertical field shown in re[fie:vgp.., Histograms made from the horizontal components of the extrapolated fields at varying heights look very similar to those of the vertical field shown in \\ref{fig:vgp}.
+ The Gaussian fit to the central core of each distribution can be used (o equantilv the degree (o which the actual fiekl is a Gaussian variable., The Gaussian fit to the central core of each distribution can be used to quantify the degree to which the actual field is a Gaussian variable.
+ The aimplituce of the Gaussian component ranges rom 0.97 al z=0 to 0.31 higher up (see relfig:sigbb)., The amplitude of the Gaussian component ranges from $0.97$ at $z=0$ to $0.81$ higher up (see \\ref{fig:sig}b b).
+ That is to sav the field is between 81% and 97%. Gaussian., That is to say the field is between $81\%$ and $97\%$ Gaussian.
+ While a significant component of the magnetic fiekl has a Gaussian distribution. it is nol spatial white noise.," While a significant component of the magnetic field has a Gaussian distribution, it is not spatial white noise."
+ Longcope. Brown and Priest (LDD03)) show that a potential field extrapolated [rom a magnetoeram will isotropic power spectral densily (psd) SA) has vertical field with variance hdk = 292. = σοι9 ," Longcope, Brown and Priest ) \nocite{Longcope2003b} show that a potential field extrapolated from a magnetogram with isotropic power spectral density (psd) $S(k)$ has vertical field with variance _z^2(z) = dk = _x^2(z) = _y^2(z) ."
+"Spatial white noise is uncorrelated [rom pixel to pixel and therefore has a flat. psd (I,−hr)"," Spatial white noise is uncorrelated from pixel to pixel and therefore has a flat psd (k_c-k) ,"
+The complex cluster morphology seen in N-ray images ancl in galaxy distributions gave support to the hypothesis that large structures form hierarchically: that is. (hat small groups ol galaxies merge to form low-mass subclusters. which then merge to form a massive rich cluster with the infalling groups aligned along laree filaments.,"The complex cluster morphology seen in X-ray images and in galaxy distributions gave support to the hypothesis that large structures form hierarchically; that is, that small groups of galaxies merge to form low-mass subclusters, which then merge to form a massive rich cluster with the infalling groups aligned along large filaments."
+ Galaxies and eroups are the building blocks of the observable Universe and contain the bulk of the observable barvyons., Galaxies and groups are the building blocks of the observable Universe and contain the bulk of the observable baryons.
+ Groups are estimated to contain a significant fraction.20-3056.. of the total matter in the Universe.," Groups are estimated to contain a significant fraction, of the total matter in the Universe."
+ Thus groups are important. cosmological indicators of the distribution aud properties of the dark matter., Thus groups are important cosmological indicators of the distribution and properties of the dark matter.
+ However. because thev are not as luminous as clusters. they have received less study than (heir more massive cousins.," However, because they are not as luminous as clusters, they have received less study than their more massive cousins."
+ Observations of rich clusters show many examples of pending or ongoing mergers of subcelusters., Observations of rich clusters show many examples of pending or ongoing mergers of subclusters.
+ In particular in X-ray cluster catalogs. about of rich clusters show substructure(Jones&Forman1984.1999:Mohrefaf1995).," In particular in X-ray cluster catalogs, about of rich clusters show substructure\citep{jon84, jon99, moh95}."
+.. Virtually all stages of cluster mergers have been thoroughly investigated with both theChandra and NMM-Newton observatories (Markeviteh.efaf.2000:Vikhlinin.Markevitch.&Murray2001:Briel.Finoguenov.Tlenry2004:Henry.Finoguenov.&Briel 2004).," Virtually all stages of cluster mergers have been thoroughly investigated with both the and XMM-Newton observatories \citep{mar00,vik01,bri04,hen04}."
+. Llowever. less attention has been paid to the merging of groups and the formation of low-mass clusters due to their lower X-ray Iuminosity and the paucity of examples in the local Universe.," However, less attention has been paid to the merging of groups and the formation of low-mass clusters due to their lower X-ray luminosity and the paucity of examples in the local Universe."
+ To our knowledge. the only nearby example of the early stages of the merger of (wo roughly equal mass groups is the NGC 499/NGC 507 pair (xin&Fabbiano1995;Kraltοἱaf2003).," To our knowledge, the only nearby example of the early stages of the merger of two roughly equal mass groups is the NGC 499/NGC 507 pair \citep{kim95,kra03}."
+. In the hierarchical scenario. group mergers represent a critical transitional phase in the formation of larger scale structure.," In the hierarchical scenario, group mergers represent a critical transitional phase in the formation of larger scale structure."
+ An understanding of the group merger process is (hus fundamental to understanding the growth ol structure., An understanding of the group merger process is thus fundamental to understanding the growth of structure.
+" In this paper. we report results from analvsis of a short Chandra/ACIS-S observation of the nearby elliptical galaxy NGC 7618 (2=0.017309 or d; 274.1 Mpe for WALAP cosmology (Spergelefal.2003) - 1""—350 pc)."," In this paper, we report results from analysis of a short /ACIS-S observation of the nearby elliptical galaxy NGC 7618 $z$ =0.017309 or $_L$ =74.1 Mpc for WMAP cosmology \citep{spe03} - $''$ =350 pc)."
+" The X-ray huninosity of NGC 7618 is ~7x LOP eres ! in the 0.1-10 keV αμα, typical of groups. not isolated elliptical galaxies. although it appears to be optically isolated (Colbert.Mulehaey.anclZabhicdoll 2001).."," The X-ray luminosity of NGC 7618 is $\sim$ $\times$ $^{42}$ ergs $^{-1}$ in the 0.1-10 keV band, typical of groups, not isolated elliptical galaxies, although it appears to be optically isolated \citep{col01}. ."
+ We find a sharp surface, We find a sharp surface
+For the simple case of 2 clumps colliding. the velocity dispersion increases with size scale approximately as a power law. independent. of the impact parameter.,"For the simple case of 2 clumps colliding, the velocity dispersion increases with size scale approximately as a power law, independent of the impact parameter."
+ Increasing the impact parameter transposes the power law to smaller size scales. since the size of the shocked. region decreases.," Increasing the impact parameter transposes the power law to smaller size scales, since the size of the shocked region decreases."
+ For size-scales Tr6ον the velocity. dispersion. includes. eas which does not enter the shock and the velocity size-scale relation Uattens.," For size-scales $>r-b/2$, the velocity dispersion includes gas which does not enter the shock and the velocity size-scale relation flattens."
+" The power law for the shocked gas is somewhat steeper than the 0xricon"""" pelation.. observed. rather aoxrd."," The power law for the shocked gas is somewhat steeper than the $\sigma \propto r^{0.5}$ relation observed, rather $\sigma~\propto~r^1$."
+ However. the power law exponent may x expected to decrease. when considering larger regions ancl structure on multiple scales composed of many clumps (Section 4.2).," However, the power law exponent may be expected to decrease when considering larger regions and structure on multiple scales composed of many clumps (Section 4.2)."
+ This power law is also independent of the comparative radii of the 2 spheres., This power law is also independent of the comparative radii of the 2 spheres.
+ Fig., Fig.
+ 16 shows the οσοι of varying he racius of the second. clump (r2)., 16 shows the effect of varying the radius of the second clump $r_{2}$ ).
+ Similarly to varving he impact parameter. the velocity size-scale relation shifts o smaller scales as re decreases. and levels olf at size-cales comparable to r».," Similarly to varying the impact parameter, the velocity size-scale relation shifts to smaller scales as $r_{2}$ decreases, and levels off at size-scales comparable to $r_{2}$."
+ Scaling both clump radii up or down will extend the velocity. size-scale relation to larger or smaller size-scales., Scaling both clump radii up or down will extend the velocity size-scale relation to larger or smaller size-scales.
+ For example. if the radius decreases by a factor of 10. the column density and therefore masseweighted velocity dispersion will also decrease by the same factor.," For example, if the radius decreases by a factor of 10, the column density and therefore mass-weighted velocity dispersion will also decrease by the same factor."
+ Assuming that the clumps have velocities ey=1. the velocity. dispersion. for an impact parameter of bz0.5 reaches a maximum of ~ 0.3.," Assuming that the clumps have velocities $v_0=1$, the velocity dispersion for an impact parameter of $b\ge 0.5$ reaches a maximum of $\thicksim$ 0.3."
+ Since the initial velocity ο can be scaled up or down. the maximum velocity dispersion for an initial velocity of voc will be approximately 0.32 ουὃς (e.g. 3 ος iM= 20).," Since the initial velocity $v$ can be scaled up or down, the maximum velocity dispersion for an initial velocity of $v_0 c_s$ will be approximately 0.3 $v_0 c_s$ (e.g. 3 $c_s$ if ${\cal M}=20$ )."
+ When the impact parameter is small. the velocity dispersion is unlikely to be supersonic (e.g. the maximum is 0.5 ὃς when 6=0.1 and A4= 20).," When the impact parameter is small, the velocity dispersion is unlikely to be supersonic (e.g. the maximum is $\thicksim 0.5$ $c_s$ when $b=0.1$ and ${\cal M}=20$ )."
+ This is expected. since if b=0 the clumps collide head on ancl the velocity dispersion is O evervwhere.," This is expected, since if $b=0$ the clumps collide head on and the velocity dispersion is 0 everywhere."
+ The velocity. size-scale relation obtained for 2) clumps, The velocity size-scale relation obtained for 2 clumps
+αἲ.|? Since equatious 9 and LL are imulti-valued. we restrict 4 to be on the range |0.2).,", Since equations \ref{eq:angle1} and \ref{eq:angle2} are multi-valued, we restrict $\chi$ to be on the range $[0,\pi)$."
+ The special case where «=0 and 4 remains finite is very müportant to the analysis that follows., The special case where $u=0$ and $q$ remains finite is very important to the analysis that follows.
+ For this case. the polarization angle is given by 4=0 when g20 and 4=7/2 when 4<0.," For this case, the polarization angle is given by $\chi=0$ when $q>0$ and $\chi=\pi/2$ when $q<0$."
+ We proceed to calculate the inteusitv of the euission. which requires the extinction cross-section fron equation 3.11 of LDsh: ipm where C; aud oj are defined as Note that àj docs not vary with position. sce we asmue that AR; aud Fy are coustauts for cach eran species.," We proceed to calculate the intensity of the emission, which requires the extinction cross-section from equation 3.14 of LD85: ) ], where $\Cp$ and $\alphaj$ are defined as = = Note that $\alphaj$ does not vary with position, since we assume that $R_j$ and $F_j$ are constants for each grain species."
+ The coutribution of eram species j to the intensity J is eiven by where Mods the surface density aud X» is a related quantity defined by <= pds," The contribution of grain species $j$ to the intensity $I$ is given by c_j ), where $\Sigma$ is the surface density and $\Sigma_2$ is a related quantity defined by = ds"
+et al. 1992)).,et al. \cite{pin92}) ).
+ Ileuce. it is of cousiderable interest to deteriunue the masses of the stars.," Hence, it is of considerable interest to determine the masses of the stars."
+" This can be done by comparing aand absolute magnitude, AZ. with mass tracks from stellar evolution calculations."," This can be done by comparing and absolute magnitude, $M_V$, with mass tracks from stellar evolution calculations."
+ Using the apparent magnitudes eiven in Table 1 aud parallaxes frou The Uipparcos and Tycho Catalogues (ESA 1997)) the absolute magnitudes are calculated (Table 6). and the stars are plotted in the loeTig-Asy- diagram (Fig. 8)).," Using the apparent magnitudes given in Table 1 and parallaxes from The Hipparcos and Tycho Catalogues (ESA \cite{esa97}) ) the absolute magnitudes are calculated (Table 6), and the stars are plotted in the $M_V$ diagram (Fig. \ref{fig.8}) )."
+ The mass tracks shown are from the new. a-element euhanced. evolutionary inodels of VandenBere et al. (19993).," The mass tracks shown are from the new, $\alpha$ -element enhanced, evolutionary models of VandenBerg et al. \cite{van99}) )."
+ Tuterpolation between the mass tracks (taking iuto account their dependence on Fe/II) leads to the masses 8eiveu in Table 6. aud from the corresponding isochrones the stellar ages given are obtained.," Interpolation between the mass tracks (taking into account their dependence on ) leads to the masses given in Table 6, and from the corresponding isochrones the stellar ages given are obtained."
+ The errors of the masses aud ages giveu in Table 6 are standard errors corresponding to the adopted errors of aand My, The errors of the masses and ages given in Table 6 are standard errors corresponding to the adopted errors of and $M_V$.
+ Additional errors may be present due to inadequate stellar models auc uucertaiuties in the Cibration of aud the boloiietrie correction., Additional errors may be present due to inadequate stellar models and uncertainties in the calibration of and the bolometric correction.
+ Such errors are. however. nore systematic and are expected to affect all stars with about the same amount.," Such errors are, however, more systematic and are expected to affect all stars with about the same amount."
+ Heuce. we couclude from Table G that the two stars for which thas been detected aud 130551) have sienificautly higher masses than the three stars with no ypresent in their atmospheres.," Hence, we conclude from Table 6 that the two stars for which has been detected and ) have significantly higher masses than the three stars with no present in their atmospheres."
+" This makes seuse. because he depth of the convection zouc of a star ou the main sequence decreases rapidly as a fiction of increasiug mass,"," This makes sense, because the depth of the convection zone of a star on the main sequence decreases rapidly as a function of increasing mass."
+ Hence. according to standard stellar models without uuxing. the depletion of lis less severe iu the more massive stars.," Hence, according to standard stellar models without mixing, the depletion of is less severe in the more massive stars."
+ Iu this connection we note that although Lis already on the subeiaut branch and the coolest of the stars. it las speut most of its life as a main sequence star at Tig26300 Ts. Tuterpretation of the novel result of this paper - the detection aud quantitative. iieasurenment of the ?Li abuudance in two old metal-poor disk stars - is continecut onu two factors: (1) the expected evolution of the interstellar ΟΤΑ abuudauce with metallicity. and (3) the depletion of the stellar Li abundance by the couvective mixing thet occurs iu the pre-main sequence phase. aud the additional depletion occurring ou the main sequence.," In this connection we note that although is already on the subgiant branch and the coolest of the stars, it has spent most of its life as a main sequence star at $\teff \simeq 6300$ K. Interpretation of the novel result of this paper - the detection and quantitative measurement of the $^6$ Li abundance in two old metal-poor disk stars - is contingent on two factors: (i) the expected evolution of the interstellar $^6$ Li abundance with metallicity, and (ii) the depletion of the stellar $^6$ Li abundance by the convective mixing that occurs in the pre-main sequence phase, and the additional depletion occurring on the main sequence."
+ As is all too wel known. prediction of Li depletion by mein sequence stars and subeiauts is an duprecise art.," As is all too well known, prediction of Li depletion by main sequence stars and subgiants is an imprecise art."
+ Standard models bv Pinsouneault et al. (1992)), Standard models by Pinsonneault et al. \cite{pin92}) )
+ predict loss of lithinu iu the pre-main sequence phase and no subsequent loss for stars of the mass of our «quiutet., predict loss of lithium in the pre-main sequence phase and no subsequent loss for stars of the mass of our quintet.
+ Depletious for masses of up to O.85M.. andl iietallicities. corresponding to |Fe/I]| = —2.6 aud 1.6 are computed by Piusouncault et al., Depletions for masses of up to ${\cal M}_{\sun}$ and metallicities corresponding to [Fe/H] = $-2.6$ and $-1.6$ are computed by Pinsonneault et al.
+ For their 0.85: , For their ${\cal M}_{\sun}$ 
+Table 2. summarizes the effect of adding gravitational radiation.,Table \ref{coreresults} summarizes the effect of adding gravitational radiation.
+ In general the effect is ereater al higher masses because gravitational radiation is stronger for à given orbit., In general the effect is greater at higher masses because gravitational radiation is stronger for a given orbit.
+ Because ol the extra energv sink. the binaries merge with fewer encounters. fewer black holes are ejected. and the fraction of sequences in which a binary is ejected is smaller.," Because of the extra energy sink, the binaries merge with fewer encounters, fewer black holes are ejected, and the fraction of sequences in which a binary is ejected is smaller."
+ The most dramatic change is in the duration of the sequence. which gravitational radiation reduces by to40%.," The most dramatic change is in the duration of the sequence, which gravitational radiation reduces by to."
+.. The distributions of final semimajor axes (Fig. 2)), The distributions of final semimajor axes (Fig. \ref{ahisto}) )
+ and final eccentricities (Fig. 3)), and final eccentricities (Fig. \ref{ehisto}) )
+ have similar shapes to the Newtonian only distributions., have similar shapes to the Newtonian only distributions.
+ Due to the circularizing effect of gravitational radiation. binaries of all mass ratios merge with a smaller (67) than Newionian only sequences with the largest difference αἱ high mass ratios.," Due to the circularizing effect of gravitational radiation, binaries of all mass ratios merge with a smaller $\left<e_{f}\right>$ than Newtonian only sequences with the largest difference at high mass ratios."
+ Guavitational radiation also produces a smaller (ap) for mop2300M..., Gravitational radiation also produces a smaller $\left<a_{f}\right>$ for $m_{0} \ga 300~\msun$.
+ This can be seen in Figure 2. where (he gravitational radiation simulations displav an excess number of sequences with low ay. which is a consequence of the binaries’ lower e.," This can be seen in Figure \ref{ahisto}
+ where the gravitational radiation simulations display an excess number of sequences with low $a_{f}$, which is a consequence of the binaries' lower $e_{f}$ ."
+ We can use (hese simulations to test the Miller&Hamilton(20022). model of IAIBIL formation., We can use these simulations to test the \citet{mh02} model of IMBH formation.
+ We assume that a 50M. seed black hole with a 10AM. companion will uidergo repeated. (hree-bocly encounters with 10M. interloping black holes in a globular cluster with eus—50kms.| and η=10?pe7., We assume that a $50~\msun$ seed black hole with a $10~\msun$ companion will undergo repeated three-body encounters with $10~\msun$ interloping black holes in a globular cluster with $v_{\mathrm{esc}}=50\kms$ and $n=10^{5}~\mathrm{pc}^{-3}$.
+ We also assume that the density of the cluster core remains constant as the IMDBII grows., We also assume that the density of the cluster core remains constant as the IMBH grows.
+ We then test whether (he model of οιw build up to IMDII masses. which we take to be 10*M... 1) without ejecting too manv black holes [rom the cluster. 2) without ejecting the IMDBII from the cluster. and 3) within (he lifetime of the globular cluster.," We then test whether the model of \citet{mh02} can build up to IMBH masses, which we take to be $10^{3}~\msun$, 1) without ejecting too many black holes from the cluster, 2) without ejecting the IMBH from the cluster, and 3) within the lifetime of the globular cluster."
+ We also test how these depend on escape velocity and seed mass., We also test how these depend on escape velocity and seed mass.
+ If the number of black holes ejected is greater than the total munber of black holes in the cluster core. then the IMDII cannot buikl up to the required mass by accreting black holes alone.," If the number of black holes ejected is greater than the total number of black holes in the cluster core, then the IMBH cannot build up to the required mass by accreting black holes alone."
+" To caleulate the total number of black holes ejected while building up to large Inasses, we sum (he average number of ejections using a linear interpolation of the values in Table 2.."," To calculate the total number of black holes ejected while building up to large masses, we sum the average number of ejections using a linear interpolation of the values in Table \ref{coreresults}."
+" Assuming a cluster escape velocity of v,=50kms! we find that the total number of black holes ejected when building up to 1000AL. is approximately. 6800 [or our Newtonian only and 5300 for gravitational radiation simulations."," Assuming a cluster escape velocity of $v_{\mathrm{esc}} = 50\kms$, we find that the total number of black holes ejected when building up to $1000~\msun$ is approximately 6800 for our Newtonian only and 5300 for gravitational radiation simulations."
+ This is far greater than the estimated 107 to 10* black holes available (PortegiesZwart&MeMillan2000)., This is far greater than the estimated $10^{2}$ to $10^{3}$ black holes available \citep{pzm00}.
+. Hf there were initially one thousand 10AL. black holes in the cluster. mergers of the massive black hole with a series of LOAL. black holes would exhaust half of the black holes in ~2.6x105vr and would ultimately. produce a 240AL. black hole.," If there were initially one thousand $10~\msun$ black holes in the cluster, mergers of the massive black hole with a series of $10~\msun$ black holes would exhaust half of the black holes in $\sim 2.6 \times
+10^{8}~\mathrm{yr}$ and would ultimately produce a $240~\msun$ black hole."
+ Increasing (he seed mass increases the final mass of the IMDBII when half of the field black holes run out., Increasing the seed mass increases the final mass of the IMBH when half of the field black holes run out.
+ If theseed mass were, If theseed mass were
+complex DNO behaviour that. occurs near the end. of outburst. during the final transition to quiescence. where frequeney doubling and tripling are seen (Paper IV).,"complex DNO behaviour that occurs near the end of outburst, during the final transition to quiescence, where frequency doubling and tripling are seen (Paper IV)."
+ All standard. photometric data reductions were performed inIRAL?.. including gain corrections of the different amplifiers (SALTICAAL has two CCDs with two amplifiers each). merging of the cdillerent amplifiers. a superllat correction to remove the ellects of vignetting bv the spherical abberation corrector in slot-mocde observations. and aperture photometry usingDAODPIIOT.," All standard photometric data reductions were performed in, including gain corrections of the different amplifiers (SALTICAM has two CCDs with two amplifiers each), merging of the different amplifiers, a superflat correction to remove the effects of vignetting by the spherical abberation corrector in `slot-mode' observations, and aperture photometry using."
+. Relative photometry was obtained for. all runs. by inclucling a reference star in the field of view., Relative photometry was obtained for all runs by including a reference star in the field of view.
+ This cllectively removes the effects of the moving pupil and varying aperture of SALT during the course of a photometric run., This effectively removes the effects of the moving pupil and varying aperture of SALT during the course of a photometric run.
+ The light curves for the four runs are shown in order of 1 and in compacted form in Fig. 2.., The light curves for the four runs are shown in order of $T$ and in compacted form in Fig. \ref{lightcurves}.
+ In the 9 December 2005 run the large gap is caused by an interruption of observation. during which VW Livi decreased in brightness by ~ 0.4 mag.," In the 9 December 2005 run the large gap is caused by an interruption of observation, during which VW Hyi decreased in brightness by $\sim$ 0.4 mag."
+ There is à noticeable change in the amplitude of rapid. variations from run to run., There is a noticeable change in the amplitude of rapid variations from run to run.
+ This is almost entirely due to the dillering prominence of DNOs., This is almost entirely due to the differing prominence of DNOs.
+ This lighteurye. at the end of a superoutburst. is irregular and no DNOs are directly. visible in it but the detailed analysis given below shows that DNOs at very low amplitude are actually occasionally. present.," This lightcurve, at the end of a superoutburst, is irregular and no DNOs are directly visible in it – but the detailed analysis given below shows that DNOs at very low amplitude are actually occasionally present."
+ Short trains of QPOs are sometimes visible., Short trains of QPOs are sometimes visible.
+ From the AAVSO archive there are only two magnitude estimates available for the final descending part of this superoutburst. which makes estimation of T uncertain.," From the AAVSO archive there are only two magnitude estimates available for the final descending part of this superoutburst, which makes estimation of $T$ uncertain."
+ However. there are independent. indicators that lead to the values quoted in Table 1..," However, there are independent indicators that lead to the values quoted in Table \ref{dno7tab1}."
+ Phe change of mean magnitude between this and the next night shows that VW Livi was on a rapidly descending part of the outburs iehteurve., The change of mean magnitude between this and the next night shows that VW Hyi was on a rapidly descending part of the outburst lightcurve.
+" Phe DNOs that appear in the second half of 16 5 January run have a period ~ 22 s. which is at the vinning of the steep climb in period seen in figure |. of ""aper IV. and shows that we have for the first time caugh 10 transition from rare to frequent DNOs at the start of 16 rapid. deceleration."," The DNOs that appear in the second half of the 5 January run have a period $\sim$ 22 s, which is at the beginning of the steep climb in period seen in figure 1 of Paper IV, and shows that we have for the first time caught the transition from rare to frequent DNOs at the start of the rapid deceleration."
+ In fact. we have not seen. DNOs a is stage of a superoutburst before (c.g. run 81594 in table lof Paper IV).," In fact, we have not seen DNOs at this stage of a superoutburst before (e.g. run S1594 in table 1 of Paper IV)."
+ These properties indicate 7 ~ 0.0, These properties indicate $T$ $\sim$ 0.0.
+ Towards the end of the run there are 24.8 s DNOs arc 400 $ QPOs. which have a ratio of 16.1. close to the expectec relationship.," Towards the end of the run there are 24.8 s DNOs and 400 s QPOs, which have a ratio of 16.1, close to the expected relationship."
+ In the rapid variation regime the immediately obvious features of the light curve. (Fig. 2)), In the rapid variation regime the immediately obvious features of the light curve (Fig. \ref{lightcurves}) )
+ are the presence of persistent large amplitude (~ 0.1 mag) DNOs and. several Cycles of similar amplitude QPOs., are the presence of persistent large amplitude $\sim$ 0.1 mag) DNOs and several cycles of similar amplitude QPOs.
+ A Fourier transform (ET) of the entire longer section of the run provides an average, A Fourier transform (FT) of the entire longer section of the run provides an average
+nmücrovariabilitv.,microvariability.
+ Of. these. 47 are classified as Sevfert galaxies. 64 as QSOs. and 6 as BALQSOs.," Of these, 47 are classified as Seyfert galaxies, 64 as QSOs, and 6 as BALQSOs."
+ In their entire sample 21.4 per cent of the objects were found to exhibit micerovariabilitv. but among objects classified as Sevfert galaxies. OSOs and BALOSOs. microvariability was seen in 17 per cent. 23 per cent and 50 per cent. respectively (Carini et 22007).," In their entire sample 21.4 per cent of the objects were found to exhibit microvariability, but among objects classified as Seyfert galaxies, QSOs and BALQSOs, microvariability was seen in 17 per cent, 23 per cent and 50 per cent, respectively (Carini et 2007)."
+ In actelition. Rabbette et (1998) have noted that two radio-quiet. BALQSOs cdisplaved. shor term X-ray variabilitv.," In addition, Rabbette et (1998) have noted that two radio-quiet BALQSOs displayed short term X-ray variability."
+ Phe observed high fraction of microvariations in BALOSOs suggests that it might. be worthwhile to expend more of the observing time devotec to microvariability on the I2ALOSO class if one wants to uncerstand physical processes in or near the accretion disc., The observed high fraction of microvariations in BALQSOs suggests that it might be worthwhile to expend more of the observing time devoted to microvariability on the BALQSO class if one wants to understand physical processes in or near the accretion disc.
+ Clearly. the present sample size of BALQSOs is very smal compared to those of the non-BALQSO classes. and. no useful conclusions about their nature can be drawn from them.," Clearly, the present sample size of BALQSOs is very small compared to those of the non-BALQSO classes, and no useful conclusions about their nature can be drawn from them."
+ Therefore it is important to increase the sample of BALQSOs. so as to be able to arrive at firmer conclusions about the fraction showing microvariabilitv.," Therefore it is important to increase the sample of BALQSOs, so as to be able to arrive at firmer conclusions about the fraction showing microvariability."
+ We note tha if BALQSOs do really show a substantially higher. duty evele for microvariability than do non-BAL RQQSOs. this would. shed light on the question of whether or not. raclio-quict BALOSOs are special cases of the RQQSOs. especially in terms of their microvariabilitv properties.," We note that if BALQSOs do really show a substantially higher duty cycle for microvariability than do non-BAL RQQSOs, this would shed light on the question of whether or not radio-quiet BALQSOs are special cases of the RQQSOs, especially in terms of their microvariability properties."
+ For instance. if even weak jets dominate the rapid. variabilitv (e.g... CGopal-Ixrishna ct 22003) then a higher duty. evele [or microvariabilitv will give indirect support for the hypothesis that BALQSOs are viewed at angles nearly perpendicular to their aceretion dises (e.g. Ghosh Punsly 2007).," For instance, if even weak jets dominate the rapid variability (e.g., Gopal-Krishna et 2003) then a higher duty cycle for microvariability will give indirect support for the hypothesis that BALQSOs are viewed at angles nearly perpendicular to their accretion discs (e.g., Ghosh Punsly 2007)."
+ This is because jet [luctuations originating in relativistic jets pointing close to our line-of-sight are amplified in magnituce and compressed. in timescale (e.g... Gopal-Ixrishna. et 22003).," This is because jet fluctuations originating in relativistic jets pointing close to our line-of-sight are amplified in magnitude and compressed in timescale (e.g., Gopal-Krishna et 2003)."
+ Whereas. if BALOSOs show only the usual low microvariability. duty eveles of normal RQQSOs (around 20 per cent) and the fluctuations still arise in weak jets. that would provide indirect support for alternative mocels. such as those where the BAL outllows come out closer to the dise plane (e.g.. Elvis 2000).," Whereas, if BALQSOs show only the usual low microvariability duty cycles of normal RQQSOs (around 20 per cent) and the fluctuations still arise in weak jets, that would provide indirect support for alternative models, such as those where the BAL outflows come out closer to the disc plane (e.g., Elvis 2000)."
+ In conjunction with X-ray and optical spectral properties. such variation information is very useful in constraining various physical mocels for the origin of microvariability (c.e.. Czerny et 22008) and the nature of BALQSOs itself (e.g.. Wevmann et 11991: Elvis et al.," In conjunction with X-ray and optical spectral properties, such variation information is very useful in constraining various physical models for the origin of microvariability (e.g., Czerny et 2008) and the nature of BALQSOs itself (e.g., Weymann et 1991; Elvis et al."
+ 2000)., 2000).
+ To address these (questions. we have recently started. a pilot program to make an extensive search for optical microvariabilitv of BALOQSOs.," To address these questions, we have recently started a pilot program to make an extensive search for optical microvariability of BALQSOs."
+ This paper is organized as follows., This paper is organized as follows.
+ In. Section 2 we describe the main aspects of our sample selection. criteria. while Section 3 brielly describes our observations ancl the data reductions.," In Section 2 we describe the main aspects of our sample selection criteria, while Section 3 briefly describes our observations and the data reductions."
+ In Sections 4 and 5 we present our analysis and. results. respectively.," In Sections 4 and 5 we present our analysis and results, respectively."
+ Section 6 gives a discussion. and our conclusions., Section 6 gives a discussion and our conclusions.
+ Our sample is chosen from the DALOSO catalogues compiled by Trump et ((2006). Scaringi et ((2009) and Gibson et ((2009) which are based on Sloan Digital Sky Survey (SDSS) Data Releases 3 and 5 (DRS: Schneider et 22005: DR5: Adclman-AIcCarthy et 22007: Schneider et 22007).," Our sample is chosen from the BALQSO catalogues compiled by Trump et (2006), Scaringi et (2009) and Gibson et (2009) which are based on Sloan Digital Sky Survey (SDSS) Data Releases 3 and 5 (DR3: Schneider et 2005; DR5: Adelman-McCarthy et 2007; Schneider et 2007)."
+ In addition. we also included. one brighter BALQSO from the compilation bv Wevmann et ((1991).," In addition, we also included one brighter BALQSO from the compilation by Weymann et (1991)."
+ Most. of the sources were selected in such a way that both optical ancl X-ray spectral data are available for. them in archives., Most of the sources were selected in such a way that both optical and X-ray spectral data are available for them in archives.
+ All were at declinations that allowed. for. the observations to be made at relatively low air masses., All were at declinations that allowed for the observations to be made at relatively low air masses.
+ We also required. our candidate sources to have ο;17., We also required our candidate sources to have $_{i} \leq 17$.
+ This constraint means that even with a I1-m class telescope we could obtain a good enough signal to noise ratio to detect μισάτοις of «0.02 mae with a reasonably good time resolution of <10 minutes., This constraint means that even with a 1-m class telescope we could obtain a good enough signal to noise ratio to detect fluctuations of $<0.02$ mag with a reasonably good time resolution of $< 10$ minutes.
+ We also limit the. BALOQSOs to have absolute magnitudes M;</—24.5. so that the Dux contribution from the host galaxy can be assumed to be negligible (Miller et 11990).," We also limit the BALQSOs to have absolute magnitudes $_{i} < -24.5$, so that the flux contribution from the host galaxy can be assumed to be negligible (Miller et 1990)."
+ Our final sample consists of a total of 19 BALQSOs. as listed. in Table 1..," Our final sample consists of a total of 19 BALQSOs, as listed in Table \ref{tab:sample}."
+ Xmong these BALQSOs 8 are classified as Hil3ALs. 10 are LoDB.ALs and 1 is a mini BAL.," Among these BALQSOs 8 are classified as HiBALs, 10 are LoBALs and 1 is a mini BAL."
+" Phe whole sample covers a redshift range of 0.39<so),«2.9.", The whole sample covers a redshift range of $ 0.39< z_{em} < 2.9$.
+ Our observations of each of the BALQSOs were carried out continuously for  4h in the Ro passband. mainly using the 1.04-m Sampurnanancl telescope located. at. the Arvabhatta Research Institute. of observational sciencEs (ARLES). Nainital. bndia.," Our observations of each of the BALQSOs were carried out continuously for $\sim$ 4h in the R passband, mainly using the 1.04-m Sampurnanand telescope located at the Aryabhatta Research Institute of observational sciencES (ARIES), Nainital, India."
+ I0 has Ritchey-Chretien (RC) optics with a £/13 beam and is equipped with a ervogenicallv cooled CCD detector with a 2048 pixel 2048 pixel chip mounted at the Cassegrain focus (Sagar 1909)., It has Ritchey-Chretien (RC) optics with a $/$ 13 beam and is equipped with a cryogenically cooled CCD detector with a 2048 pixel $\times$ 2048 pixel chip mounted at the Cassegrain focus (Sagar 1999).
+ Phe readout noise of the CCD chip is 5.3 /pixel and it has a gain of 10ο Analog to Digital Unit (ADU)., The readout noise of the CCD chip is 5.3 $^{-}$ /pixel and it has a gain of 10 $^{-}$$/$ Analog to Digital Unit (ADU).
+ Each pixel of the CCD chip has a dimension of 24 sam>. corresponding to 0.37 arcsec? on the sky. and so covers a total Ποια of ~ 13). 137.," Each pixel of the CCD chip has a dimension of 24 $\mu$ $^{2}$, corresponding to 0.37 $^{2}$ on the sky, and so covers a total field of $\sim$ $^{\prime}$ $\times$ $^{\prime}$."
+ Vo improve the signal to noise ratio. observations were carried out in a 2 pixel 2 pixel binning mode.," To improve the signal to noise ratio, observations were carried out in a 2 pixel $\times$ 2 pixel binning mode."
+" The tvpical seeing during our observing runs at ARLES was ~ 3"".", The typical seeing during our observing runs at ARIES was $\sim$ $^{\prime\prime}$.
+ In addition. two sources were observed. with 2.01-m Himalayan Chandra Telescope (LICL) located at the Indian Astronomical Observatory (LAO). LHanle. India.," In addition, two sources were observed with 2.01-m Himalayan Chandra Telescope (HCT) located at the Indian Astronomical Observatory (IAO), Hanle, India."
+ It is also of the RC design with a [/9 beam at the Casseerain cust., It is also of the RC design with a $/$ 9 beam at the Cassegrain .
+. The detector was a cevogenically cooled 2048 . 4096 chip. of which the central 2048 2048 pixels were used.," The detector was a cryogenically cooled 2048 $\times$ 4096 chip, of which the central 2048 $\times$ 2048 pixels were used."
+ The pixel size is 15 fam? so that the image scale of 0.29 aresec/pixel. covers an area of about. 10∕ 10 on the sk., The pixel size is 15 $\mu$ $^{2}$ so that the image scale of 0.29 $/$ pixel covers an area of about $^{\prime}$ $\times$ ${^\prime}$ on the sky.
+ The readout noise of this CCD is 4.87 /pixel and the gain is L226 /XDU., The readout noise of this CCD is 4.87 $^{-}$ /pixel and the gain is 1.22 $^{-}$$/$ ADU.
+ The CCD was used in an unbinned modo., The CCD was used in an unbinned mode.
+" The tvpical seeing during our observations at LAO was ον 1.5"".", The typical seeing during our observations at IAO was $\sim$ $^{\prime\prime}$.
+ We chose an . filter for this observational program because it is at the maximum response of the CCD system: thus the time resolution achievable for cach object is maximized., We chose an R filter for this observational program because it is at the maximum response of the CCD system; thus the time resolution achievable for each object is maximized.
+ As most of our sources have g;1617. the best time resolution we could achieve was of the order of 3 minutes. and we almost always nianaged data points spaced less than S minutes apart. so very rapid Iuctuations could be picked up.," As most of our sources have $g_{i}\sim 16-17$, the best time resolution we could achieve was of the order of 3 minutes, and we almost always managed data points spaced less than 8 minutes apart, so very rapid fluctuations could be picked up."
+ We also took care to select sources and fields of view so às to ensure availability of at least two. but usually more. comparison stars on the CCD frame that werewithin around | mag of the QSO's brightness.," We also took care to select sources and fields of view so as to ensure availability of at least two, but usually more, comparison stars on the CCD frame that werewithin around 1 mag of the QSO's brightness."
+ This allowed us to, This allowed us to
+"decades later. are curiously negative. decreasing between the phases 0.719 aud 0,825 from -6.5 to -16.6 kins +.","decades later, are curiously negative, decreasing between the phases 0.749 and 0.825 from -6.5 to -16.6 km $^{-1}$."
+ Tn this wav appears to be more redshifted with respect to the intercombination lines thanCiv., In this way appears to be more redshifted with respect to the intercombination lines than.
+ The redshift relative to the intercombination lues. deduced from tle UST/STIS spectra. decreases with the ionization potential of the corresponding ions. being mach less for Ov] than for ΟΠΗ aud iv].," The redshift relative to the intercombination lines, deduced from the HST/STIS spectra, decreases with the ionization potential of the corresponding ions, being much less for ] than for ] and ]."
+ The 1610 Hine appears to be for all but the first observation strongly redshifted with respect to the systemic radial velocity of 12.1 or 12.9 1. but its radial velocity way be compatible with that of the compact compoucut at the," The 1640 line appears to be for all but the first observation strongly redshifted with respect to the systemic radial velocity of 12.4 or 12.9 $^{-1}$, but its radial velocity may be compatible with that of the compact component at the"
+illustratecl case. the jet never. catches up to the winds boundary since this jet is launched 50 Myr after the wind and is. moreover. cnereised for only LO Myr.,"illustrated case the jet never catches up to the wind's boundary since this jet is launched 50 Myr after the wind and is, moreover, energised for only 10 Myr."
+ In the other three cases. the jets do catch up to the bubble boundary. despite having a shorter lifetime (in one case) ancl lower total energies (in two cases).," In the other three cases, the jets do catch up to the bubble boundary, despite having a shorter lifetime (in one case) and lower total energies (in two cases)."
+ Once the jets pass through the bubble ancl enter the IGM their velocities of advance are seen to change., Once the jets pass through the bubble and enter the IGM their velocities of advance are seen to change.
+ With our assumptions that the bubble continues to inflate even after the ACGN. wind is shut. oll. while the jet-fed cocoons stop growing right after the jet is switched olf we see that in all three of these displaved cases the bubble boundary could eventually once again overtake the jet: however. thanks to the confinement exerted by the cocoon on the bubble in the interim. the times when this occurs that can be read from 11 are unclerestinmates.," With our assumptions that the bubble continues to inflate even after the AGN wind is shut off, while the jet-fed cocoons stop growing right after the jet is switched off, we see that in all three of these displayed cases the bubble boundary could eventually once again overtake the jet; however, thanks to the confinement exerted by the cocoon on the bubble in the interim, the times when this occurs that can be read from 1 are underestimates."
+" We can now define the catch-up time. /,. and catch-up radius. A2.. through the relation Ay.)ο=Re."," We can now define the catch-up time, $t_c$, and catch-up radius, $R_c$, through the relation $R_j(t_c) = R_S(t_c) \equiv R_c$."
+ The formula in ((10) is valid for /;</f. but at later times the jets are ploughing through the constant density IGM. and we have for f«£f;|τ In 22 values for A are plotted against /;. the time interval between the winds onset and the jets? launching. for substantial ranges of the most important input. parameters.," The formula in (10) is valid for $t_j < t < t_c$, but at later times the jets are ploughing through the constant density IGM, and we have for $t_c < t < t_j+\tau_j$ In 2 values for $R_c$ are plotted against $t_j$, the time interval between the wind's onset and the jets' launching, for substantial ranges of the most important input parameters."
+ Each of the panels shows the results for two or three cillerent jet powers., Each of the panels shows the results for two or three different jet powers.
+" The values of A (imposed by the ICM: 22) are shown in 22a and 2b for Ly.=1.5810 org and by=1.582107 ere. respectively. while the «corresponding values of Rp for 4,=3.16«107 erg 220) and ££,=3.16«10"" ere 22d) exceed 2 Alpe and are not shown."," The values of $R_P$ (imposed by the IGM; 2) are shown in 2a and 2b for $E_w = 1.58 \times 10^{57}$ erg and $E_w = 1.58 \times 10^{58}$ erg, respectively, while the corresponding values of $R_P$ for $E_w = 3.16 \times 10^{58}$ erg 2c) and $E_w = 3.16 \times 10^{60}$ erg 2d) exceed 2 Mpc and are not shown."
+" The convention used in 22 is that if the jet never catches up to the bubble's boundary then 2, is plotted as 2),", The convention used in 2 is that if the jet never catches up to the bubble's boundary then $R_c$ is plotted as $R_p$.
+ Under most circumstances we see that £2). rises as fj increases. às expected. since a jet should. take longer to catch up to a bubble which has had a greater. heacl-start.," Under most circumstances we see that $R_c$ rises as $t_j$ increases, as expected, since a jet should take longer to catch up to a bubble which has had a greater head-start."
+ In all cases. more powerful jets obviously produce. lower values of £. for fixed. values of /;.," In all cases, more powerful jets obviously produce lower values of $R_c$ for fixed values of $t_j$ ."
+" The less powerful jets (og. L;/L,=0.2 in 22a. 2b and 2d: Lj/L,,=1.0 in 22b and 2) and even very powerful. but. short-lived. jets (LjfL,—3.0 in 22d) may never catch up to the wind boundary if the jets launch is sullicicntby delaved."," The less powerful jets (e.g., $L_j/L_w = 0.2$ in 2a, 2b and 2d; $L_j/L_w = 1.0$ in 2b and 2d), and even very powerful, but short-lived, jets $L_j/L_w = 3.0$ in 2d) may never catch up to the wind boundary if the jet's launch is sufficiently delayed."
+ The assumption that pi is fixed at the value it has attained ab { leads to the peculiar results plotted at. low values of f; in 22a. 2b. and 2d for some of the adopted: values of jet. power.," The assumption that $\rho_{b}$ is fixed at the value it has attained at $t_j$ leads to the peculiar results plotted at low values of $t_j$ in 2a, 2b, and 2d for some of the adopted values of jet power."
+ Because these nominal densities can be quite high for small values of /; (see SS) the jets appear to never cateh up to the wind. boundary: since. in actuality the bubble density would continue to decline. unless thejets were very short-Iived they would indeed eventually catch up to the winds boundary.," Because these nominal densities can be quite high for small values of $t_j$ (see 8) the jets appear to never catch up to the wind boundary; since, in actuality the bubble density would continue to decline, unless the jets were very short-lived they would indeed eventually catch up to the wind's boundary."
+ We shall now briefly examine some broad implications of the above analvtical scheme. taking a fiducial values of Evo10 erg For the kinetic οποιον of the wind outLlow and the same value for the kinetic energies of the relativistic jets.," We shall now briefly examine some broad implications of the above analytical scheme, taking a fiducial values of $E_w \sim 10^{59}$ erg for the kinetic energy of the wind outflow and the same value for the kinetic energies of the relativistic jets."
+" From ((1) and 22. it is seen that ifthe bubble's expansion is allowed to continue much after cessation of the wind outllow (747107 vr) then in the course of ~10° vears it would have grown to a quasi-spherical volume of size around 1 Alpe if it encounters an ambient LAL pressure «107 dene 7: this will be the case. provided. the AGN is located the field. rather than inside a dense eluster environment. and is not at too high a redshift so that the combination 9,,(1|z)«150."," From (1) and 2, it is seen that if the bubble's expansion is allowed to continue much after cessation of the wind outflow $\tau_{\rm act} \approx 10^8$ yr), then in the course of $\sim 10^9$ years it would have grown to a quasi-spherical volume of size around 1 Mpc if it encounters an ambient IGM pressure $< 10^{-16}$ dyne $^{-2}$; this will be the case provided the AGN is located the field, rather than inside a dense cluster environment, and is not at too high a redshift so that the combination $\delta_m (1+z)^3 < 150$."
+ However. if a relativistic jet is subsequently LlatuacheL. then once it crosses the bubble’s boundary. the much higher pressure of its radio lobe would begin to act as a wall. blocking.e and even reversinge the bubble’s expansiona alonge the jets direction.," However, if a relativistic jet is subsequently launched, then once it crosses the bubble's boundary, the much higher pressure of its radio lobe would begin to act as a wall, blocking, and even reversing the bubble's expansion along the jet's direction."
+ The cartoon in 332 shows this as the last phase of this dual mode ACN energy output., The cartoon in 3 shows this as the last phase of this dual mode AGN energy output.
+" The pressure in the radio lobe is indeed: much higher than the LGAL pressure. since it is given by LonecS10H7, dvne em7. assuming the minimum. energy condition which signifies an energy equipartition between magnetic ancl relativistic particles: here Boy is the lobe ficld in microgauss and values of B5 between 1 ancl 30 are typically estimated for FR LE ICs."," The pressure in the radio lobe is indeed much higher than the IGM pressure, since it is given by $P_{\rm lobe} \simeq 8 \times 10^{-14} B^2_{-6}$ dyne $^{-2}$, assuming the minimum energy condition which signifies an energy equipartition between magnetic and relativistic particles; here $B_{-6}$ is the lobe field in microgauss and values of $B_{-6}$ between 1 and 30 are typically estimated for FR II RGs."
+ Inserting this value of Lane into ((2) vields an equilibrium. radius of the bubble. {ές20.22Mpebt:ABvU. which is in the range from a few to hundreds of kiloparsecs.," Inserting this value of $P_{\rm lobe}$ into (2) yields an equilibrium radius of the bubble, $R_{eq} \simeq 
+0.22~{\rm Mpc}~ E^{1/3}_{w,59} B^{-2/3}_{-6}$, which is in the range from a few to hundreds of kiloparsecs."
+ These values are also displaved in H1 and 2 as grev bands corresponding to the range of values of B mentioned above., These values are also displayed in 1 and 2 as grey bands corresponding to the range of values of $B$ mentioned above.
+ Note that the higher pressure in the radio lobes may even reverse the expansion of the bubble boundary in the areas of contact with the lobes iJ?cI4.," Note that the higher pressure in the radio lobes may even reverse the expansion of the bubble boundary in the areas of contact with the lobes if $R_c >
+R_{eq}$."
+" Phus. the actual width of the gap between the radio lobes should be between 2h, and 2/4. assuming that the jets live long enough and. are launched soon enough to overtake the bubble."," Thus, the actual width of the gap between the radio lobes should be between $2R_{eq}$ and $2R_c$, assuming that the jets live long enough and are launched soon enough to overtake the bubble."
+ The consisteney of these estimates with the observed. widths of the radio gaps in SDs argues for the ellicacy of the wind's confinement by the lobes. as envisioned. in our model.," The consistency of these estimates with the observed widths of the radio gaps in SDs argues for the efficacy of the wind's confinement by the lobes, as envisioned in our model."
+" We can estimate the time scale over which this recompression along the jet axis from £2.to F4, could occur as", We can estimate the time scale over which this recompression along the jet axis from $R_c$to $R_{eq}$ could occur as
+"(C,=27!m«107ns","$C_\ell = 2.7^{+1.1}_{-2.6} \times10^{-5} \,\mu {\rm K}^2$."
+ IIowever. these numbers should be compared with caution as the two experiments lave different flux cuts for source masking. aud the excess power will depend ou the fix to which sources have been maskecl.," However, these numbers should be compared with caution as the two experiments have different flux cuts for source masking, and the excess power will depend on the flux to which sources have been masked."
+ We also investigate the effects of allowing the auplitude of the SZE power spectrum to float freely., We also investigate the effects of allowing the amplitude of the SZE power spectrum to float freely.
+ The SZE power spectrum template is based on the simulations in ?.., The SZE power spectrum template is based on the simulations in \cite{shaw2009}.
+" The simulations are for a WALAPS cosmolosv with a,=0.77.", The simulations are for a WMAP5 cosmology with $\sigma_8 = 0.77$.
+ The amplitude of the SZE power spectrum is expected to scale approximately as og. so the derived SZE amplitudecan be related to os.," The amplitude of the SZE power spectrum is expected to scale approximately as $\sigma_8^7$, so the derived SZE amplitudecan be related to $\sigma_8$."
+ Iu practice. the APENX-SZ data set lacks the sensitivity to make a detection of SZE power: however. the results cau be used to place an upper limit on ox.," In practice, the APEX-SZ data set lacks the sensitivity to make a detection of SZE power; however, the results can be used to place an upper limit on $\sigma_8$."
+ The upper liits on σς and the poiut source amplitudes for the joiut fit are reported in Table 2.., The upper limits on $\sigma_8$ and the point source amplitudes for the joint fit are reported in Table \ref{tab:params}.
+ We assume a flat prior on os., We assume a flat prior on $\sigma_8$.
+ The exact amplitude of the SZE spectrum is only poorly nuderstood. leading to a systematic uncertaintv on σς (2). ," The exact amplitude of the SZE spectrum is only poorly understood, leading to a systematic uncertainty on $\sigma_8$ \citep{komatsu2002}. ."
+This systematic uncertainty is not included iu, This systematic uncertainty is not included in
+radius fixed at Ai. yields shorter timescales and lower accretion luminosities.,radius fixed at $R_{\rm dis}$ yields shorter timescales and lower accretion luminosities.
+" For the black hole cases. I choose /244,=LOCAL/c. so just outside the innermost stable cireular orbit of a non-rotating (Schwarzsehild-) black hole at Bisco=6674/7."," For the black hole cases, I choose $R_{\rm dis}= 10 GM/c^2$, so just outside the innermost stable circular orbit of a non-rotating (Schwarzschild-) black hole at $R_{\rm ISCO}= 6GM/c^2$."
+" The details of the fallback times and energies change slightly with /?4;,. but none of the conclusions of this paper depends on the exact numerical value of /?,4.."," The details of the fallback times and energies change slightly with $R_{\rm dis}$, but none of the conclusions of this paper depends on the exact numerical value of $R_{\rm dis}$."
+ As an illustration. I show in Fig.," As an illustration, I show in Fig."
+ 2. a set of fallback trajectories., \ref{fig2} a set of fallback trajectories.
+ The initial conditions were taken from the double neutron star merger calculation with I.1 and 1.6 that is shown in Fig. 2.., The initial conditions were taken from the double neutron star merger calculation with 1.1 and 1.6 that is shown in Fig. \ref{fig2}.
+ 200 randomly chosen trajectories out of the more than 5800 fallback particles are plotted., 200 randomly chosen trajectories out of the more than 5800 fallback particles are plotted.
+ There is a broad distribution of eccentricities and fallback times. while some particles return from their initial position directly towards the central object. others follow highly eccentricity orbits out to distances of many 10 000 km before they come to a halt and fall back towards the centre.," There is a broad distribution of eccentricities and fallback times, while some particles return from their initial position directly towards the central object, others follow highly eccentricity orbits out to distances of many 10 000 km before they come to a halt and fall back towards the centre."
+ Fig., Fig.
+ 3 shows the accretion luminosities. Lice=dIdl. derived for the various DNS and NSBH systems.," \ref{fig3} shows the accretion luminosities, $L_{\rm acc}= dE_{\rm fb}/dt$, derived for the various DNS and NSBH systems."
+" Here. fy, /denotes the difference between the potential plus kinetic energ at the start radius. ης and the potential energ' at the dissipation radius. £71..."," Here, $E_{\rm fb}$ denotes the difference between the potential plus kinetic energy at the start radius, $r_i$, and the potential energy at the dissipation radius, $R_{\rm dis}$."
+" The curves have been obtained by binning the energies contained in the fallback material. μι. according to the corresponding fallback times. Τι, see eq. (2)."," The curves have been obtained by binning the energies contained in the fallback material, $E_{\rm fb}$ , according to the corresponding fallback times, $\tau_{i}$, see eq. \ref{fallback_time}) )."
+ A fraction ος of this energy is channelled into X-rays. Ly=exLace. The double neutron star cases form a rather homogeneous class with respect to their fallback accretion. in all cases the fallback material is approximately 0.03.. see Tab. |..," A fraction $\epsilon_{\rm X}$ of this energy is channelled into X-rays, $L_X = \epsilon_{\rm X} L_{\rm acc}$ The double neutron star cases form a rather homogeneous class with respect to their fallback accretion, in all cases the fallback material is approximately 0.03, see Tab. \ref{tab:runs}. ."
+ After an initial. short-lived plateau. the luminosity falls off with time close to the expected 5/3-power law (Rees1988:Phinney1989).," After an initial, short-lived plateau, the luminosity falls off with time close to the expected 5/3-power law \citep{rees88,phinney89}."
+ It has to be pointed out that the last point in these curves is determined by the numerical mass resolution in the hydrodynamies simulations and should therefore be interpreted with some caution., It has to be pointed out that the last point in these curves is determined by the numerical mass resolution in the hydrodynamics simulations and should therefore be interpreted with some caution.
+ All other points should be a fair representation of the overall fallback activity., All other points should be a fair representation of the overall fallback activity.
+ Typically. the X-ray luminosity about one hour after the coalescence is Ly~(5X):tort erg/s. For the investigated mass range. the spread in the luminosities one hour after the coalescence is about one order of The neutron star black hole cases show a larger diversity.," Typically, the X-ray luminosity about one hour after the coalescence is $L_X \sim \left(\frac{\epsilon_{\rm X}}{0.1}\right) \cdot 10^{44}$ erg/s. For the investigated mass range, the spread in the luminosities one hour after the coalescence is about one order of The neutron star black hole cases show a larger diversity."
+ The mass in the fallback material of different mass ratios varies by about a factor of 500. see Tab. 1..," The mass in the fallback material of different mass ratios varies by about a factor of 500, see Tab. \ref{tab:runs},"
+ an hour after the merger the accretion luminosities of the different NSBH systems differ by about two ordersof magnitude., an hour after the merger the accretion luminosities of the different NSBH systems differ by about two ordersof magnitude.
+ Also the involved time scales change strongly, Also the involved time scales change strongly
+are mostly weak with small Mach number of AZ<3. because thev form in the hot gas of AT= keV. In this paper. we have studied DSA at weak cosmological shocks.,"are mostly weak with small Mach number of $M\la 3$, because they form in the hot gas of $kT \ga$ keV. In this paper, we have studied DSA at weak cosmological shocks."
+ Since the test-particle solutions could provide a simple vet reasonable description for weak shocks. we first suggested analvtie solutions which describe the DSA in the test-particle regime. including both thepre-existing aud injected CI. populations.," Since the test-particle solutions could provide a simple yet reasonable description for weak shocks, we first suggested analytic solutions which describe the DSA in the test-particle regime, including both the and injected CR populations."
+ We adopted a thermal leakage injection model to emulate the acceleration of suprathermal particles into the CR population. along with a simple transport model in whieh Alfvénn waves self-excited by the CHR. streaming instability drift relative to the bulk plasma upstream of the gas subshock.," We adopted a thermal leakage injection model to emulate the acceleration of suprathermal particles into the CR population, along with a simple transport model in which Alfvénn waves self-excited by the CR streaming instability drift relative to the bulk plasma upstream of the gas subshock."
+ We then performed kinetic DSA simulations and compared the analvtie and numerical solutions for wide ranges of model parameters relevant for shocks in ICAIs and cluster outskirts: (he shock Mach number Af=1.5—5. the slope of the pre-existing CR spectrum 5=41—5. the ratio of the upstreamCR to gas pressure 2=0.01—0.1. the injection parameter ej=0.25—B 0.3. and the Allvénnie speed parameter ὁ=0—0.12.," We then performed kinetic DSA simulations and compared the analytic and numerical solutions for wide ranges of model parameters relevant for shocks in ICMs and cluster outskirts: the shock Mach number $M = 1.5 - 5$, the slope of the pre-existing CR spectrum $s=4-5$, the ratio of the upstreamCR to gas pressure $R=0.01-0.1$, the injection parameter $\epsilon_B=0.25-0.3$ , and the Alfvénnic speed parameter $\delta=0-0.42$."
+ The upstream gas was assumed (o be fully ionized with 7i101 Kk. The main results can be sunmiarized as follows: 1) For weak shocks with M.3. the test-particle solutions given in Equation (13)) should provide a good approximation lor the time-dependent CR spectrum at the shock location.," The upstream gas was assumed to be fully ionized with $T_0=10^7$ K. The main results can be summarized as follows: 1) For weak shocks with $M \la 3$, the test-particle solutions given in Equation \ref{ftest}) ) should provide a good approximation for the time-dependent CR spectrum at the shock location."
+ We note (hat the test-particle slope. q. in Equation (2)) and the maximum momenta. Paasl/). in Equation (3)) max include the Alfvénnie drift effect.," We note that the test-particle slope, $q$ , in Equation \ref{qtp}) ) and the maximum momentum, $p_{\rm max}(t)$, in Equation \ref{pmax}) ) may include the Alfvénnic drift effect."
+ 2) Forthe injection parameter considered here. ej;=0.25—0.3. the injection Iraction is rather low. typically €~5x10? to 10/7 for M<3.," 2) Forthe injection parameter considered here, $\epsilon_B=0.25-0.3$, the injection fraction is rather low, typically $\xi \sim 5\times10^{-5}$ to $10^{-3}$ for $M \la 3$."
+ The pre-existing CR population provides more particles for DSA than the freshly injected population., The pre-existing CR population provides more particles for DSA than the freshly injected population.
+ Hence. the pre-existing population dominates over the injected population.," Hence, the pre-existing population dominates over the injected population."
+ If there exist no CRs upstream (J?= 0). the downstream: CR pressure. absorbs typically much less than ~1 of the shock ram pressure for M.X3.," If there exist no CRs upstream $R=0$ ), the downstream CR pressure absorbs typically much less than $\sim 1$ of the shock ram pressure for $M \la 3$."
+ With pre-existing CRs that accounts for 5% of the gas thermal pressure in (he upstream flow. the CR acceleration efficiency increases to a few to 10 for those weak shocks.," With pre-existing CRs that accounts for 5 of the gas thermal pressure in the upstream flow, the CR acceleration efficiency increases to a few to 10 for those weak shocks."
+ 3) For the pre-exisiting population. (he enhancement of the distribution function across the shock. fs(p)/fip). at a given momentum is substantially larger than that expected [rom the simple adiabatic compression.," 3) For the pre-exisiting population, the enhancement of the distribution function across the shock, $f_2(p)/f_1(p)$, at a given momentum is substantially larger than that expected from the simple adiabatic compression."
+ Hence. with amplified] magnetic fields downstream. the re-acceleration of pre-existing CR. electrons can result in a substantial svnchrotron radiation behind the shock.," Hence, with amplified magnetic fields downstream, the re-acceleration of pre-existing CR electrons can result in a substantial synchrotron radiation behind the shock."
+ We estimated that the enhancement in svnchrotron radiation across the shock. ο)Ja). is about a few to several for M.~ 1.5. while itcould reach to LO?—10? for AZ~ 3. depending on (he detail model parameters.," We estimated that the enhancement in synchrotron radiation across the shock, $J_2({\nu})/J_0({\nu})$, is about a few to several for $M \sim 1.5$ , while itcould reach to $10^2-10^3$ for $M\sim 3$ , depending on the detail model parameters."
+ This is substantially larger than the, This is substantially larger than the
+at different epochs will give us important clues on the evolution of galaxies.,at different epochs will give us important clues on the evolution of galaxies.
+ The first quantitative SFRs were derived from evolutionary synthesis models of galaxy colors (Tinsley 1968. 1972. Searle et al.," The first quantitative SFRs were derived from evolutionary synthesis models of galaxy colors (Tinsley 1968, 1972, Searle et al."
+ 1973). confirming the trends in SFRs and star formation histories along the Hubble sequence. and giving the first predictions of the evolution of the SFR with cosmic lookback time.," 1973), confirming the trends in SFRs and star formation histories along the Hubble sequence, and giving the first predictions of the evolution of the SFR with cosmic lookback time."
+ The development of more precise direct SFR diagnostics includes the integrated emission-line fluxes (Cohen 1976. Kennicutt 1983). near-ultraviolet continuum fluxes (Donas & Deharveng 1984). and infrared continuum fluxes (Harper & Low 1973. Rieke & Lebofsky 1978. Telesco & Harper 1980): see Kennicutt (1998) for a review.," The development of more precise direct SFR diagnostics includes the integrated emission–line fluxes (Cohen 1976, Kennicutt 1983), near-ultraviolet continuum fluxes (Donas $\&$ Deharveng 1984), and infrared continuum fluxes (Harper $\&$ Low 1973, Rieke $\&$ Lebofsky 1978, Telesco $\&$ Harper 1980); see Kennicutt (1998) for a review."
+ The hydrogen Balmer line i8 currently the most reliable tracer of star formation. since in regions and star-forming galaxies. the Balmer emission-line luminosity scales directly with the total ionizing flux of the embedded stars.," The hydrogen Balmer line is currently the most reliable tracer of star formation, since in regions and star-forming galaxies, the Balmer emission-line luminosity scales directly with the total ionizing flux of the embedded stars."
+ A widely known calibration of the line as SER tracer is the one devised by Kennicutt (1998)., A widely known calibration of the line as SFR tracer is the one devised by Kennicutt (1998).
+ However. it 15 important to take into account corrections for stellar absorption and reddening to obtain SFRs in agreement with the ones derived using other wavelengths (e.g. Rosa-Gonzállez et al.," However, it is important to take into account corrections for stellar absorption and reddening to obtain SFRs in agreement with the ones derived using other wavelengths (e.g. Rosa-Gonzállez et al."
+ 2002. Charlot et al.," 2002, Charlot et al."
+ 2002. Dopita et al.," 2002, Dopita et al."
+ 2002)., 2002).
+ In parallel. other diagnostics have been developed using the oxygen doublet [Ou] 13726. 3729 for the redshift range z0.4—1.5 (e.g. Gallagher et al.," In parallel, other diagnostics have been developed using the oxygen doublet ] $\lambda$ 3726, 3729 for the redshift range $z \sim 0.4-1.5$ (e.g. Gallagher et al."
+ 1989. Kennicutt 1998. Rosa-Gonzállez et al.," 1989, Kennicutt 1998, Rosa-Gonzállez et al."
+ 2002. Kewley et al.," 2002, Kewley et al."
+ 2004)., 2004).
+" Moreover. this ""Siagnostic is usefull when the line is not easily observable at higher redshifts (ς>O.4 in the optical)."," Moreover, this diagnostic is usefull when the line is not easily observable at higher redshifts $z \gtrsim 0.4$ in the optical)."
+" However. the On| doublet presents problems in reddening and abundance ""Sependence (Jansen. Franx & Fabricant 2001. Charlot et al."," However, the ] doublet presents problems in reddening and abundance dependence (Jansen, Franx $\&$ Fabricant 2001, Charlot et al."
+ 2002)., 2002).
+ Alternatively. tt is possible to estimate the SFR from the soft X-ray luminosity. which ts comparable to that determinec from the luminosity (Rosa Gonzállez et al.," Alternatively, it is possible to estimate the SFR from the soft X-ray luminosity, which is comparable to that determined from the luminosity (Rosa Gonzállez et al."
+ 2009. Rovilos et al.," 2009, Rovilos et al."
+ 2009)., 2009).
+ A strong dependence of the SFR and the stellar mass and its evolution with redshift has been found. with the bulk of star formation occurring first in massive galaxies. and later i1 less massive systems (e.g. Guzmánn et al.," A strong dependence of the SFR and the stellar mass and its evolution with redshift has been found, with the bulk of star formation occurring first in massive galaxies, and later in less massive systems (e.g. Guzmánn et al."
+ 1997. Brinchmant & Ellis 2000. Juneau et al.," 1997, Brinchmann $\&$ Ellis 2000, Juneau et al."
+ 2005. Bauer et al.," 2005, Bauer et al."
+ 2005. Bell et al.," 2005, Bell et al."
+ 2005. Pérrez-Gonzalez et al.," 2005, Pérrez-Gonzalez et al."
+ 2005. Feulner et al.," 2005, Feulner et al."
+ 2005. Papovich et al.," 2005, Papovich et al."
+ 2006. Caputi et al.," 2006, Caputi et al."
+ 2006. Reddy et al 2006. Erb et al.," 2006, Reddy et al 2006, Erb et al."
+ 2006. Noeske et al.," 2006, Noeske et al."
+ 2007a. Buat et al.," 2007a, Buat et al."
+ 2008)., 2008).
+ In the local universe. several studies have illustrated a relationship between the SFR and stellar mass. identifying two populations: galaxies on a star-forming sequence. and “quenched” galaxies. with little or no detectable star formation (Brinchmann et al.," In the local universe, several studies have illustrated a relationship between the SFR and stellar mass, identifying two populations: galaxies on a star-forming sequence, and $``$ quenched"" galaxies, with little or no detectable star formation (Brinchmann et al."
+ 2004. Salim et al.," 2004, Salim et al."
+ 2005. Lee 2006).," 2005, Lee 2006)."
+ At higher redshift. Noeske et al. (," At higher redshift, Noeske et al. ("
+"2007a) showed the existence of a ""main sequence” (MS) for SF galaxies in the SEFR-stellar mass relation over the redshift range 0.2«zI.I.","2007a) showed the existence of a $``$ main sequence"" (MS) for SF galaxies in the SFR–stellar mass relation over the redshift range $0.2 < z < 1.1$."
+ From the galaxies considered in this study. it was shown that the slope of the MS remains constant to zΣα. while the MS as a whole moves to higher SER às z increases.," From the galaxies considered in this study, it was shown that the slope of the MS remains constant to $z>1$, while the MS as a whole moves to higher SFR as $z$ increases."
+ Metallicity is another important property of galaxies. and its study is crucial for a deep understanding of galaxy formation and evolution. since it is related to the whole past history of the galaxy.," Metallicity is another important property of galaxies, and its study is crucial for a deep understanding of galaxy formation and evolution, since it is related to the whole past history of the galaxy."
+ Metallicity 1s a tracer of the fraction of baryonic mass already converted into stars and is sensitive to the metal losses due to stellar winds. supernovae and active nuclei feedbacks.," Metallicity is a tracer of the fraction of baryonic mass already converted into stars and is sensitive to the metal losses due to stellar winds, supernovae and active nuclei feedbacks."
+ A detailed description of the different metallicity methods and calibrations are given in Lara-Lóppez et al. (, A detailed description of the different metallicity methods and calibrations are given in Lara-Lóppez et al. (
+2009a.b).,"2009a,b)."
+ Stellar mass and metallicity are strongly correlated in SF galaxies. with massive galaxies showing higher metallicities than less massive galaxies.," Stellar mass and metallicity are strongly correlated in SF galaxies, with massive galaxies showing higher metallicities than less massive galaxies."
+ This relationship provides crucial insight into galaxy formation and evolution., This relationship provides crucial insight into galaxy formation and evolution.
+ The mass-metallicity (M.—Z) relation was first observed by Lequeux et al. (, The mass-metallicity $M-Z$ ) relation was first observed by Lequeux et al. (
+1979). has been intensively studied (Skillman et al.,"1979), has been intensively studied (Skillman et al."
+ 1989; Brodie & Huchra 1991: Zaritsky et al., 1989; Brodie $\&$ Huchra 1991; Zaritsky et al.
+ 1994: Richer & McCall 1995: Garnett et al., 1994; Richer $\&$ McCall 1995; Garnett et al.
+ 1997; Pilyugin & Ferrint 2000. among others). and it is well established by the work of Tremonti et al. (," 1997; Pilyugin $\&$ Ferrini 2000, among others), and it is well established by the work of Tremonti et al. ("
+2004. hereafter TO4) for the local universe (z ~ 0.1) using SDSS data.,"2004, hereafter T04) for the local universe (z $\sim$ 0.1) using SDSS data."
+ The study of the redshift evolution of the M—Z relation has provided us with erucial information on the cosmic evolution of star formation., The study of the redshift evolution of the $M-Z$ relation has provided us with crucial information on the cosmic evolution of star formation.
+ Regarding the evolution of the M—Z relation for SF galaxies at z « l. Savaglio et al. (," Regarding the evolution of the $M-Z$ relation for SF galaxies at z $<$ 1, Savaglio et al. ("
+2005). have investigated the mass-metallicity relations using galaxies at 04 « z « I. finding that metallicity is lower at higher redshift by ~ 0.15 dex.,"2005), have investigated the mass–metallicity relations using galaxies at 0.4 $<$ z $<$ 1, finding that metallicity is lower at higher redshift by $\sim$ 0.15 dex."
+ Moreover. Maier at al. (," Moreover, Maier at al. ("
+2005). Hammer et al. (,"2005), Hammer et al. ("
+2005). and Liang et al. (,"2005), and Liang et al. ("
+2006) found that emission line galaxies were poorer in metals at z ~ 0.7 than present-day spirals.,2006) found that emission line galaxies were poorer in metals at z $\sim$ 0.7 than present–day spirals.
+ A study of Lamareille et al. (, A study of Lamareille et al. (
+2009) focused on the evolution of the M—Z relation up to z ~ 0.9. suggesting that the M—Z relation is flatter at higher redshifts.,"2009) focused on the evolution of the $M-Z$ relation up to z $\sim$ 0.9, suggesting that the $M-Z$ relation is flatter at higher redshifts."
+ However. Carollo & Lilly (2001). from emission-line ratios of 15 galaxies in a range of 0.5 « z « I. found that their metallicities appear to be remarkably similar to those of local galaxies selected with the same criteria.," However, Carollo $\&$ Lilly (2001), from emission–line ratios of 15 galaxies in a range of 0.5 $<$ z $<$ 1, found that their metallicities appear to be remarkably similar to those of local galaxies selected with the same criteria."
+ Also. Lilly et al. (," Also, Lilly et al. ("
+2003). from a sample of 66 SF galaxies with 0.47 < z « 0.92. claim a smaller variation in metallicity of ~ 0.08 dex compared with the metallicity observed locally. showing only modest evolutionary effects (for more details about the M—Z relation. see Lara-Lóppez et al.,"2003), from a sample of 66 SF galaxies with 0.47 $<$ z $<$ 0.92, claim a smaller variation in metallicity of $\sim$ 0.08 dex compared with the metallicity observed locally, showing only modest evolutionary effects (for more details about the $M-Z$ relation, see Lara-Lóppez et al."
+ 2009b)., 2009b).
+ In a recent study. Calura et al. (," In a recent study, Calura et al. ("
+2009) have demonstrated the importance on the morphology of galaxies when deriving the M—Z relation since. at any redshift. elliptical galaxies present the highest stellar masses and the highest metallicities. whereas the irregulars are the least massive galaxies. characterised by the lowest O abundances.,"2009) have demonstrated the importance on the morphology of galaxies when deriving the $M-Z$ relation since, at any redshift, elliptical galaxies present the highest stellar masses and the highest metallicities, whereas the irregulars are the least massive galaxies, characterised by the lowest O abundances."
+ In this paper. we consistently approach several topics. starting with the introduction. of the S2N2 as a reliable diagram to classify galaxies. the analysis of the metallicity evolution of galaxies in the three BPT diagrams. and. for a better understanding of the processes involved in the observed evolution of galaxies at low redshift. we studied the mass. metallicity and SFR relations. such as the M—Z. metallicity-SFR and mass-SFR relations.," In this paper, we consistently approach several topics, starting with the introduction of the S2N2 as a reliable diagram to classify galaxies, the analysis of the metallicity evolution of galaxies in the three BPT diagrams, and, for a better understanding of the processes involved in the observed evolution of galaxies at low redshift, we studied the mass, metallicity and SFR relations, such as the $M-Z$, metallicity-SFR and mass-SFR relations."
+ We also point out that the morphology of galaxies play an important role when deriving conclusions. since late-type galaxies will result in. lower metallicity estimates and higher SERs than early-type (Calura et al.," We also point out that the morphology of galaxies play an important role when deriving conclusions, since late–type galaxies will result in lower metallicity estimates and higher SFRs than early–type (Calura et al."
+ 2009)., 2009).
+ This paper is structured as follows. in Sect.," This paper is structured as follows, in Sect."
+ 2 we detail the data used for this study. the dust extinction correction and the metallicity estimates for our sample of galaxies. in Sect.," 2 we detail the data used for this study, the dust extinction correction and the metallicity estimates for our sample of galaxies, in Sect."
+ 3 we introduce the S2N2 as a reliable diagram to segregate, 3 we introduce the S2N2 as a reliable diagram to segregate
+brief plateaus or inflection points.,brief plateaus or inflection points.
+ In Section 2 we show that the Dainotti relation Ly;x-. as originally envisioned may be due to the observational bias against detecting and characterizing faint plateaus., In Section 2 we show that the Dainotti relation $L^*_{\rm II} \simpropto {t^*}_{\rm II}^{-1}$ as originally envisioned may be due to the observational bias against detecting and characterizing faint plateaus.
+" Nevertheless, the accreted mass estimates 6M inferred using the and values from Dainotti et al. ("," Nevertheless, the accreted mass estimates $\delta M$ inferred using the $L^*_{\rm II}$ and $t^*_{\rm II}$ values from Dainotti et al. ("
+"2010) are interestingLy; in £jthe context of this work, and one might legitimately ask what theoretical prediction for the Ly,(tj;) relation the fall-back accretion hypothesis would make, given a hypothetical physical constraint in which 6M were held constant and the initial radius of the fall-back disk were allowed to vary.","2010) are interesting in the context of this work, and one might legitimately ask what theoretical prediction for the $L^*_{\rm II}(t^*_{\rm II})$ relation the fall-back accretion hypothesis would make, given a hypothetical physical constraint in which $\delta M$ were held constant and the initial radius of the fall-back disk were allowed to vary."
+ Figure 6 shows the results of seven runs in which we fix 5M=10-4+Mo and vary the initial outer radius of the fall back disk., Figure 6 shows the results of seven runs in which we fix $\delta M= 10^{-4}\msun$ and vary the initial outer radius of the fall back disk.
+ One can see clear inverse relation between the duration of the plateau £j;a and the luminosity at the end of the plateau Lr;., One can see a clear inverse relation between the duration of the plateau $t^*_{\rm II}$ and the luminosity at the end of the plateau $L^*_{\rm II}$.
+" The low—z behavior of 6M seen in Figure 1, however, indicates that a spread in óM may be more realistic, in which case one would not expect to be able to use the theoretical prediction of the Lr;(f1;) relation as a useful discriminant for the theory."," The $-z$ behavior of $\delta M$ seen in Figure 1, however, indicates that a spread in $\delta M$ may be more realistic, in which case one would not expect to be able to use the theoretical prediction of the $L^*_{\rm II}(t^*_{\rm II})$ relation as a useful discriminant for the theory."
+ -0.065 cm We have presented time dependent calculations of the fall-back disk scenario to account for the long-term X-ray light curves for GRBs., -0.065 cm We have presented time dependent calculations of the fall-back disk scenario to account for the long-term X-ray light curves for GRBs.
+" For IGRBs, an initial radial scale of a stellar radius ~10!! cm gives a natural viscous evolution time of —10* s to redistribute the matter into a quasi-steady disk, roughly consistent with the observed plateau duration."," For lGRBs, an initial radial scale of a stellar radius $\sim10^{11}$ cm gives a natural viscous evolution time of $\sim10^4$ s to redistribute the matter into a quasi-steady disk, roughly consistent with the observed plateau duration."
+" GRB 060729 which we model in the context of a IGRB had one of the longest plateaus ever observed, thus more generally we anticipate progenitor radii, or more specifically, circularization radii, given by ~4/jtot/JeritRx, to lie in the range ~101?—1011 cm."," GRB 060729 which we model in the context of a lGRB had one of the longest plateaus ever observed, thus more generally we anticipate progenitor radii, or more specifically, circularization radii, given by $\sim \sqrt{j_{\rm tot}/j_{\rm crit}}R_*$, to lie in the range $\sim 10^{10} - 10^{11}$ cm."
+" The rate of decay for GRB 060729 is close but does not match the models precisely, which may hint at time-variable emission processes that would affect our adopted efficiencies €,-, and f."," The rate of decay for GRB 060729 is close but does not match the models precisely, which may hint at time-variable emission processes that would affect our adopted efficiencies $\epsilon_{\rm net}$ and $f$."
+" It is interesting to note that our effective values of the efficiencies, accretion plus to match the observed X-ray flux levels are beaming,comparable requiredbetween IGRBs and sGRBs: fed.c&10-7? for GRB 060729 and fe,c3x107? for GRB 0512214. If the accretion efficiency έποι is about the same for IGRBs and sGRBs, this may indicate that SGRBs are less beamed by factor of ~3 compared to IGRBs, roughly in line with previousa results (Watson et al."," It is interesting to note that our effective values of the efficiencies, accretion plus beaming, required to match the observed X-ray flux levels are comparable between lGRBs and sGRBs: $f \epsilon_{\rm net}^{-1} \simeq 10^{-2}$ for GRB 060729 and $f \epsilon_{\rm net}^{-1} \simeq 3\times 10^{-2}$ for GRB 051221A. If the accretion efficiency $\epsilon_{\rm net}$ is about the same for lGRBs and sGRBs, this may indicate that sGRBs are less beamed by a factor of $\sim3$ compared to lGRBs, roughly in line with previous results (Watson et al."
+" 2006, Grupe et al."," 2006, Grupe et al."
+ 2006)., 2006).
+" This similarity between IGRBs and sGRBs afterglows is not entirely unexpected, given the general similarities in"," This similarity between lGRBs and sGRBs afterglows is not entirely unexpected, given the general similarities in"
+"The complex. nebulosity surrounding 1) Carinae represents the remnants of ejecta from this Luminous Blue Variable star (10"" L. Davidson et al. 1986) during irregular outbursts over. the last 300 ves (Walborn Liller 1977).","The complex nebulosity surrounding $\eta$ Carinae represents the remnants of ejecta from this Luminous Blue Variable star $^{6.6}$ $_{\odot}$ – Davidson et al, 1986) during irregular outbursts over the last 300 yrs (Walborn Liller 1977)."
+ Investigations of the kinematies. morphologies and ionization mechanisms of the distinct regions identified in the sketch in Fig.," Investigations of the kinematics, morphologies and ionization mechanisms of the distinct regions identified in the sketch in Fig."
+ 1 that comprise this nebulosity have been reported in a series of papers (hereafter referred to as Papers ] 31e., 1 that comprise this nebulosity have been reported in a series of papers (hereafter referred to as Papers 1 – 3 ie.
+ Meaburn. Wolstenceroft Walsh (1987). Aleaburn. Walsh Wolsteneroft (1993 a) ancl Aleaburn et al (1993 b) respectively).," Meaburn, Wolstencroft Walsh (1987), Meaburn, Walsh Wolstencroft (1993 a) and Meaburn et al (1993 b) respectively)."
+ There is the. clusty bipolar. Homunculus expanding at z 650 kms (Paper 2) surrounded. by a shocked. nitrogen enriched inner shell (identified by Walborn et al. LOTS as 5 ridge ancl Ware) with receding. expansive motions of z 350 kms.! (Paper 1 and see Dufour. 1989).," There is the dusty bi–polar Homunculus expanding at $\approx$ 650 $\kms$ (Paper 2) surrounded by a shocked, nitrogen enriched inner shell (identified by Walborn et al, 1978 as S ridge and W arc) with receding, expansive motions of $\approx$ 350 $\kms$ (Paper 1 and see Dufour, 1989)."
+ A knotty jet projects through the NE quadrant of this inner shell (Hester et al 1991) whose episodic. outllow velocity is « 1500 kms+ (Paper 3).," A knotty `jet' projects through the NE quadrant of this inner shell (Hester et al 1991) whose episodic, outflow velocity is $\leq$ 1500 $\kms$ (Paper 3)."
+ Outside these inner regions a variety of. distinctly different phenomena are clearly apparent on Ulx Schmidt plates. Anglo-Australian ��Telescope (AAT) prime focus plates and the most recent LIST images (Fig.," Outside these inner regions a variety of distinctly different phenomena are clearly apparent on UK Schmidt plates, Anglo-Australian Telescope (AAT) prime focus plates and the most recent HST images (Fig."
+ 2 ancl see Sect., 2 and see Sect.
+ 2)., 2).
+ There is the irregular. olfset outer ‘shell’ of knots xwtially identified by Walborn et al (LOTS).," There is the irregular, off–set outer `shell' of knots partially identified by Walborn et al (1978)."
+ A feature now markecl ‘spike’ in Fig., A feature now marked `spike' in Fig.
+" 1 (prematurely identified as a jet in ""aper L also see fig.", 1 (prematurely identified as a jet in Paper 1 – also see fig.
+" (ο) in Paper 2) projects through this irregular outer ""shell.", 1(c) in Paper 2) projects through this irregular outer `shell'.
+ This can be seen in the LST image in Fig., This can be seen in the HST image in Fig.
+ 2 to be & 0.37 wide., 2 to be $\approx$ $^{\prime\prime}$ wide.
+ Also a potentially interesting eature of the outer nebulosity is the faint arc of nebulosity. marked ‘are’ in Fie.," Also a potentially interesting feature of the outer nebulosity is the faint arc of nebulosity, marked `arc' in Fig."
+ P and outlined by a dashed line., 1 and outlined by a dashed line.
+" This can be seen on a deep AAT prime focus plate to extend well xvond the perimeter of the outer shell in the vicinity of the ""spike! and to originate apparently in the vicinity of knot I5 (Walborn et al LOTS).", This can be seen on a deep AAT prime focus plate to extend well beyond the perimeter of the outer shell in the vicinity of the `spike' and to originate apparently in the vicinity of knot E5 (Walborn et al 1978).
+ The investigation of the dramatic kinematics of 1) some of the prominent knots of the outer ‘shell. ii) the spike’ (alreacly indicated in Paper 1) and ii) the ‘are’ is reported in the present work.," The investigation of the dramatic kinematics of i) some of the prominent knots of the outer `shell', ii) the `spike' (already indicated in Paper 1) and iii) the `arc' is reported in the present work."
+ Spatially resolved: profiles of the La and 1]A6548 ancl 6584 lines have now been obtained. over these regions with the LEAIALL spectrometer on the ESO New Technology Telescope (NEL) in Chile., Spatially resolved profiles of the $\alpha$ and $\lambda$ 6548 and 6584 lines have now been obtained over these regions with the EMMI spectrometer on the ESO New Technology Telescope (NTT) in Chile.
+ Incidentally. a cistance to 5 Carinae of 2.6 kpe will be adopted here to be consistent with Papers 2 3.," Incidentally, a distance to $\eta$ Carinae of 2.6 kpc will be adopted here to be consistent with Papers 2 3."
+ This value is within the error bars of both the measurement of 2.5 kpc, This value is within the error bars of both the measurement of 2.5 kpc
+Consider black holes with oneelectric charge qy,"instance, if only one of the $D_{ABC}$ 's (up to permutations) is nonzero,all $\gamma^A$ 's may be chosen For the case of either equal $\gamma^A$ $A=1,2,3$ ) or only $D_{123}\neq0$, the auxiliary field $T_{ab}^-$ takes the form where Solving the stabilization equations, one obtains the moduli on the horizon in terms of the charges and the moduli at infinity."
+andp*(A= 1...NV)magnetic charges. One AC," The Bekenstein-Hawking entropy takes the form where This has the same form as the extremal entropy, with the charges $(q_0,p^A)$ replaced by $(\qs_0,\ps^A)$."
+ h'pgusmh4fel utlactcoshs act (nosumni," Note that, unlike the extremal case, the entropy depends on the values of the moduli at infinity."
+ation) TE where (qu =hy sinh-ση =Ouf. (25) altractor ," In addition, the non-extremal entropy has a different functional dependence on the original charges since the parameters $\alpha^A,\alpha_0$ themselves depend on the charges for given asymptotic moduli."
+- (Fils)E) I. (26) where 17.Hj ave harmonic [unctions IH! pi LE Th= hj., The near-extremal black holes are described by adding to the extremal black holes the leading terms in $\mu$ while holding the physical charges fixed.
+ (27), One gets
+"(20)) do not account for observational error, and hence we do not use them to solve for m and s estimates for our relatively low signal-to-noise BOSS survey data.","\ref{eq:var_stack}) ) do not account for observational error, and hence we do not use them to solve for $m$ and $s$ estimates for our relatively low signal-to-noise BOSS survey data."
+" However, our subsample of high signal-to-noise 7-hour observations allows us to test them, which we do before proceeding to the verification of our Bayesian analysis framework."," However, our subsample of high signal-to-noise 7-hour observations allows us to test them, which we do before proceeding to the verification of our Bayesian analysis framework."
+" First, we use Equation (19)) with a mean of 2.33 and an intrinsic scatter of 0.07 to predict a value of stack=219kms-!, which is in very good agreement with the result of +kms! that we obtain by fitting the stacked spectrum(222 12)of this set of 125 galaxies directly."," First, we use Equation \ref{eq:sigma_stack}) ) with a mean of 2.33 and an intrinsic scatter of 0.07 to predict a value of $\sigma_{\rm stack}=\rm 219\,km\,s^{-1}$, which is in very good agreement with the result of $\rm (222 \pm 12) \,km\,s^{-1}$ that we obtain by fitting the stacked spectrum of this set of 125 galaxies directly."
+" Similarly, we predict [Var(o2.,,.)]/4=40kms-! from Equation (20)), which is in reasonable agreement with the value of 46kms~! obtained through a bootstrap resampling process."," Similarly, we predict $[\mathrm{Var}(\sigma_{\mathrm{stack}}^2)]^{1/4} = 40 \rm\,km\,s^{-1}$ from Equation \ref{eq:var_stack}) ), which is in reasonable agreement with the value of $46 \rm \,km\,s^{-1}$ obtained through a bootstrap resampling process."
+" In both cases the agreement is not exact because there is stil some observational error even in the 7-hour data, but as mentioned above, we will pass to the Bayesian framework to quantify these effects."," In both cases the agreement is not exact because there is still some observational error even in the 7-hour data, but as mentioned above, we will pass to the Bayesian framework to quantify these effects."
+" We next carry out the estimation of the m and s parameters of the selected sub-sample of objects, using the Bayesian approach described above, for both the 7-hour and the 1.75-hour data sets."," We next carry out the estimation of the $m$ and $s$ parameters of the selected sub-sample of objects, using the Bayesian approach described above, for both the 7-hour and the 1.75-hour data sets."
+ Figure 4 shows the resulting posterior probability density for these parameters as estimated from both data sets.," Figure \ref{fig:bayesVerify}
+ shows the resulting posterior probability density for these parameters as estimated from both data sets."
+" As expected, we see that the posterior PDF is tighter for the 7-hour data."," As expected, we see that the posterior PDF is tighter for the 7-hour data."
+" More importantly, we see no significant bias in the posterior PDF between the low-SNR and high-SNR data sets."," More importantly, we see no significant bias in the posterior PDF between the low-SNR and high-SNR data sets."
+" This is especially significant for the estimation of the s parameter: if we were not handling our observational uncertainties correctly, we might expect to infer a broader intrinsic distribution s from the noisier data, but this not the case."," This is especially significant for the estimation of the $s$ parameter: if we were not handling our observational uncertainties correctly, we might expect to infer a broader intrinsic distribution (higher $s$ value) from the noisier data, but this not the case."
+ (higherWe also value)see that the parameters used to engineer the subsample are recovered with no significant bias in m., We also see that the parameters used to engineer the subsample are recovered with no significant bias in $m$.
+ We see slight offset of the 7-hour maximum posterior s value froma the input value used to engineer the sample., We see a slight offset of the 7-hour maximum posterior $s$ value from the input value used to engineer the sample.
+" This is in the direction and of the size to expected given the observational error of the 7-hour individual-spectrum velocity dispersion measurements, which have an RMS signal-to-noise of about 17."," This is in the direction and of the size to expected given the observational error of the 7-hour individual-spectrum velocity dispersion measurements, which have an RMS signal-to-noise of about 17."
+" This corresponds to an observational broadening of about ddex in the engineered histogram of Figure 3,, which is deconvolved by the Bayesian parameter estimation procedure to give the lower recovered s value seen in Figure 4.."," This corresponds to an observational broadening of about dex in the engineered histogram of Figure \ref{fig:vdTestHist}, which is deconvolved by the Bayesian parameter estimation procedure to give the lower recovered $s$ value seen in Figure \ref{fig:bayesVerify}."
+" Another concern is that there might be a systematic bias with redshift, since the spectral regions used by in fitting for velocity dispersions (rest frame wavelength range from tto move to the redder and noisier parts of the 6,800A))spectrum as the redshift gets higher."," Another concern is that there might be a systematic bias with redshift, since the spectral regions used by in fitting for velocity dispersions (rest frame wavelength range from to ) move to the redder and noisier parts of the spectrum as the redshift gets higher."
+" In order to test this, we construct another controlled sub-sample with 152 galaxies of redshift z<0.2 and very high SNR."," In order to test this, we construct another controlled sub-sample with 152 galaxies of redshift $z < 0.2$ and very high SNR."
+" Then we take the best-fit template combination models of those 152 galaxies returned by and redshift them to progressively higher redshift bins, giving them a uniform random distribution over a bin width of Az=0.04 in each case match our actual "," Then we take the best-fit template combination models of those 152 galaxies returned by and redshift them to progressively higher redshift bins, giving them a uniform random distribution over a bin width of $\Delta z = 0.04$ in each case (to match our actual binning)."
+"At each new redshift, the (tomodel spectra are added to binning).sky-subtracted BOSS sky fibers to simulate realistic survey noise, and scaled individually in flux to give a typical median SNR at that redshift bin."," At each new redshift, the model spectra are added to sky-subtracted BOSS sky fibers to simulate realistic survey noise, and scaled individually in flux to give a typical median SNR at that redshift bin."
+ We then analyse the simulated redshifted samples with our Bayesian method to estimate the posterior PDF of m and s., We then analyse the simulated redshifted samples with our Bayesian method to estimate the posterior PDF of $m$ and $s$.
+" The results are shown in Figure 5 and Figure 6,, for 5 separate redshift bins."," The results are shown in Figure \ref{fig:vdisp_verify1} and Figure \ref{fig:vdisp_verify2}, for 5 separate redshift bins."
+" We see that the recovered parameters are consistent within observational error across all redshifts, with no apparent redshift-dependent bias."," We see that the recovered parameters are consistent within observational error across all redshifts, with no apparent redshift-dependent bias."
+" Finally, to rule out any significant dependence of our measurement on airmass and fiber position within the BOSS spectrographs, we make use of data from plates 3615, 3647, 4238 and 4239."," Finally, to rule out any significant dependence of our measurement on airmass and fiber position within the BOSS spectrographs, we make use of data from plates 3615, 3647, 4238 and 4239."
+" These four plates cover roughly the same area of sky, but with different plate drillings that place the same objects in very different fibers within the spectrograph system."," These four plates cover roughly the same area of sky, but with different plate drillings that place the same objects in very different fibers within the spectrograph system."
+ They were also observed over a range of different airmasses on multiple nights., They were also observed over a range of different airmasses on multiple nights.
+" From these plates, we construct several sub-samples of spectra, all of which include the same galaxies, but are drawn from different plates and/or observations."," From these plates, we construct several sub-samples of spectra, all of which include the same galaxies, but are drawn from different plates and/or observations."
+" As with the previous tests, we recover consistent estimates of m and s from the analysis of all these samples."," As with the previous tests, we recover consistent estimates of $m$ and $s$ from the analysis of all these samples."
+Although the M dwarfs constitute a large fraction of the detectable baryonic matter. we still lack a great deal of knowledge about our low-mass ἐς 0.6M..) hydrogen-burning neighbours.,"Although the M dwarfs constitute a large fraction of the detectable baryonic matter, we still lack a great deal of knowledge about our low-mass $<0.6M_{\odot}$ ) hydrogen-burning neighbours."
+ Studies (??) suggest that as much as of all stars in the solar vicinity (~10ppe) are M. dwarfs which nakes these objects essential when deriving quantities such as the initial mass function (IMF.?) and the present day nass function.," Studies \citep{1998ASPC..134...28H,2003PASP..115..763C} suggest that as much as of all stars in the solar vicinity $\sim$ pc) are M dwarfs which makes these objects essential when deriving quantities such as the initial mass function \citep[IMF,][]{1955ApJ...121..161S} and the present day mass function."
+ These frequently used functions are derived using the luminosity., These frequently used functions are derived using the luminosity.
+ The transformation to mass. based upon stellar evolution theories. is sensitive to chemical composition.," The transformation to mass, based upon stellar evolution theories, is sensitive to chemical composition."
+ Aoreover. a study of the possible time dependence of the IMF function for the low-mass part. needs good stellar metallicity criteria for M dwarfs.," Moreover, a study of the possible time dependence of the IMF function for the low-mass part, needs good stellar metallicity criteria for M dwarfs."
+ Thus. if we are to create a realistic model of the galactic evolution and present day status. detailed studies of these faint but numerous objects are of great interest.," Thus, if we are to create a realistic model of the galactic evolution and present day status, detailed studies of these faint but numerous objects are of great interest."
+ The M dwarfs are needed in the understanding of main-sequence stellar evolution and to define a limit between stellar and substellar objects., The M dwarfs are needed in the understanding of main-sequence stellar evolution and to define a limit between stellar and substellar objects.
+ Finally. a well defined metallicity scale for M-dwarfs 1s essential to determine whether or not the general trend towards supersolar metallicities among FGK-stars planet hosts (e.g.?) holds also for cooler objects.," Finally, a well defined metallicity scale for M-dwarfs is essential to determine whether or not the general trend towards supersolar metallicities among FGK-stars planet hosts \citep[e.g.][]{2005ApJ...622.1102F}
+ holds also for cooler objects."
+ Spectroscopic studies of M dwarfs at high resolution have proven to be a difficult task., Spectroscopic studies of M dwarfs at high resolution have proven to be a difficult task.
+ In the low-temperature regime occupied by these targets (2000z Ty.4100 KK). the optical spectrum is covered by a forest of molecular lines. hiding or blending most of the atomic lines used in spectral analysis.," In the low-temperature regime occupied by these targets $\lesssim$ $_{\rm eff}\lesssim$ K), the optical spectrum is covered by a forest of molecular lines, hiding or blending most of the atomic lines used in spectral analysis."
+ However. models of low-mass late-type stars have undergone continuous improvements. from the early work by ?.. ? and ?2?.. to the work by ?.. 22. 2.. ?. with their extensive grid of NextGen models. and the improved MARCS models (?)..," However, models of low-mass late-type stars have undergone continuous improvements, from the early work by \citet{1969lls..symp..457T}, , \citet{1969ApJ...157..799A} and \citet{1975A&A....38..283M,1976A&A....48..443M}, to the work by \citet{1993PASAu..10..250B}, \citet{1995A&A...295..736B,1995A&AS..109..263B}, \citet{1995ApJ...445..433A}, , \citet{1999ApJ...512..377H} with their extensive grid of NextGen models, and the improved MARCS models \citep{2008A&A...486..951G}."
+ New molecular and atomic line data are collected and organised in laree databases such as VALD (?) and VAMDC (?) and will improve the situation further., New molecular and atomic line data are collected and organised in large databases such as VALD \citep{1999A&AS..138..119K} and VAMDC \citep{2010JQSRT.111.2151D} and will improve the situation further.
+ The faintness characterising the M dwarfs has limited the number of high-resolution. high signal-to-noise studies.," The faintness characterising the M dwarfs has limited the number of high-resolution, high signal-to-noise studies."
+ Non-sufficient resolutiol and dominant molecular features make it difficult to derive accurate atomic line strengths needed for a reliable metallicity determination., Non-sufficient resolution and dominant molecular features make it difficult to derive accurate atomic line strengths needed for a reliable metallicity determination.
+ Metallicities have been derived via photometric calibrations (2).. studies using molecular indices (?).. as well as spectrum synthesis (??) anc spectroscopic calibrations in the K band (?)..," Metallicities have been derived via photometric calibrations \citep{2005A_A...442..635B}, studies using molecular indices \citep{2006PASP..118..218W}, as well as spectrum synthesis \citep{2006ApJ...653L..65B,2006ApJ...652.1604B} and spectroscopic calibrations in the K band \citep{2010ApJ...720L.113R}."
+ In this paper we present a detailed spectroscopic study i1 the J-band (1100—1400 nnm). a spectral region relatively free from molecular lines.," In this paper we present a detailed spectroscopic study in the J-band $1100-1400$ nm), a spectral region relatively free from molecular lines."
+ In the near infrared. atomic lines cat be isolated and the lack of molecular lines allows a precise continuum placement.," In the near infrared, atomic lines can be isolated and the lack of molecular lines allows a precise continuum placement."
+ This spectral region was exploited for abundance analysis in the pioneering study of Betelgeuse. based on FTS spectra. by ?.. but has not been used much for late-type dwarfs until the last decade due to lack of efficient IR. spectrometers at large telescopes.," This spectral region was exploited for abundance analysis in the pioneering study of Betelgeuse, based on FTS spectra, by \citet{VieiraThesis}, but has not been used much for late-type dwarfs until the last decade due to lack of efficient IR spectrometers at large telescopes."
+ The present generation of IR echelles at large telescopes have. however. opened new possibilities. see. e.g. ??..," The present generation of IR echelles at large telescopes have, however, opened new possibilities, see, e.g. \citet{2000ApJ...533L..45M,2007ApJ...658.1217M}."
+ Similarly to previous studies (??).. we select a number of binary systems with a solar-type primary and an M-dwarf companion to assess the accuracy of the atmospheric models used in the analysis and to verify the atomic line treatment.," Similarly to previous studies \citep{2006ApJ...652.1604B,2005A_A...442..635B}, we select a number of binary systems with a solar-type primary and an M-dwarf companion to assess the accuracy of the atmospheric models used in the analysis and to verify the atomic line treatment."
+ We then apply this result to a number of non-binary M dwarfs., We then apply this result to a number of non-binary M dwarfs.
+ Our choice of spectral region makes a careful analysis possible. for a sample of stars that are thought to have a high metal content (—0.35 to 0.5 dex). avoiding the large contribution of molecules such as TiO to the spectrum at optical wavelengths.," Our choice of spectral region makes a careful analysis possible, for a sample of stars that are thought to have a high metal content $-$ 0.35 to 0.5 dex), avoiding the large contribution of molecules such as TiO to the spectrum at optical wavelengths."
+ The paper is organized as follows., The paper is organized as follows.
+ In Section 2. we briefly summarize previous metallicity studies of M dwarfs., In Section \ref{sec:prevstud} we briefly summarize previous metallicity studies of M dwarfs.
+ In Section 3.. we describe the programme stars. the observations. and some aspects of the data reduction.," In Section \ref{sect:observations}, we describe the programme stars, the observations, and some aspects of the data reduction."
+ In Section. 4.. the ingredients of the spectrum analysis are presented — compilation and derivation of spectral line data. stellar atmospheric parameters. and the procedure for metallicity estimation.," In Section \ref{sect:analysis}, the ingredients of the spectrum analysis are presented -- compilation and derivation of spectral line data, stellar atmospheric parameters, and the procedure for metallicity estimation."
+ In Section 5 we discuss the results of the analysis. and Section 6 concludes the paper.," In Section \ref{sect:results} we discuss the results of the analysis, and Section \ref{sect:conclusions} concludes the paper."
+ The stars in our sample have been investigated with various methods differentfrom ours byseveral authors., The stars in our sample have been investigated with various methods differentfrom ours byseveral authors.
+ In. the following. we give a brief description of these studies.," In the following, we give a brief description of these studies."
+As one can see that the current cnerey of electrons is independent. of their injection energy. fy and. equals 1/2 for long enough /.,As one can see that the current energy of electrons is independent of their injection energy $E_0$ and equals $E\simeq 1/\beta t$ for long enough $t$.
+ From LEq.(18)) it follows that the radio flux just after 10 injection of positrons equals 38+10 Jv for the Nyc:107 in each capture events that is necessary to produce 1¢ observed annihilation Dux from the GC., From \ref{nu_sh}) ) it follows that the radio flux just after the injection of positrons equals $\sim 3 \cdot 10^7$ Jy for the $N_0\simeq 2\cdot 10^{55}$ in each capture events that is necessary to produce the observed annihilation flux from the GC.
+ Llowever. for the time /2LO’ vears (the average period between str captures) re peak of intensity shifts from the frequency. LOO MIIZ to several Alllz where this radio llux cannot be observed because of absorption in the interstcllar gas. ancl just this cHeet is a key point of our analysis presented. below.," However, for the time $t\simeq 10^5$ years (the average period between str captures) the peak of intensity shifts from the frequency 100 MHz to several MHz where this radio flux cannot be observed because of absorption in the interstellar gas, and just this effect is a key point of our analysis presented below."
+ As an example we show in Lig., As an example we show in Fig.
+ 2. the spectrum racio emission from the GC region produced. by secondary electrons and positrons at the time 10? vr after the capture., \ref{s_radio} the spectrum radio emission from the GC region produced by secondary electrons and positrons at the time $10^5$ yr after the capture.
+ The magnetic field strength in the GC is 2 and 3 mG. As one can see for long time alter the capture the peak of emission shifts from the frequcney of hundreds Mllz to several MllIz. and the intensity of radio emission is negligible at 330 MlIIz.," The magnetic field strength in the GC is 2 and 3 mG. As one can see for long time after the capture the peak of emission shifts from the frequency of hundreds MHz to several MHz, and the intensity of radio emission is negligible at 330 MHz."
+ For calculations of racio emission we used the accurate equation (frome.g.Berezinskiietal.1990).. where ANGE) is the total number of electrons. and positrons with the energy £. d=SN kpe is the distance to the GC. and AuGr) is the MeDonald function.," For calculations of radio emission we used the accurate equation \citep[from e.g.][]{ber90}, where $N(E)$ is the total number of electrons and positrons with the energy $E$, $d=8$ kpc is the distance to the GC, and $K_\alpha(x)$ is the McDonald function."
+ In Fig., In Fig.
+ 3. we presented the limitations of positron injection energy fy derived from radio data for cillerent values of the magnetic field. strength. and. dillerent. times J passed from the last star capture., \ref{lim_radio} we presented the limitations of positron injection energy $E_0$ derived from radio data for different values of the magnetic field strength and different times T passed from the last star capture.
+ We considered. two possible cases: the injection in the form of delta-function. Eq. (113).," We considered two possible cases: the injection in the form of delta-function, Eq. \ref{delta_inj}) ),"
+ shown in the figure by solid lines ancl power-LIaw injection spectrum (dashed line): One can see that if the injection is non-stationary. then the injection cnerey of positrons can be extremely hieh in case of strong magnetic fields in the GC.," shown in the figure by solid lines and power-law injection spectrum (dashed line): One can see that if the injection is non-stationary, then the injection energy of positrons can be extremely high in case of strong magnetic fields in the GC."
+ Indeed for almost stationary situation when characteristic period between injections is small Z7<1107 vr the maximum allowed. energy in case of ff1 mC cannot exceed: 30 MeV clue to radio limitations., Indeed for almost stationary situation when characteristic period between injections is small $T \leq 1\times10^4$ yr the maximum allowed energy in case of $H \geq 1$ mG cannot exceed 30 MeV due to radio limitations.
+ In case of power-law spectrum situation is even worse., In case of power-law spectrum situation is even worse.
+" However if period is long enough and magnetic field is strong the situation changes: for =1.10"" vr απ {4=2 mG the maximum energy is about 1 GeV and it can be even higher for longer periods and higher values of the magnetic field.", However if period is long enough and magnetic field is strong the situation changes: for $T=1\times 10^5$ yr and $H=2$ mG the maximum energy is about 1 GeV and it can be even higher for longer periods and higher values of the magnetic field.
+" ‘There are restrictions for our model which follow from racio and egamma-observations: ‘These restrictions are presented in ""Table 1..", There are restrictions for our model which follow from radio and gamma-observations: These restrictions are presented in Table \ref{param_st1}. .
+pixels aud 0.17 0. for the smallest bins used here.,pixels and 0.17 $\rm e^-$ for the smallest bins used here.
+Noise The averaged nightly dat used here is comprised of 12 doime-flats with total counts about 810.000 οfpixcl. vieldiug a photon noise ckwe toJor 0.1 epixel| equivalent skv backeround counts.," The averaged nightly flat used here is comprised of 12 dome-flats with total counts about 840,000 $\rm e^-/pixel$, yielding a photon noise close to, or 330.1 $\rm e^{-} pixel^{-1}$ equivalent sky background counts."
+ Civen that that final nuage includes 150 shifted. averaged flats the error per pixel is reduced to0.," Given that that final image includes 150 shifted, averaged flats the error per pixel is reduced to."
+009%.. This part of the flat field error is random in nature and trauslates to a formal error of for the largest sample bius used here. to for the siallest bius.," This part of the flat field error is random in nature and translates to a formal error of for the largest sample bins used here, to for the smallest bins."
+ The stability of the large-scale flat field was checked by comparing the dome-flats from adjacent mights., The stability of the large-scale flat field was checked by comparing the dome-flats from adjacent nights.
+ The average results of these tests iudicate an error of0.03:4.. or close to 90 o0pixel4 for our measured sky backeround.," The average results of these tests indicate an error of, or close to 90 $\rm e^{-}
+pixel^{-1}$ for our measured sky background."
+ In our CCD system. the opening aud closing time of shutter is 20 milliseconds aud the exposure time of cach dome-fat is 150 seconds.," In our CCD system, the opening and closing time of shutter is 20 milliseconds and the exposure time of each dome-flat is 150 seconds."
+ So there exists a gradient due to finite shutter speed in the each dome-flat frame., So there exists a gradient due to finite shutter speed in the each dome-flat frame.
+ Combining these two sets of errors. we estimate the error from systematic variatious iu the large-scale flat-field is 0.033%... equivaleut to 100 6pix:el+ of ομι ," Combining these two sets of errors, we estimate the error from systematic variations in the large-scale flat-field is equivalent to 100 $\rm e^{-} pixel^{-1}$ of sky background."
+Following MDBII (Morrisonotal. 19913) and Zheug et al.," Following MBH \cite{MBH94}) ) and Zheng et al.,"
+ we can calculate the raudon eror ¢uc to intrinsic variation in the surface brightuess of the galaxy., we can calculate the random error due to intrinsic variation in the surface brightness of the galaxy.
+ Iu the terminology introduced by Toury ancl Schueider (1988). wil for our image is 20.30 mae. the exposure time is 15[0123 sec. the distance is 15 AIpe and Moooo is adopted as 0 (Zhengctal. 1999)).," In the terminology introduced by Tonry and Schneider (1988), m1 for our image is 20.30 mag, the exposure time is 154,043 sec, the distance is 14.5 Mpc and $\rm \bar{M}_{6660}$ is adopted as 0 \cite{Zheng99}) )."
+ In the case of 300 οpixel31 from the ealaxy. (equd toa surface briehtuess of 28.23 mag aresee 7). this error is 53.8 οpixelο. or an error of l8translates QO all average error of 1.20 ο for the larecst bins used here. to 10.98 e. for tlre sinallest bius.," In the case of 300 $\rm e^- pixel^{-1}$ from the galaxy (equal to a surface brightness of 28.23 mag $\rm arcsec^{-2}$ ), this error is 53.8 $\rm e^{-} pixel^{-1}$, or an error of translates to an average error of 1.29 $\rm e^- $ for the largest bins used here, to 10.98 $\rm e^- $ for the smallest bins."
+ There are two types of backerotud subtraction errors. both dominaed by systematic effects;," There are two types of background subtraction errors, both dominated by systematic effects."
+ One is the accuracy of the sky background subraction. for which we adopt the vaue of 115 obtained above.," One is the accuracy of the sky background subtraction, for which we adopt the value of 115 $^{-}$ obtained above."
+ The second comes frou imperfect star stbtraction. especially for saturated stars.," The second comes from imperfect star subtraction, especially for saturated stars."
+ For this we adopt an average error of 100 /pixel by checking tlie regions in which stars are subtraced., For this we adopt an average error of 100 $^{-}$ /pixel by checking the regions in which stars are subtracted.
+ We acknowledge this error can be siguificautly higher for selected sars. especially for the saturated stars which are so close to NGC 1565.," We acknowledge this error can be significantly higher for selected stars, especially for the saturated stars which are so close to NGC 4565."
+ For the true errors dwolved in star subtraction. we must look to the coasisteney of the luninosity profiles at faint surface brightucss levels.," For the true errors involved in star subtraction, we must look to the consistency of the luminosity profiles at faint surface brightness levels."
+ All the sources of errors discussed are listed in Table 2., All the sources of errors discussed are listed in Table 2.
+ We assume that the mean count per pixel from the object galaxy is 300 (1.0. 1n its faint halo).a sky level of 302.500 and bin sizes of «50 pixels," We assume that the mean count per pixel from the object galaxy is 300 $^{-}$ (i.e., in its faint halo),a sky level of 302,500 $^{-}$ and bin sizes of $\times$ 50 pixels"
+The generation of internal gravity waves depends ou the structure of the stellar convective envelope.,The generation of internal gravity waves depends on the structure of the stellar convective envelope.
+ For the stars we consider here this region is relatively extended as can be seen on Fig., For the stars we consider here this region is relatively extended as can be seen on Fig.
+ 1 (see also Table 1) where some of their characteristics are plotted as a fiction of effective temperature both ou the zero age main sequence (zauis: open squares) and at LO Cir (black squares)., 1 (see also Table 1) where some of their characteristics are plotted as a function of effective temperature both on the zero age main sequence (zams; open squares) and at 10 Gyr (black squares).
+ At a given age. the size of the couvective envelope decreases when one goes to higher effective temperature.," At a given age, the size of the convective envelope decreases when one goes to higher effective temperature."
+" For a given stellar mass. as the star evolves the effective temperature ries and thus the thickness of the convective envelope decreases,"," For a given stellar mass, as the star evolves the effective temperature rises and thus the thickness of the convective envelope decreases."
+ Let us uote that not all the stars for which we made the present computations will lie on the Lthimm plateau itself., Let us note that not all the stars for which we made the present computations will lie on the lithium plateau itself.
+ If we consider indeed an age of LO Car. only the stars with masses between e0.7M. and ~0.82ML. (for [Fe/U]= 2) have an effective tempcrature in the plateau range (.6.. between ~ 5700 aud 6GS5O0KN). while cooler oues correspond to stars that exhibit lower lithium abundances and more massive ones have already evolved off the main sequence.," If we consider indeed an age of 10 Gyr, only the stars with masses between $\sim 0.7~{\rm M}_{\odot}$ and $\sim 0.82~{\rm M}_{\odot}$ (for $[{\rm Fe/H}]=-2$ ) have an effective temperature in the plateau range (i.e., between $\sim$ 5700 and 6500K), while cooler ones correspond to stars that exhibit lower lithium abundances and more massive ones have already evolved off the main sequence."
+ In this limited mass range Geliüch is dudicated by the shaded. area in Fig., In this limited mass range (which is indicated by the shaded area in Fig.
+ 3 where we show the corresponding rauge in effective teniperature on the zanis). the thickness of the couvective envelope varies only modestly: from —20% to —1:( of the stellar radius for the gams models.," 3 where we show the corresponding range in effective temperature on the zams), the thickness of the convective envelope varies only modestly: from $\sim 20 \%$ to $\sim 12 \%$ of the stellar radius for the zams models."
+" Iun the sue mass rauge aud still ou the zin. the characteristic convective frequency varies from vo ~L.6 plz to ~12 plz while the spherical harmonic munber (,. corresponding to the largest couvective scale varies from 80 to ~1 LO."," In the same mass range and still on the zams, the characteristic convective frequency varies from $\nu_c \sim$ 1.6 $\mu$ Hz to $\sim$ 12 $\mu$ Hz while the spherical harmonic number $\ell_c$ corresponding to the largest convective scale varies from $\sim$ 80 to $\sim$ 140."
+ Iu order to investigate how the wave spectruni produced by these convection zones varies as a fiction of stellar nass. let us first discuss waves characteristics based on the stellar models on the zero age nain sequence.," In order to investigate how the wave spectrum produced by these convection zones varies as a function of stellar mass, let us first discuss waves characteristics based on the stellar models on the zero age main sequence."
+ The corresponding moment spectra are shown in Fie., The corresponding momentum spectra are shown in Fig.
+ 2 where C;=ἐπD isB the momentum Duuunuositv..- ro and Fy beine respectively the radius at the vase of the convective envelope. aud the momentum flux.," 2 where ${\cal L}_J=4\pi r_{\rm cz}^2 {\cal F}_J$ is the momentum luminosity, $r_{\rm cz}$ and ${\cal F}_J$ being respectively the radius at the base of the convective envelope and the momentum flux."
+ We sec that the spectra characteristics. which depend on the structure of the convection zone. evolve with sJar mass.," We see that the spectrum characteristics, which depend on the structure of the convection zone, evolve with stellar mass."
+ Iu particular. low frequency waves disappear when the stellar nass dmuoreases due to a larger thermal diffusivity below he thinner convective cuvelope (see Fig.," In particular, low frequency waves disappear when the stellar mass increases due to a larger thermal diffusivity below the thinner convective envelope (see Fig."
+ 3) that leads to stronger damping (Eq., 3) that leads to stronger damping (Eq.
+ 1)., 1).
+ In addition. the fiux associated with a eiven frequency also rises with stellar lass together with the convective flux.," In addition, the flux associated with a given frequency also rises with stellar mass together with the convective flux."
+ We then calculate the net momentum extraction associated with the waves for an initial differeutia rotation of 6Q=0.01pz over 0.058.2...," We then calculate the net momentum extraction associated with the waves for an initial differential rotation of $\delta \Omega = 0.01 \,\mu {\rm Hz}$ over $0.05\,R_*$."
+ Figure 3 oresents the corresponding net average huuinositv at 0.038 below the surface convection zone as a function of Tag.," Figure 3 presents the corresponding net average luminosity at $0.03\,R_*$ below the surface convection zone as a function of $T_{\rm eff}$ ."
+ We see that on the ziuus the net momentum Iuiinosityv modestly increases with the stellar mass (and thus with Z;4) up to (.8AL. and then slishtlvy. decreases for moremassive, We see that on the zams the net momentum luminosity modestly increases with the stellar mass (and thus with $T_{\rm eff}$ ) up to $\sim 0.8~{\rm M}_{\odot}$ and then slightly decreases for moremassive
+and we can check that they [ie on the relation (Ίσα.,and we can check that they lie on the relation (Eq.
+ 1)., 1).
+ Επι» is à sell-consistency. check on the photometric metallicity calibration., This is a self-consistency check on the photometric metallicity calibration.
+ We decided as a simple exercise to check the self-consistency of the method ancl vielded a remarkable result: the position of a Ix dwarf within the main sequence on a colour-magnituce cliagram isvery tightly correlated. with metallicity., We decided as a simple exercise to check the self-consistency of the method and yielded a remarkable result: the position of a K dwarf within the main sequence on a colour-magnitude diagram is tightly correlated with metallicity.
+ The spectroscopic sample of 15 stars with which we made the check is shown in Table 4., The spectroscopic sample of 15 stars with which we made the check is shown in Table 4.
+ Most. of the stars come from Flynn ancl Morell. (L997). ancl a super-solar metallicity Ix cbwarf. has been taken [rom Feltzine and Gonzalez (2001).," Most of the stars come from Flynn and Morell (1997), and a super-solar metallicity K dwarf has been taken from Feltzing and Gonzalez (2001)."
+" All the stars are ""Ix. ebwarfs. ie. in the absolute magnitude range 5.5«Ady<7.3."," All the stars are “K dwarfs”, i.e. in the absolute magnitude range $5.5 < M_V < 7.3$."
+" Phe table shows the HD number. D.V. colour. absolute magnitude A, (from the Llipparcos parallax) and the dillerence between the absolute magnitude of the star and the absolute magnitude of the JLINO isochrone (at the colour of the star)."," The table shows the HD number, $B-V$ colour, absolute magnitude $M_V$ (from the Hipparcos parallax) and the difference between the absolute magnitude of the star and the absolute magnitude of the JFK0 isochrone (at the colour of the star)."
+" We then show the spectroscopically determined metallicity Fe/1l].,,. and the metallicity Fe/H]gr. derived from Eq."," We then show the spectroscopically determined metallicity $_{\mathrm spec}$ and the metallicity $_{\mathrm KF}$, derived from Eq."
+ 1., 1.
+ A colour magnitude diagram (2V. versus AA) for the stars in table 4 is shown in panel (a) of figure 15.. along with the JEIX0 isochrone.," A colour magnitude diagram $B-V$ versus $M_V$ ) for the stars in table 4 is shown in panel (a) of figure \ref{CMDkdwarfs}, along with the JFK0 isochrone."
+ Panel (b) shows the metallicities of the stars as a function. of DPV colour., Panel (b) shows the metallicities of the stars as a function of $B-V$ colour.
+ Panel. (c) shows the absolute magnitude cilference ACAA-) between the stars Ady and the absolute magnitude of the «110 isochrone (at the stars 2V. colour). as a function of the spectroscopically determined. metallicity. Fe/11]..," Panel (c) shows the absolute magnitude difference $\Delta(M_V)$ between the star's $M_V$ and the absolute magnitude of the JFK0 isochrone (at the star's $B-V$ colour), as a function of the spectroscopically determined metallicity, $_{\mathrm spec}$."
+ Phere is avery tight relationship between CMS) and Fell]. he scatter in a least squares fit is a mere 0.08 magnitudes)., There is a tight relationship between $\Delta(M_V)$ and $_{\mathrm spec}$ (the scatter in a least squares fit is a mere 0.08 magnitudes).
+If the geometry of the star-lormine environment. (he spectral distribution of stellar radiation. (he temperature distribution ancl nature of the dust. and (he gas to cust ratio were (he same for all starbursts. anv one of these parameters shoukl give the same result Lor SFR.,"If the geometry of the star-forming environment, the spectral distribution of stellar radiation, the temperature distribution and nature of the dust, and the gas to dust ratio were the same for all starbursts, any one of these parameters should give the same result for SFR."
+ Of course. all starbursts are not the same among these many characteristics. so different indicators of SER. give dillerent results depending on these characteristics.," Of course, all starbursts are not the same among these many characteristics, so different indicators of SFR give different results depending on these characteristics."
+ Although we do nol attempt here an evaluation of the relative merits of various estimates of SED. we can use the data lor our 10 Jy sample of starbursts (to estimate the dispersion (hat arises among these three different methods for measuring SER.," Although we do not attempt here an evaluation of the relative merits of various estimates of SFR, we can use the data for our 10 mJy sample of starbursts to estimate the dispersion that arises among these three different methods for measuring SFR."
+ For this comparison. we utilize the PAIL feature. the Νο feature. and the strength of the dust continuum atum.," For this comparison, we utilize the PAH feature, the [NeIII] feature, and the strength of the dust continuum at."
+. The flux ντ jm) at the peak of the PAIL feature measures the luminosity in the photodissociation region. the flux of (he Neon line measures (he ionizing Iuminosi(v within the IHE region. and the fIux density of the continmun measures the emission from warm dust.," The flux $\nu$ $_{\nu}$ $\mu$ m) at the peak of the PAH feature measures the luminosity in the photodissociation region, the flux of the Neon line measures the ionizing luminosity within the HII region, and the flux density of the continuum measures the emission from warm dust."
+ The measures used for comparison are given in Tables 2 and 3 and illusirated in Figures 12 and 13., The measures used for comparison are given in Tables 2 and 3 and illustrated in Figures 12 and 13.
+ Parameters are compared to um)) which determines if à source spectrum arises strictly [rom a starburst without AGN contamination. as discussed in section 3.1.," Parameters are compared to ) which determines if a source spectrum arises strictly from a starburst without AGN contamination, as discussed in section 3.1."
+" In Figure 12. the distribution of the ratio vf,(7.7ya) to ΠΝΟΗ} is shown for the sources in Tables 2 and 3 with PAIL features."," In Figure 12, the distribution of the ratio $\nu$ $_{\nu}$ $\mu$ m) to f([NeIII]) is shown for the sources in Tables 2 and 3 with PAH features."
+ This is a measure of huninosily arising in the PDR compared to that arising in the II] region., This is a measure of luminosity arising in the PDR compared to that arising in the HII region.
+ Including the limits which are shown. there is a weak trend in Figure 12 for PAIL EW to increase as the ratio of PALL to Νο) increases.," Including the limits which are shown, there is a weak trend in Figure 12 for PAH EW to increase as the ratio of PAH to [NeIII] increases."
+ This trend is as expected if the weak PAI] sources contain an AGN contributioΕν to the Νο luminosity., This trend is as expected if the weak PAH sources contain an AGN contribution to the [NeIII] luminosity.
+ The comparison of SFRs is intended to apply only for sourceV. which are pure starbursts. without an AGN contribution.," The comparison of SFRs is intended to apply only for sources which are pure starbursts, without an AGN contribution."
+ These pure starbursts are taken as defined in Figure 6. based on a criterion that EW(6.2jmi))) >0.," These pure starbursts are taken as defined in Figure 6, based on a criterion that ) $>$."
+4jan.. For these pure starbursts. the median and dispersion in Αλ ratio iΕν Figure 12 are log μμ τιμη)Νο) = 2.840.3.," For these pure starbursts, the median and dispersion in the PAH to [NeIII] ratio in Figure 12 are log $\nu$ $_{\nu}$ $\mu$ m)/f([NeIII])] = $\pm$ 0.3."
+ This result indicates that we should expect a dispersion by a factor of & 2.0 between SFR. estimates from PAIL compared to those from [Nell]. even for pure starbursts.," This result indicates that we should expect a dispersion by a factor of $\pm$ 2.0 between SFR estimates from PAH compared to those from [NeIII], even for pure starbursts."
+" In Figure 13. the distribution of the ratio F,(24jn) to Fr. rj) is shown for the sources in Tables 2 and 3 with PAIL features."," In Figure 13, the distribution of the ratio $_{\nu}$ $\mu$ m) to $_{\nu}$ $\mu$ m) is shown for the sources in Tables 2 and 3 with PAH features."
+ This is a measure of luminosity arising [rom warm dust within both the ILLI region and PDR compared to Iuminosity arising in the PDR., This is a measure of luminosity arising from warm dust within both the HII region and PDR compared to luminosity arising in the PDR.
+ There is a weak trend for this ratio to change with EW(6.2;0n))., There is a weak trend for this ratio to change with ).
+ This trend is expected. simply because EW(6.2j0n)) is also a measure of PAIL strength. compared to the continuum.," This trend is expected, simply because ) is also a measure of PAH strength compared to the continuum."
+ For the pure starbursts with EW(6.2;an)) >0.4jm.. the median and dispersion of the ratio are," For the pure starbursts with ) $>$, the median and dispersion of the ratio are"
+of heavy accretion may not be so dramatic.,of heavy accretion may not be so dramatic.
+ Phe models of ? show that non-accreting isochrones may only overestimate true ages for PAIS stars with Tir3500 IXIx. (about half of our sample)., The models of \citet{hosokawa11} show that non-accreting isochrones may only overestimate true ages for PMS stars with $T_{\rm eff}>3500$ K (about half of our sample).
+ ? argues that the early accretion phase probably adds significant energy to the protostar. perhaps increasing rather than decreasing the radius.," \citet{hartmann11}
+ argues that the early accretion phase probably adds significant energy to the protostar, perhaps increasing rather than decreasing the radius."
+ Nevertheless they also argue that plausible variations in initial protostellar radii and accretion histories could give rise to O.8clelex clispersions in apparent stellar. age a significant. proportion of that observed., Nevertheless they also argue that plausible variations in initial protostellar radii and accretion histories could give rise to dex dispersions in apparent stellar age – a significant proportion of that observed.
+ If our approach and assumptions are valid then ages from the LEES diagram cannot be used reliably to trace the history of star formation in the ONC as attempted. by ο and 7.., If our approach and assumptions are valid then ages from the HR diagram cannot be used reliably to trace the history of star formation in the ONC as attempted by \citet{palla99} and \citet{huff06}.
+ Furthermore. our work implies that the bulk of the stars in the ONC are formed over a timescale shorter than the median lifetime of circumstellar material.," Furthermore, our work implies that the bulk of the stars in the ONC are formed over a timescale shorter than the median lifetime of circumstellar material."
+ Our basic quantitative model suggests that if the ONC has a mean age of about MMyr. (see Fable. 1). then at least of its stars must have ages of MMsyr (based on a l-sigma dispersion. of O.ltcelex and a log normal age distribution) with a more likely range that is smaller than this.," Our basic quantitative model suggests that if the ONC has a mean age of about Myr (see Table 1), then at least of its stars must have ages of Myr (based on a 1-sigma dispersion of dex and a log normal age distribution) with a more likely range that is smaller than this."
+ Note though that this mean age. ancl hence age range. are dependent on the adopted evolutionary. models. which have significant svstematic uncertainties at these voung ages (?)..," Note though that this mean age, and hence age range, are dependent on the adopted evolutionary models, which have significant systematic uncertainties at these young ages \citep{baraffe02}."
+ The ONC age has been previously estimated at. MMwyr using PAIS isochrones and the recently revised ONC distance of &400 ppe (?7??77)..," The ONC age has been previously estimated at Myr using PMS isochrones and the recently revised ONC distance of $\simeq 400$ pc \citep{jeffries07a, sandstrom07, menten07, mayne08}."
+ DIXIO obtained an ONC age of MAlvr based the isochrones of ?.., DR10 obtained an ONC age of Myr based the isochrones of \citet{siess00}.
+ 2? provides a uniform recalibration of voung PAIS ages based. on fitting of the upper main-sequence. finding an age for the ONC of MMvr and there is a kinematic age of zc2.5 MMwr. ound by tracing back three runaway stars to their estimated origin as a single stellar svstem in the ONC (?)..," \citet{naylor09} provides a uniform recalibration of young PMS ages based on fitting of the upper main-sequence, finding an age for the ONC of Myr and there is a kinematic age of $\geq 2.5$ Myr, found by tracing back three runaway stars to their estimated origin as a single stellar system in the ONC \citep{hoogerwerf01}."
+ On the other hand a vounger mean age would agree better with he conclusions of 2.. who argued that the close similarity oetween the kinematic structure of the stars and gas in the ONC indicate that the cluster is no more than a crossing ime old (~1Myr.seealso2)..," On the other hand a younger mean age would agree better with the conclusions of \citet{tobin09}, who argued that the close similarity between the kinematic structure of the stars and gas in the ONC indicate that the cluster is no more than a crossing time old \citep[$\sim 1$\,Myr -- see
+also][]{proszkow09}."
+ A vounger age might also ος indicated by the voung Cz1 MMyr) age deduced for the ligh-mass stars in the Trapezium (???) Our analysis could »v consistent. with the adoption of any of these estimates or the mean ONC age because we only require that any age spread is smaller than the disc dispersal timescale.," A younger age might also be indicated by the young $\la 1$ Myr) age deduced for the high-mass stars in the Trapezium \citep{storzer99, clarke07, mann09} Our analysis could be consistent with the adoption of any of these estimates for the mean ONC age because we only require that any age spread is smaller than the disc dispersal timescale."
+ Whilst the ages of individual voung PAIS stars derived rom the LR diagram may be unreliable. this does. not necessarily invalidate mean ages deduced for whole clusters. »eeause the additional sources of luminosity scatter may o roughly svmametric.," Whilst the ages of individual young PMS stars derived from the HR diagram may be unreliable, this does not necessarily invalidate mean ages deduced for whole clusters, because the additional sources of luminosity scatter may be roughly symmetric."
+ Even if they were not. there is no reason to suppose that the rank ordering by mean age of well-observed. nearby clusters is incorrect and so there is no contradiction in our use of mean cluster ages to argue for a monotonic decay. of disc indicators with age (although the absolute timescale must be uncertain). whilst at the same time arguing that individual stellar ages are so uncertain as to preclude using them to estimate star formation histories.," Even if they were not, there is no reason to suppose that the rank ordering by mean age of well-observed nearby clusters is incorrect and so there is no contradiction in our use of mean cluster ages to argue for a monotonic decay of disc indicators with age (although the absolute timescale must be uncertain), whilst at the same time arguing that individual stellar ages are so uncertain as to preclude using them to estimate star formation histories."
+ 7T suggest that star formation began in the ONC at a low level as long ago as LOAIAIvr ancl indeed. there are stars in the LR diagram that are as old or older than this., \citet{huff06} suggest that star formation began in the ONC at a low level as long ago as Myr and indeed there are stars in the HR diagram that are as old or older than this.
+ However. as our mocdelling is insensitive to the tails of the age distributions (see Section 4.1). it is reasonable to ask whether a small population of older stars might still be consistent with our analysis.," However, as our modelling is insensitive to the tails of the age distributions (see Section 4.1), it is reasonable to ask whether a small population of older stars might still be consistent with our analysis."
+" But there is a problem with this idea when confronted with the infrared disc diagnostics. because too many of these ""old"" stars still possess. discs."," But there is a problem with this idea when confronted with the infrared disc diagnostics, because too many of these “old” stars still possess discs."
+ For instance. in the Spitzer sample there are 56. stars of the sample) with an apparent age ο5 MMyr and 36 of them have dises based on their 33.6/8.0] colour. a fraction consistent with the overall clise fraction ofGN%.," For instance, in the Spitzer sample there are 56 stars of the sample) with an apparent age $>5$ Myr and 36 of them have discs based on their [3.6]-[8.0] colour, a fraction consistent with the overall disc fraction of."
+". If the ages of these ""old stars were accurate. then for a reasonable exponential disc lifetime of sav MMyr. (sce Section 4.1). we would only expect to see. 13 stars with discs — inconsistent. with the observed. value at very high gaienificance."," If the ages of these “old” stars were accurate, then for a reasonable exponential disc lifetime of say Myr (see Section 4.1), we would only expect to see 13 stars with discs – inconsistent with the observed value at very high significance."
+ Some of the more extreme objects (in terms of their apparent age) could be examples of stars occulted by their disces and observed in scattered light (2).., Some of the more extreme objects (in terms of their apparent age) could be examples of stars occulted by their discs and observed in scattered light \citep{slesnick04}.
+ 7. find tha 10 components of (presumably coeval) PALS binary systenis frequently exhibit apparent age dilferences of OA dex. anc more. which they also attribute to systematic problems in estimating the luminosities of PAIS stars with discs anc accretion.," \citet{kraus09} find that the components of (presumably coeval) PMS binary systems frequently exhibit apparent age differences of 0.4 dex and more, which they also attribute to systematic problems in estimating the luminosities of PMS stars with discs and accretion."
+ A similar argument applies to the 56 stars in the Spitzer sample that are apparently. vounger than MMwyr., A similar argument applies to the 56 stars in the Spitzer sample that are apparently younger than Myr.
+ Only 32 of these possess discs compared to the 51 expectec for a clise lifetime of Myr and an initial clise fraction of unity., Only 32 of these possess discs compared to the 51 expected for a disc lifetime of Myr and an initial disc fraction of unity.
+ In summary. the disc frequencies also argue agains the accuracy of the ages of objects in both. the voung ane old tails of the age distribution and. like the bulk of the population. their luminosities must too be explained in terms of observational uncertanties or physical ellects tha alter the Iuminositv-age relationship.," In summary, the disc frequencies also argue against the accuracy of the ages of objects in both the young and old tails of the age distribution and, like the bulk of the population, their luminosities must too be explained in terms of observational uncertanties or physical effects that alter the luminosity-age relationship."
+ Assuming that most stars are born with circumstellar material and that the infrared signatures of this material decay. on average. monotonically with time. then a wide spread of ages (210 MMsr) in the ONC should: manifest itself by marked. differences in the age distributions of stars with and without infrared. disc signatures.," Assuming that most stars are born with circumstellar material and that the infrared signatures of this material decay, on average, monotonically with time, then a wide spread of ages $\simeq 10$ Myr) in the ONC should manifest itself by marked differences in the age distributions of stars with and without infrared disc signatures."
+ “Vhis hypothesis has been tested using a large. homogeneous ssumple from the ONC catalogue of DILIO. using three independent. means of diagnosing disc presence and ages derived. from. the UR. diagram.," This hypothesis has been tested using a large, homogeneous sample from the ONC catalogue of DR10, using three independent means of diagnosing disc presence and ages derived from the HR diagram."
+ We have found no significant evidence for differences in the apparent age distributions of stars with and without discs and their means and medians are very similar., We have found no significant evidence for differences in the apparent age distributions of stars with and without discs and their means and medians are very similar.
+ This is consistent with the conclusion that any real age spread in the ONC must be smaller than the median lifetime of the cireumstellar disces., This is consistent with the conclusion that any real age spread in the ONC must be smaller than the median lifetime of the circumstellar discs.
+ A simple quantitative model has been developed to interpret these results., A simple quantitative model has been developed to interpret these results.
+ This model suggests that. for plausible disc lifetimes. the contribution of any real age spread to the apparent age dispersion. inferred. from the UR diagram must indeed be very small «0.14 cdex ispersion in a log-normal age distribution at 99 per cent confidence. compared. with an observed. age dispersion. of c- (klddex.," This model suggests that for plausible disc lifetimes, the contribution of any real age spread to the apparent age dispersion inferred from the HR diagram must indeed be very small; $<0.14$ dex dispersion in a log-normal age distribution at 99 per cent confidence, compared with an observed age dispersion of $\simeq 0.4$ dex."
+ Even stars in the tails of the apparent age istribution have disc frequencies incompatible with their ipparent ages and consistent with coevality with the rest of 10 population., Even stars in the tails of the apparent age distribution have disc frequencies incompatible with their apparent ages and consistent with coevality with the rest of the population.
+ These results argue strongly against cluster, These results argue strongly against cluster
+in dynamicallyP voung clusters (see. e.g. Schindler 1999 for Virgo: Machacek 2002 and AbclelSalam 1998 for Abell 2218: Drinkwater 2001 for Fornax).,"in dynamically young clusters (see, e.g. Schindler 1999 for Virgo; Machacek 2002 and AbdelSalam 1998 for Abell 2218; Drinkwater 2001 for Fornax)."
+ We adopt simple isothermal beta models for the surface brightuess and gas deusity of the Foruax cluster gas outside he edge such that the surface brightness S(r) is given by where S(r) is the surface brightuess at a giveu radius. Sq is the central surface brightuess. aud i. is the core radius.," We adopt simple isothermal beta models for the surface brightness and gas density of the Fornax cluster gas outside the edge such that the surface brightness $S(r)$ is given by where $S(r)$ is the surface brightness at a given radius, $S_0$ is the central surface brightness, and $r_c$ is the core radius."
+ Then the correspouding electron deusity is given by Iu Figure 6 we show two multiple-component beta moclel fits to the surface brightness prolile ol NGC 1399. with point sources aud emission [rom NCC L101 excluded. where the fit to the data is the stun of individual components given by Equation 2..," Then the corresponding electron density is given by In Figure \ref{fig:n1399rprof} we show two multiple-component beta model fits to the surface brightness profile of NGC 1399, with point sources and emission from NGC 1404 excluded, where the fit to the data is the sum of individual components given by Equation \ref{eq:betamodel}."
+ Iu each case the tuuermost central regiou, In each case the innermost central region
+"depend on the details of the subgrid models; the dynamics of the marginally balanced cavity shells are strongly influenced by the stellar evolution model and the dust model, particularly the treatment of sublimation and evaporation.","depend on the details of the subgrid models; the dynamics of the marginally balanced cavity shells are strongly influenced by the stellar evolution model and the dust model, particularly the treatment of sublimation and evaporation."
+" In contrast to this epoch of marginal Eddington equilibrium, the extent of the outflow cavity in the RT + FLD simulations increases rapidly in time."," In contrast to this epoch of marginal Eddington equilibrium, the extent of the outflow cavity in the RT + FLD simulations increases rapidly in time."
+" With the exception of the More=50Mo case, in which a correspondingly lower mass star (yielding much lower luminosity) is formed, the outflow cavity increases in size monotonically in the RT + FLD simulations."," With the exception of the $M_\mathrm{core}=50\Msol$ case, in which a correspondingly lower mass star (yielding much lower luminosity) is formed, the outflow cavity increases in size monotonically in the RT + FLD simulations."
+ The difference in the growth rate for the various radiation transport schemes is most prominent in the More=100Mo and pοςr case (Fig., The difference in the growth rate for the various radiation transport schemes is most prominent in the $M_\mathrm{core}=100\Msol$ and $\rho \propto r^{-2}$ case (Fig.
+" 4 left panel), i.e. in the case of the initially highest mass in the central core region, allowing for a rapid formation and growth of the central star."," \ref{fig:OutflowRadiusvsTime} left panel), i.e. in the case of the initially highest mass in the central core region, allowing for a rapid formation and growth of the central star."
+ In the Moore=50Μο case (Fig., In the $M_\mathrm{core}=50\Msol$ case (Fig.
+" 4 middle panel), the much lower stellar luminosity results in an initially unstable cavity region even in the RT+FLD case (but with a smaller extent than 100 AU)."," \ref{fig:OutflowRadiusvsTime} middle panel), the much lower stellar luminosity results in an initially unstable cavity region even in the RT+FLD case (but with a smaller extent than 100 AU)."
+ In the p«r7!? case (Fig., In the $\rho \propto r^{-1.5}$ case (Fig.
+" 4 right panel), the mass is initially more concentrated at larger radii and hence the cavity shell has to sweep up far more mass during its increase."," \ref{fig:OutflowRadiusvsTime} right panel), the mass is initially more concentrated at larger radii and hence the cavity shell has to sweep up far more mass during its increase."
+" Therefore, the growth phase of the cavity takes much longer than in the cases of initially steeper density profiles."," Therefore, the growth phase of the cavity takes much longer than in the cases of initially steeper density profiles."
+" To unveil the physical background of the qualitative and quantitative difference of the outflow cavity structure in the simulations using either the RT+FLD or the FLD method, we now investigate the morphology and dynamics of the cavity shell depending on the radiation transport method."," To unveil the physical background of the qualitative and quantitative difference of the outflow cavity structure in the simulations using either the RT+FLD or the FLD method, we now investigate the morphology and dynamics of the cavity shell depending on the radiation transport method."
+" In this section, we analyze quantitatively the difference in the morphology and the dynamics of the cavity shell in the RT+FLD and FLD cases."," In this section, we analyze quantitatively the difference in the morphology and the dynamics of the cavity shell in the RT+FLD and FLD cases."
+ We focus on the data of the simulation with an initial core mass Meore=100Mo and a radial density slope of B=—2., We focus on the data of the simulation with an initial core mass $M_\mathrm{core} = 100\Msol$ and a radial density slope of $\beta=-2$.
+ The cavity in the FLD case increases in its first growth phase up to roughly 1400 AU before the expansion along the polar axis stops and the mass flow is reversed by gravity., The cavity in the FLD case increases in its first growth phase up to roughly 1400 AU before the expansion along the polar axis stops and the mass flow is reversed by gravity.
+" We analyze the gas dynamics at the point in time when in both simulations (with RT+FLD and with FLD only) the expansion of the outflow cavity has arrived at the same location, namely at t=16.7 kyr."," We analyze the gas dynamics at the point in time when in both simulations (with RT+FLD and with FLD only) the expansion of the outflow cavity has arrived at the same location, namely at $t=16.7$ kyr."
+" The gas density, temperature, radial velocity, and the radial mass loss rate along the polar axis at this time are shown in Fig. 5.."," The gas density, temperature, radial velocity, and the radial mass loss rate along the polar axis at this time are shown in Fig. \ref{fig:1Danalysis}."
+" The upper left panel of this figure shows that there is almost no difference in the radial extent, compactness, and peak density of the swept-up mass in the cavity shell."," The upper left panel of this figure shows that there is almost no difference in the radial extent, compactness, and peak density of the swept-up mass in the cavity shell."
+" The peaks in the density distributions correspond to 4.1x107!”gcm? and 3.7x107!7gcm? in the RT+FLD and FLD cases, respectively."," The peaks in the density distributions correspond to $4.1\times10^{-17}\rhocgs$ and $3.7\times10^{-17}\rhocgs$ in the RT+FLD and FLD cases, respectively."
+ The upper right panel of Fig., The upper right panel of Fig.
+" 5 shows that the RT+FLD run results in a continuously slightly higher temperature in the cleared cavity, the cavity shell, and beyond."," \ref{fig:1Danalysis} shows that the RT+FLD run results in a continuously slightly higher temperature in the cleared cavity, the cavity shell, and beyond."
+ The temperature distribution in the pure FLD run stays to below the temperature distribution of the RT--FLD run., The temperature distribution in the pure FLD run stays to below the temperature distribution of the RT+FLD run.
+" In contrast to these similarities, the two lower panels of Fig."," In contrast to these similarities, the two lower panels of Fig."
+ 5 highlight the striking difference of both runs in their radial velocities and mass flux along the polar axis., \ref{fig:1Danalysis} highlight the striking difference of both runs in their radial velocities and mass flux along the polar axis.
+" In the RT+FLD run, the swept-up shell material moves into the interstellar medium with a speed slightly higher than 100 km s!, which is roughly a factor of three higher than in the corresponding FLD run."," In the RT+FLD run, the swept-up shell material moves into the interstellar medium with a speed slightly higher than 100 km $\mbox{s}^{-1}$ , which is roughly a factor of three higher than in the corresponding FLD run."
+" At larger radii (r>3000 AU), the velocity slope is still dominated by gravitational infall and hence independent of the radiation transport method."," At larger radii $r > 3000$ AU), the velocity slope is still dominated by gravitational infall and hence independent of the radiation transport method."
+" Up to the onset of the Eddington equilibrium epoch in the FLD run, the mass of the central star in both simulations (RT+FLD and FLD only) closely match, leading to the same strength of gravitational attraction for the matter in the cavity and shell towards the direction of the star."," Up to the onset of the Eddington equilibrium epoch in the FLD run, the mass of the central star in both simulations (RT+FLD and FLD only) closely match, leading to the same strength of gravitational attraction for the matter in the cavity and shell towards the direction of the star."
+" Furthermore, the density structure is roughly the same (Fig. 5,,"," Furthermore, the density structure is roughly the same (Fig. \ref{fig:1Danalysis},"
+" upper left panel) and the deviation in temperature within the cavity is smaller than (Fig. 5,,"," upper left panel) and the deviation in temperature within the cavity is smaller than (Fig. \ref{fig:1Danalysis},"
+" upper right panel), hence the slightly higher thermal pressure cannot be the main driver of the enormous difference in the radial velocities and mass flux (Fig. 5,,"," upper right panel), hence the slightly higher thermal pressure cannot be the main driver of the enormous difference in the radial velocities and mass flux (Fig. \ref{fig:1Danalysis},"
+ lower panels)., lower panels).
+" As a consequence, the radiation pressure acting on the swept-up mass on top of the outflow cavity has to differ in the two radiation transport methods."," As a consequence, the radiation pressure acting on the swept-up mass on top of the outflow cavity has to differ in the two radiation transport methods."
+" First, the FLD approximation assumes that the dust grains of the swept-up mass on top of the cavity are embedded in a local radiation field, whereas the RT+FLD method takes into account the (frequency-dependent) absorption probability caused by direct stellar irradiation."," First, the FLD approximation assumes that the dust grains of the swept-up mass on top of the cavity are embedded in a local radiation field, whereas the RT+FLD method takes into account the (frequency-dependent) absorption probability caused by direct stellar irradiation."
+ The resulting difference in the radiative acceleration at the top of the outflow cavity is analytically estimated in the following section., The resulting difference in the radiative acceleration at the top of the outflow cavity is analytically estimated in the following section.
+" Secondly, since the FLD approximation is mathematically a moment method, the derivation of the FLD approximation includes the integral over the angular distribution of the radiative flux; hence the emitted photons of the central star do not move along straight rays until they are absorbed."," Secondly, since the FLD approximation is mathematically a moment method, the derivation of the FLD approximation includes the integral over the angular distribution of the radiative flux; hence the emitted photons of the central star do not move along straight rays until they are absorbed."
+ The radiative flux in the FLD approximation instead follows a path that minimizes the optical depth., The radiative flux in the FLD approximation instead follows a path that minimizes the optical depth.
+" The FLD method was originally introduced to describe spherically symmetric flows, hence the integral over the propagation direction did not result in unphysical behavior (seee.g.?).."," The FLD method was originally introduced to describe spherically symmetric flows, hence the integral over the propagation direction did not result in unphysical behavior \citep[see e.g.][]{Bruenn:1985p18798}."
+" In a multi-dimensional environment with a non-isotropic optical depth, this assumption of the FLD approximation breaks down."," In a multi-dimensional environment with a non-isotropic optical depth, this assumption of the FLD approximation breaks down."
+ Including higher-order terms in the derivation would minimize this inaccuracy and potentially yield a sufficient tracing of the correct photon path., Including higher-order terms in the derivation would minimize this inaccuracy and potentially yield a sufficient tracing of the correct photon path.
+" However, in the RT+FLD scheme the irradiation by the central star is computed via a tay-tracing equation, which takes into account the correct photon propagation direction, i.e. the stellar irradiationflux directly impinges on the swept-up mass shell at the top of the optically thin cavity."," However, in the RT+FLD scheme the irradiation by the central star is computed via a tay-tracing equation, which takes into account the correct photon propagation direction, i.e. the stellar irradiationflux directly impinges on the swept-up mass shell at the top of the optically thin cavity."
+the earlier models ??).,the earlier models .
+. Note that this assumption can be easily dropped in all modern line models., Note that this assumption can be easily dropped in all modern line models.
+" In order to allow for a comparison of the line shapes with earlier results, we use the law of?,, even though for lines caused by fluorescence due to the irradiation of a disc with hard X-rays from above, a limb-brightening law would be more appropriate(?)."," In order to allow for a comparison of the line shapes with earlier results, we use the limb-darkening law of, even though for lines caused by fluorescence due to the irradiation of a disc with hard X-rays from above, a limb-brightening law would be more appropriate."
+". In order to illustrate the pure frame-dragging effects of these different spins onto the line shape, the inner disc radius was set to 9rg for all three spins."," In order to illustrate the pure frame-dragging effects of these different spins onto the line shape, the inner disc radius was set to $9\,r_\mathrm{g}$ for all three spins."
+" The Figure shows that in this case the major difference between the different spins is the relative strength of the core of the line to the red wing, which decreases with decreasing a."," The Figure shows that in this case the major difference between the different spins is the relative strength of the core of the line to the red wing, which decreases with decreasing $a$."
+" For this case of a large inner radius, the most significant differences in line shape are seen for low values of 0, while the red tails are virtually indistinguishable."," For this case of a large inner radius, the most significant differences in line shape are seen for low values of $\theta_\mathrm{o}$ while the red tails are virtually indistinguishable."
+" The slight increase in line flux at the lowest energies is due to the increased Doppler boosting in the case of a«0 (for a given radius, u* increases with decreasing a, cf."," The slight increase in line flux at the lowest energies is due to the increased Doppler boosting in the case of $a<0$ (for a given radius, $u^t$ increases with decreasing $a$, cf."
+ Eq. 12))., Eq. \ref{eq:four_velocity}) ).
+ The difference in energy shift of photons emerging from an accretion disc between maximal positive and maximal negative spin of the black hole is shown in Fig. 4.., The difference in energy shift of photons emerging from an accretion disc between maximal positive and maximal negative spin of the black hole is shown in Fig. \ref{fig:diff_image}.
+" As these differences are highest close to inner edge of the disc, a higher emissivity pronounces the devations in the line profiles."," As these differences are highest close to inner edge of the disc, a higher emissivity pronounces the devations in the line profiles."
+ Moreover this figure shows that for an accretion disc with an inner radius larger than 30rz no significant differences would be expected.," Moreover this figure shows that for an accretion disc with an inner radius larger than $30\,r_\mathrm{g}$ no significant differences would be expected."
+ Figure 5 shows line profiles for different spins of the black hole for the more realistic case that the disc extends down to the marginally stable orbit., Figure \ref{fig:line_risco} shows line profiles for different spins of the black hole for the more realistic case that the disc extends down to the marginally stable orbit.
+" Since the inner edge of the disc is closer to the black hole for positively spinning black holes, more strongly redshifted photons emerge."," Since the inner edge of the disc is closer to the black hole for positively spinning black holes, more strongly redshifted photons emerge."
+" As already noted by?,, this leads to broader lines in these systems, especially for discs with an emissivity that is strongly peaked towards rin."," As already noted by, this leads to broader lines in these systems, especially for discs with an emissivity that is strongly peaked towards $r_\mathrm{in}$."
+" Maximally negatively spinning black holes have the smallest width, although the line will still be detectable as being broad even at CCD resolution (depending on inclination, typical widths of the main peak are around eeV)."," Maximally negatively spinning black holes have the smallest width, although the line will still be detectable as being broad even at CCD resolution (depending on inclination, typical widths of the main peak are around eV)."
+ Lines from counterrotating black holes will therefore be more difficult to detect than lines from positively rotating black holes., Lines from counterrotating black holes will therefore be more difficult to detect than lines from positively rotating black holes.
+ The major difference of line shapes for discs around black holes with a=0 and counterrotating discs (see Fig. 5)), The major difference of line shapes for discs around black holes with $a=0$ and counterrotating discs (see Fig. \ref{fig:line_risco}) )
+" lies in the strength of the blue peak, since the skew symmetric shape is mainly due to frame dragging effects and the small inner radii."," lies in the strength of the blue peak, since the skew symmetric shape is mainly due to frame dragging effects and the small inner radii."
+ Detecting these lines observationally is therefore more difficult than detecting lines from discs around positively rotating black hole., Detecting these lines observationally is therefore more difficult than detecting lines from discs around a positively rotating black hole.
+" In addition, as shown by limb-darkening/-brighteninga affects the strength of the red wing."," In addition, as shown by limb-darkening/-brightening affects the strength of the red wing."
+" For counterrotating black holes, this results in a possible degeneracy, as for different limb-darkening"," For counterrotating black holes, this results in a possible degeneracy, as for different limb-darkening"
+"practice, these quantities are much more tightly correlated with Note that these are degenerate with solutions for which £ differs by 180°, i.e., €=1289*17735., and for which the line-of- of the is reversed, i.e., 0—. sightWe do not comp","practice, these quantities are much more tightly correlated with Note that these are degenerate with solutions for which $\xi$ differs by $180^\circ$ , i.e., $\xi={128^\circ}^{+17^\circ+33^\circ}_{-15^\circ-24^\circ}$, and for which the line-of-sight component of the spin is reversed, i.e., $\theta={112^\circ}_{-5^\circ-9^\circ}^{+20^\circ+28^\circ}$."
+"onentattempt to spindetermine the systematic 1129320728""uncertainties associated with a accretion model.", We do not attempt to determine the systematic uncertainties associated with selecting a particular accretion model.
+" However, we note that a selectingnumber of particularefforts to fit alternative accretion flow models to the 2007 have reached consistent results despite differences in the epochmodels, suggesting that these results are robust."," However, we note that a number of efforts to fit alternative accretion flow models to the 2007 epoch have reached consistent results despite differences in the models, suggesting that these results are robust."
+" Furthermore, the of the fits to the mm-VLBI visibility and spectral data, qualityconcurrently, that the features of the BL06 model responsible for determiningsuggests the spectral and image properties are generic, and are therefore insensitive to the particulars of the accretion flow models."," Furthermore, the quality of the fits to the mm-VLBI visibility and spectral data, concurrently, suggests that the features of the BL06 model responsible for determining the spectral and image properties are generic, and are therefore insensitive to the particulars of the accretion flow models."
+" However, full studies of the systematic errors associated with the particular choices made for the accretion flow properties and the underlying spacetimes are now justified."," However, full studies of the systematic errors associated with the particular choices made for the accretion flow properties and the underlying spacetimes are now justified."
+" While beyond the purview of the present paper, we will such efforts elsewhere."," While beyond the purview of the present paper, we will report upon such efforts elsewhere."
+" Probabilityreport upondistributions for each of the spin parameters, marginalized over all others, are shown in Figure 9.."," Probability distributions for each of the spin parameters, marginalized over all others, are shown in Figure \ref{fig:marg1D}."
+" In addition to the present case, these are also shown for the analysis of the 2007 for comparison."," In addition to the present case, these are also shown for the analysis of the 2007 epoch for comparison."
+" In these the lo and 20 ranges, defined by epochthe cumulative probability, are also shown."," In these the $1\sigma$ and $2\sigma$ ranges, defined by the cumulative probability, are also shown."
+ In all cases the marginalized distributions from the combined data set are much more probabilitynarrowly peaked than their 2007 counterparts., In all cases the marginalized probability distributions from the combined data set are much more narrowly peaked than their 2007 epoch counterparts.
+" Nevertheless, are all consistent at the 1c epochlevel with those obtained from the they2007 epoch alone."," Nevertheless, they are all consistent at the $1\sigma$ level with those obtained from the 2007 epoch alone."
+" It is now possible to exclude a>0.62 at the 2c level, with a=0130321043 substantially preferring non-spinning models."," It is now possible to exclude $a>0.62$ at the $2\sigma$ level, with $a=0^{+0.32+0.62}$, substantially preferring non-spinning models."
+ Thus the high-spin island seen in Figure 7 of ? is now eliminated., Thus the high-spin island seen in Figure 7 of \citet{Brod_etal:09} is now eliminated.
+" Similarly, the position angle is now very clearly constrained, choosing the solution less favored by the 2007 epoch data (though still consistent at the 1c level)."," Similarly, the position angle is now very clearly constrained, choosing the solution less favored by the 2007 epoch data (though still consistent at the $1\sigma$ level)."
+ In this case we have £=-609*172?., In this case we have $\xi={-60^\circ}^{+14^\circ+29^\circ}_{-5^\circ-12^\circ}$.
+" Finally, the most probable viewing angle is 0=61510.2."," Finally, the most probable viewing angle is $\theta={61^\circ}^{+7^\circ+24^\circ}_{-9^\circ-15^\circ}$."
+" This is somewhat higher thanthe most likely value from the 2007 epoch alone, though well within the 1c uncertainty."," This is somewhat higher thanthe most likely value from the 2007 epoch alone, though well within the $1\sigma$ uncertainty."
+ratio should shehltly decrease with a larger Z7 due to observational bias effect aud base on model prediction (see Section 2 of their paper) which is opposed to the iuplication of Table &..,"ratio should slightly decrease with a larger $H$ due to observational bias effect and base on model prediction (see Section 2 of their paper), which is opposed to the implication of Table \ref{tbl-8}."
+ To investigate this matter further. we merge all data together. creating a large dataset with 110 NEAs (with 131 NEAs with 77<22 and 382 NEAs with J< 20). and erouped them into onc-Inaguitude-wide bius.," To investigate this matter further, we merge all data together, creating a large dataset with 449 NEAs (with 434 NEAs with $H\leq22$ and 382 NEAs with $H\leq20$ ), and grouped them into one-magnitude-wide bins."
+ For cach magnitude bin. the C/X-like:S-like ratio is computed base on the scheme described above.," For each magnitude bin, the C/X-like:S-like ratio is computed base on the scheme described above."
+ The result is listed iu Table 9.., The result is listed in Table \ref{tbl-9}.
+ Comparing the result with Morbidelli ot αἱ., Comparing the result with Morbidelli et al.
+ study. which has included 183 NEAs with JI<20. he ratio variation from magnitude to magnitude votween the two studies is similar: a lower C/X- ratio on {1«18 aud a higher ratio ou IT>leeD< Usan. which agrees the iuplication of Table 8..,"'s study, which has included 183 NEAs with $H<20$, the ratio variation from magnitude to magnitude between the two studies is similar: a lower C/X-like:S-like ratio on $H<18$ and a higher ratio on $H\geq18\Leftrightarrow D\leq1$ km, which agrees the implication of Table \ref{tbl-8}."
+ To check the statistical significance of his phenomenon. we divide the whole sample iuto wo groups by an 77=18 cutoff. with the sample n of each group to be 216 UF< I8) aud 203 (LT> 18) respectively. aud perform a A? test ou the two eroups.," To check the statistical significance of this phenomenon, we divide the whole sample into two groups by an $H=18$ cutoff, with the sample $n$ of each group to be 246 $H<18$ ) and 203 $H\geq18$ ) respectively, and perform a $\chi^{2}$ test on the two groups."
+ It has been shown that the distributions of two eroups are clifferent at a confidence level of99.5%... which is large enough to be considered as statistically significant.," It has been shown that the distributions of two groups are different at a confidence level of, which is large enough to be considered as statistically significant."
+" Although our treatment of the C/AN-like:S-like ratio of the N-complex may induct some wncertaity., we note that even we consider a lo uncertaünutv of the nunubers of X-conrplex menibers to be ποκακο (assunungi random observation errors) for both ff<18 aud II>Us. for which the nuubers of X-conipkX members cousidered as ""C/X-like would be 104:H) and 17cod respectively, the minim possible confidence level is still shown to be ~95."," Although our treatment of the C/X-like:S-like ratio of the X-complex may induct some uncertainty, we note that even we consider a $1\sigma$ uncertainty of the numbers of X-complex members to be “C/X-like” (assuming random observation errors) for both $H<18$ and $H\geq18$, for which the numbers of X-complex members considered as “C/X-like” would be $10\pm3$ and $17\pm4$ respectively, the minimum possible confidence level is still shown to be $\sim95$."
+ What Is nore. we notice that the C/X-like:S-like ratio we derived is 0.17+0.02 for NEAs with JI«Is. which is iu a good match ofthe model prediction of 0.15 according to Morbidelli et a. iplviug our treatinenut of he N-complex members with 7«18 would lead o a good match with the model.," What is more, we notice that the C/X-like:S-like ratio we derived is $0.17\pm0.02$ for NEAs with $H<18$, which is in a good match of the model prediction of 0.18 according to Morbidelli et al, implying our treatment of the X-complex members with $H<18$ would lead to a good match with the model."
+ However. thiugs are no lonecr that case when it comes to an Jf<20 cutoff line. for which our estimate of the C/N-like:S-like ratio is show to be 0.22£0.02. about ligher than Morbidelli et al," However, things are no longer that case when it comes to an $H<20$ cutoff line, for which our estimate of the C/X-like:S-like ratio is shown to be $0.22\pm0.02$, about higher than Morbidelli et al."
+s model prediction (0.17).,'s model prediction (0.17).
+" It is noteworthy that we have to cousider aff N-comiplex mecuibers with ls«JJ20 to be ""S-like as in an extreme situation for a compatible ratio with the nodel prediction (0.18 versus the models 0.1654E0.015).", It is noteworthy that we have to consider $all$ X-complex members with $18\leq H<20$ to be “S-like” as in an extreme situation for a compatible ratio with the model prediction (0.18 versus the model's $0.165\pm0.015$ ).
+ lu addition. our estimate of C/X-like:S-like ratio for lsΠ«22 (sample 50=1l8T comparing with η~50 for ls<If20 in Morbidelli et al," In addition, our estimate of C/X-like:S-like ratio for $18\leq H<22$ (sample $n=187$ comparing with $n\sim50$ for $18\leq H<20$ in Morbidelli et al."
+os sample) is 0.33£0.01.,'s sample) is $0.33\pm0.04$.
+ Iu all. it is plausible to sav that a transition point of C/N-like:S-like ratio exists at JF=18 in our sample. which is consistent with the hypothesis that small asteroids eoncrally have vouuger collision age aud teud to be ess weathered (thus. less reddish/ueutral colored) han large ones (Binzeletal.1998).. hence leads to he overabundance of C/X-like asteroids in smal NEAs.," In all, it is plausible to say that a transition point of C/X-like:S-like ratio exists at $H=18$ in our sample, which is consistent with the hypothesis that small asteroids generally have younger collision age and tend to be less weathered (thus, less reddish/neutral colored) than large ones \citep{bin98}, hence leads to the overabundance of C/X-like asteroids in small NEAs."
+ ALorbicJl: et al., Morbidelli et al.
+ suggested a factor of —1. vetween observed aud true C/X-like:S-like ratio caused by the phase angle effect at discovery., suggested a factor of $\sim1.4$ between observed and true C/X-like:S-like ratio caused by the phase angle effect at discovery.
+ Asstuuing that this nunuber remains quasi-cousta or differeut. £7. we will eet a true C/N-like:S-like ratio of 0.21 for asteroids with [7<Ls. which is in a good match with Morbidelli et al," Assuming that this number remains quasi-constant for different $H$, we will get a true C/X-like:S-like ratio of 0.24 for asteroids with $H<18$, which is in a good match with Morbidelli et al."
+ mode (0.2540.02).,'s model $0.25\pm0.02$ ).
+ However. for ff>18. the true ratio can be as high as 0.5 in our suuple.," However, for $H\geq18$, the true ratio can be as high as 0.5 in our sample."
+ Because of the coustraiuts of the target observability. only two pairs were with both components observed. they are (L979) Sakharov vs. (13732) Woodall and (11812) Kapbos vs. (228717) 2002 VID respectively.," Because of the constraints of the target observability, only two pairs were with both components observed, they are (1979) Sakharov vs. (13732) Woodall and (11842) Kap'bos vs. (228747) 2002 VH3 respectively."
+ Unfortunatelv. we were unable to obtain complete color indices for either pairs: for he first pair. (13732) Woodall is suffered. from instrument problems so the Wo7 imnaguitude is inissine. while (228717) 2002 VIT3 was just ossine near a 15 mae star when the observation was taken for the second pair. so we are only able to make a very crude classification for these wo objects.," Unfortunately, we were unable to obtain complete color indices for either pairs: for the first pair, (13732) Woodall is suffered from instrument problems so the $V-I$ magnitude is missing, while (228747) 2002 VH3 was just passing near a 4.5 mag star when the observation was taken for the second pair, so we are only able to make a very crude classification for these two objects."
+ In general. both components of both ours are classified as S-like asteroid. providing a weak evidence of the common-origim bypothesis as expected.," In general, both components of both pairs are classified as S-like asteroid, providing a weak evidence of the common-origin hypothesis as expected."
+ Our efforts have added the spectral data of sole tenus of interesting asteroids into knowledec., Our efforts have added the spectral data of some tens of interesting asteroids into knowledge.
+ For the few asteroids that had been reportedly classified. our results are ecnerally in rough agreement. indicating the goodness of our work.," For the few asteroids that had been reportedly classified, our results are generally in rough agreement, indicating the goodness of our work."
+ Further to observingC» aud classifviugOo the asteroids.," Further to observing and classifying the asteroids,"
+"intensity (J) data are required to derive the field strength, we applied an appropriate flat-field correction to all spectra.","intensity $I$ ) data are required to derive the field strength, we applied an appropriate flat-field correction to all spectra."
+" As discussedcbs by ?,, under the weak field approximation one can derive the mean longitudinal magnetic field component over the stellar disk (B) from the circular polarimetry and the gradient of the intensity spectrum with the equation: Here we used a factor Cz summarizing the physical constants like the electron charge e, the electron mass m. and the speed of the light c in the form: and an effective Landé factor gog=1 as discussed by ?.."," As discussed by \citet{2002A&A...389..191B}, under the weak field approximation one can derive the mean longitudinal magnetic field component over the stellar disk $\langle B_z \rangle$ from the circular polarimetry and the gradient of the intensity spectrum with the equation: Here we used a factor ${\rm C}_z$ summarizing the physical constants like the electron charge ${\rm e}$, the electron mass ${\rm m_e}$ and the speed of the light ${\rm c}$ in the form: and an effective Landé factor ${\rm g}_{\rm eff} = 1$ as discussed by \citet{1994A&A...291..668C}."
+ It should be clearly mentioned that there are a number of approximations and assumptions in equation 1 above., It should be clearly mentioned that there are a number of approximations and assumptions in equation \ref{eqn_reg} above.
+ One of, One of
+recover the obscured values.,recover the obscured values.
+ Finally. we describe the sSFR as a fiction of local galaxy density. fiucineg a higher fraction of star forming galaxies at the cluster outskirts. than witlin the cluster core.," Finally, we describe the sSFR as a function of local galaxy density, finding a higher fraction of star forming galaxies at the cluster outskirts, than within the cluster core."
+ Figure 9 shows the galaxv optical color aud stellar mass. as a function of SER.," Figure \ref{sfrcol} shows the galaxy optical color and stellar mass, as a function of sSFR."
+ The starburst ealaxies are blue. with (ea )« 0.7.aud are dawarts. with << LOAD...," The starburst galaxies are blue, with $^{\prime}$ $^{\prime}$ $<$ 0.7,and are dwarfs, with $\,$$<$$\,$ $^{9}$ $_{\odot}$."
+ Most sources with a specific star formation rate less than 0.005 are red and have a large stellar mass., Most sources with a specific star formation rate less than 0.005 are red and have a large stellar mass.
+ These are the red-sequence galaxies that dominate the cluster core and the FIR enüssiou is likely caused by a large population of old red giant branch stars (also known to cuit at MIPS μιας ??)).," These are the red-sequence galaxies that dominate the cluster core and the FIR emission is likely caused by a large population of old red giant branch stars (also known to emit at MIPS $\mu$ m; \citet{tem08,cal10}) )."
+ Oue weht expect that if a large population of dusty dwarfs existed. some low mass galaxies should be red by virtue of their dust extinction. rather than because of age.," One might expect that if a large population of dusty dwarfs existed, some low mass galaxies should be red by virtue of their dust extinction, rather than because of age."
+ There are a few red dwarts with normal star formation. but we do not detect any starbursting red clwarts.," There are a few red dwarfs with normal star formation, but we do not detect any starbursting red dwarfs."
+ Because dust-obseured starburst galaxies should be optically fainter than their unobscured. bhie. counterparts. if may be that our spectroscopic data is too shallow or incomplete to detect these fainter dwarfs.," Because dust-obscured starburst galaxies should be optically fainter than their unobscured, blue, counterparts, it may be that our spectroscopic data is too shallow or incomplete to detect these fainter dwarfs."
+ To investigate this issuc. we preseut FieureOo LO which malses use of a deeper. complete catalog of sources based on photometric redshifts.," To investigate this issue, we present Figure \ref{grcolmag} which makes use of a deeper, complete catalog of sources based on photometric redshifts."
+ The figure plots the color magnitude diagram of MIPS detected members., The figure plots the color magnitude diagram of MIPS detected members.
+ The spectroscopic nenibers (filled black circles) are compared to he photometric members (open circles: defined as hose with ρωσ 0.08).," The spectroscopic members (filled black circles) are compared to the photometric members (open circles; defined as those with $_{phot} <\,$ 0.08)."
+ If it was true that our spectroscopic limit introduced a bias against dusty red star forming warts. than the photometric nembers. which are much fainter. would show a sample of low-inass rec ealaxies.," If it was true that our spectroscopic limit introduced a bias against dusty red star forming dwarfs, than the photometric members, which are much fainter, would show a sample of low-mass red galaxies."
+" But. as the feure shows. there does not appear to be a missing opulation of red 21,2440 sources just bevoud our spectroscopic limit of 2—19.5."," But, as the figure shows, there does not appear to be a missing population of red $\mu$ m sources just beyond our spectroscopic limit of $^{\prime}$$\sim$ 19.5."
+ There is a caveat., There is a caveat.
+ Unfortunately. it is not possible to iuterpret the data below 1/5 20.5 as the photometric redshifts become unreliable.," Unfortunately, it is not possible to interpret the data below $^{\prime}$$\sim$ 20.5 as the photometric redshifts become unreliable."
+ Nevertheless. the lack of dusty red dawarts does not appear to be caused by the spectroscopic limit of our study.," Nevertheless, the lack of dusty red dwarfs does not appear to be caused by the spectroscopic limit of our study."
+ To derive the star formation rate for the unobscured component we measure the Πο emissiou line and use the conversion of ?.., To derive the star formation rate for the unobscured component we measure the $\alpha$ emission line and use the conversion of \citet{ken98}.
+ The Πα emission was observed through fixed-diameter fiber spectrographs. but Coma galaxies are relatively nearby. so mauv of the galaxies are lareer than the ~3” fiber diameter.," The $\alpha$ emission was observed through fixed-diameter fiber spectrographs, but Coma galaxies are relatively nearby, so many of the galaxies are larger than the $\sim$ $^{\prime\prime}$ fiber diameter."
+ Thus. we apply an aperture correction.," Thus, we apply an aperture correction."
+ Using SExtractor ou the SDSS τή dmaee. we retrieve the isophotal area for cach source.," Using SExtractor on the SDSS $^{\prime}$ image, we retrieve the isophotal area for each source."
+ We assume that the star formation is spread throughout the ealaxy area. thus the ratio of the galaxy area to fiber area gives the aperture correctiou.," We assume that the star formation is spread throughout the galaxy area, thus the ratio of the galaxy area to fiber area gives the aperture correction."
+ Corrections for the largest Coma members can be substantial. up to a factor of teu. however. the 1iedian correction is —2.," Corrections for the largest Coma members can be substantial, up to a factor of ten, however, the median correction is $\sim$ 2."
+ Star formation rates based on optical emission lines suffer C»ereatlv from extinction by surouncdiueC» dust. and so a second iuportaut correction woe apply to the Πα fux measurement is au extinction correction.," Star formation rates based on optical emission lines suffer greatly from extinction by surrounding dust, and so a second important correction we apply to the $\alpha$ flux measurement is an extinction correction."
+ Iu most sources. we can measure the emission line which we use to find the Balmer decrement aud caleulate the amount of extinction.," In most sources, we can measure the $\beta$ emission line which we use to find the Balmer decrement and calculate the amount of extinction."
+ We calculate a significant E(D-V) with a median of 0.55 and standard deviation of 0.12., We calculate a significant E(B-V) with a median of 0.55 and standard deviation of 0.42.
+ Extinction and aperture corrected star formation rates frou. the optical cunissiou lines are given in Table 2.., Extinction and aperture corrected star formation rates from the optical emission lines are given in Table \ref{elines}.
+ We can test how well the star formation rates sed ou total IR huninesities compare to values neasured from corrected. Wa emission line fluxes., We can test how well the star formation rates based on total IR luminosities compare to values measured from corrected $\alpha$ emission line fluxes.
+ Iu Figure 11.. we show the ratio of the SER derived roni the optical cussion line flux to that from SED fitting.," In Figure \ref{sfrouo}, we show the ratio of the SFR derived from the optical emission line flux to that from SED fitting."
+ The aperture correction is more important for he larger galaxies. aud ouce performed briugs a unuiforii ratio of unobseured to obscured star ormnation for galaxies with significant arp (Figure 1... central panel).," The aperture correction is more important for the larger galaxies, and once performed brings a uniform ratio of unobscured to obscured star formation for galaxies with significant $_{SED}$ (Figure \ref{sfrouo}, central panel)."
+ Iucludiug the extinction correction. performed on a source by source basis. he two rates approach each other.," Including the extinction correction, performed on a source by source basis, the two rates approach each other."
+ The median value of Log|SERμι /SFRsgp]| is 0.0. and the," The median value of $_{H\alpha}$ $_{SED}$ ] is 0.0, and the"
+The US Naval Observatory (USNO) operated the Saneh (0.2 11) Twin Astroeraph from 1998 to 2001 for an all-sky astrometric survey.,The US Naval Observatory (USNO) operated the 8-inch (0.2 m) Twin Astrograph from 1998 to 2004 for an all-sky astrometric survey.
+ About 2/3 of the skv was observed from the Cerro Tololo Tater-American Observatory (CTIO) while the rest of the northern sky was observed from the Naval Observatory Flagstaff Station (NOFS)., About 2/3 of the sky was observed from the Cerro Tololo Inter-American Observatory (CTIO) while the rest of the northern sky was observed from the Naval Observatory Flagstaff Station (NOFS).
+ The average number of completed fields per vear was a factor of 2.0 larger at C'TIO than at NOFS., The average number of completed fields per year was a factor of 2.0 larger at CTIO than at NOFS.
+" A tk bv Uk CCD with 9 jan pixel size was used in a single bandpass (579 to 613 nm) providing a flat field of view (FOV) of just over square degree. taking advantage of oulv a tiny fraction of the FOV delivered by the optical svsteii of the Twin Astrograph’s ""red lens.”"," A 4k by 4k CCD with 9 $\mu$ m pixel size was used in a single bandpass (579 to 643 nm) providing a flat field of view (FOV) of just over 1 square degree, taking advantage of only a tiny fraction of the FOV delivered by the optical system of the Twin Astrograph's “red lens.”"
+ A 2-fold overlap pattern of fields spau the eutire sky., A 2-fold overlap pattern of fields span the entire sky.
+" Each field was observed with a loug (about 125 sec) au a short (about 25 sec) exposure. thus cach star should appear on at least 2 different CCD exposures, and stars iu the mid-maguitude range (abou LO to LL) should have ft images."," Each field was observed with a long (about 125 sec) and a short (about 25 sec) exposure, thus each star should appear on at least 2 different CCD exposures, and stars in the mid-magnitude range (about 10 to 14) should have 4 images."
+ UCACS coutains just over 100 million objects. most of these are stars.," UCAC3 contains just over 100 million objects, most of these are stars."
+ It covers the magnitude range of about Ro = S to 16 (Fie., It covers the magnitude range of about R = 8 to 16 (Fig.
+ 1) with positional precision at mean epoch rangiug from 15 to 100 mas. depending ou magnitude (Fie.," 1) with positional precision at mean epoch ranging from 15 to 100 mas, depending on magnitude (Fig."
+ 2)., 2).
+ Mean position errors are shown per 1/10 mag biu with stars up fo maeuitucde 19 excluded whenever the ornal error in either one of the coordinates exceeds LOO as., Mean position errors are shown per 1/10 mag bin with stars up to magnitude 13 excluded whenever the formal error in either one of the coordinates exceeds 100 mas.
+ For fainter stars no such outlier exclusion was adopted. which explaius the discontinuity in Fie.," For fainter stars no such outlier exclusion was adopted, which explains the discontinuity in Fig."
+ 2 aud also shows what effect such a restriction has on the derived mean formal position errors., 2 and also shows what effect such a restriction has on the derived mean formal position errors.
+ The distribution of proper motions is shown iu Fig., The distribution of proper motions is shown in Fig.
+ 3. aud the proper motions errors as a function of magnitude are prescuted in Fie.," 3, and the proper motions errors as a function of magnitude are presented in Fig."
+" 1,", 4.
+ The large lucrease ο| the formal proper motion errors for stars at magnitude 8 aud brighter is caused by the saturation of the CCD data with associated larec. formal positional errors.," The large increase of the formal proper motion errors for stars at magnitude 8 and brighter is caused by the saturation of the CCD data with associated large, formal positional errors."
+ The weighted mean epoch of UCACS data for most stars is in the range of 1980 to 2002 (Fig., The weighted mean epoch of UCAC3 data for most stars is in the range of 1980 to 2002 (Fig.
+ 5). depending ou magnitude as consequence of the observing history of stars aud the positional precisions at various epochs.," 5), depending on magnitude as consequence of the observing history of stars and the positional precisions at various epochs."
+ The released catalog is based on all applicable. regular survey field observations. excliding the CCD exposures taken on extragalactic link fields and most calibration fields.," The released catalog is based on all applicable, regular survey field observations, excluding the CCD exposures taken on extragalactic link fields and most calibration fields."
+ Observations of minor planets have been extracted and will be published separately from UCAC3., Observations of minor planets have been extracted and will be published separately from UCAC3.
+ The released UCACS is a compiled catalog. similar to UCAC2.," The released UCAC3 is a compiled catalog, similar to UCAC2."
+ No individual epoch observations are given. nor are the pixel data publicly available at this poiut.," No individual epoch observations are given, nor are the pixel data publicly available at this point."
+ The Tyeho-2 catalog (Ilogetal.2000) πας used as reference star catalog to obtain UCAC3 positions ou the Iipparcos System (ESA1997). which is the ceurreut optical realization of the luternational Celestial Reference. Frame ICRF).," The Tycho-2 catalog \citep{tycho2} was used as reference star catalog to obtain UCAC3 positions on the Hipparcos System \citep{hipcat}, which is the current optical realization of the International Celestial Reference Frame (ICRF)."
+ Most stars in UCAC3 have proper motions which were derives your the astrograph CCD data combined with various earlier epoch data. including all exouud-based. catalogs used also for the Tyeho-2 project. unpublished new iieasurements of other astrograph plates. the Southern Proper Motion (SPM) first epoch plates. auc σοι] plate data through the SuperCOSAIOS project.," Most stars in UCAC3 have proper motions which were derived from the astrograph CCD data combined with various earlier epoch data, including all ground-based catalogs used also for the Tycho-2 project, unpublished new measurements of other astrograph plates, the Southern Proper Motion (SPM) first epoch plates, and Schmidt plate data through the SuperCOSMOS project."
+ A final UCACL release is planned which will utilize the new reductions of the Northern Proper Motion (NPAD) program. supplementing the SPM data. which then would allow us to derive proper motions for all UCAC stars without the use of Schinidt plate data.," A final UCAC4 release is planned which will utilize the new reductions of the Northern Proper Motion (NPM) program, supplementing the SPM data, which then would allow us to derive proper motions for all UCAC stars without the use of Schmidt plate data."
+ This goal could uot be achieved for UCAC?23 due to a production deadline alb lack of time to complete the NPM work., This goal could not be achieved for UCAC3 due to a production deadline and lack of time to complete the NPM work.
+ The 2-Micron. Al-Sky Survey. 2MASS (1) was used extensively to analyze systematic errors of UCACS3 data and to supplement the UCACS catalog with near IR. photometry.," The 2-Micron All-Sky Survey, 2MASS \citep{2mass} was used extensively to analyze systematic errors of UCAC3 data and to supplement the UCAC3 catalog with near IR photometry."
+" Optical B.R.I magnitudes were copied from the Super""OSMOS source catalog (photographic plotometiy) iuto UCACS for the benefit of the users."," Optical B,R,I magnitudes were copied from the SuperCOSMOS source catalog (photographic photometry) into UCAC3 for the benefit of the users."
+ The uuuber of UCACS objects matched with various catalogs is preseuted in Table 1., The number of UCAC3 objects matched with various catalogs is presented in Table 1.
+ For more details about the observational data and earlier reductions the reader is referred to the UCACI (Zachariasetal.2000). aud UCAC2 (Zachariasetal.2001) papers., For more details about the observational data and earlier reductions the reader is referred to the UCAC1 \citep{ucac1} and UCAC2 \citep{ucac2} papers.
+ Coutrary to those papers. which cach describe one of the earlier releases du detail. the UCACA2 effort. will be documented in a series of papers.," Contrary to those papers, which each describe one of the earlier releases in detail, the UCAC3 effort will be documented in a series of papers."
+ This paper eives the introduction aiming at the user of the UCAC3 catalog. describing the released data. lamitations. and comparisons to other catalogs.," This paper gives the introduction aiming at the user of the UCAC3 catalog, describing the released data, limitations, and comparisons to other catalogs."
+ Technical details of the reduction process will be outlined iu a paper about the new pixel processing (Zacharias2010) anc a separate paper ou the astrometric reductions leading to the, Technical details of the reduction process will be outlined in a paper about the new pixel processing \citep{pxred} and a separate paper on the astrometric reductions leading to the
+the source arising in the inner region.,the source arising in the inner region.
+" As a consequence, the relativistic effects are strongly enhanced."," As a consequence, the relativistic effects are strongly enhanced."
+ The aforementioned increase in the variance with inclination is slower and the overall variability level is lower if the black hole is not rotating., The aforementioned increase in the variance with inclination is slower and the overall variability level is lower if the black hole is not rotating.
+" In Figure 2,, we show the inclination dependence of the normalized variance for the same set of radial profiles B as in Fig. 1,,"," In Figure \ref{fig:a0}, we show the inclination dependence of the normalized variance for the same set of radial profiles $\beta$ as in Fig. \ref{fig:a95},"
+ but now with the Schwarzschild black hole., but now with the Schwarzschild black hole.
+" The ratio of the variances for 8=3.5 at 0,=81.6? and 0,=30? is equal to 3.5 for an a=0 black hole, but increases to 3.7 for an a=0.95 Kerr black hole in the presented plots."," The ratio of the variances for $\beta = 3.5$ at $\incli = 81.6^{\circ}$ and $\incli=30^{\circ}$ is equal to $3.5$ for an $a = 0$ black hole, but increases to $3.7$ for an $a = 0.95$ Kerr black hole in the presented plots."
+ The rise in the variance with the inclination angle is not a result of the statistical errors in simulations., The rise in the variance with the inclination angle is not a result of the statistical errors in simulations.
+ It is well known that a single simulation of short timescale is strongly affected by the power leaking from the lower frequencies (Vaughan 2003)., It is well known that a single simulation of short timescale is strongly affected by the power leaking from the lower frequencies \citep{vaughan03}.
+". However, in Figs."," However, in Figs."
+ 1 and 2 the sequences for fixed values of 8 were calculated for the same realizations of the same statistical distribution., \ref{fig:a95} and \ref{fig:a0} the sequences for fixed values of $\beta$ were calculated for the same realizations of the same statistical distribution.
+ The rise in variance is entirely due to the change in the viewing angle., The rise in variance is entirely due to the change in the viewing angle.
+ If we use different random realizations of the flare distribution the trend is less clear because of large statistical dispersion in the adopted length of the light curve., If we use different random realizations of the flare distribution the trend is less clear because of large statistical dispersion in the adopted length of the light curve.
+ A single variance (in simulations as well as in the actual data) is determined with an accuracy of a factor of two., A single variance (in simulations as well as in the actual data) is determined with an accuracy of a factor of two.
+" Therefore, to show the model predictions to higher accuracy, we had to extend the simulated light curves by a factor of eight."," Therefore, to show the model predictions to higher accuracy, we had to extend the simulated light curves by a factor of eight."
+" For these longer light curves, we calculated the variance enhancement, defined as the ratio of the variance seen at 81.6? to that at 30, as a function of the Kerr parameter."," For these longer light curves, we calculated the variance enhancement, defined as the ratio of the variance seen at $81.6^ \circ$ to that at $30^
+\circ$, as a function of the Kerr parameter."
+" The result is shown in Fig. 3,,"," The result is shown in Fig. \ref{fig:ratio},"
+ for three values of the flare time-scale duration., for three values of the flare time-scale duration.
+" We also plot the ratio of the observed variance, given in Eq. (6)),"," We also plot the ratio of the observed variance, given in Eq. \ref{eq:var_value}) ),"
+ to the variance expected from Eq. (, to the variance expected from Eq. (
+4) of Nikotajuk (2009).,4) of \citet{nikolajuk09}.
+". This ratio is found to equal 4.0, its standard deviation (107) error coming from the errors in both the variance and in scaling constant."," This ratio is found to equal $4.0$, its standard deviation $1\sigma$ ) error coming from the errors in both the variance and in scaling constant."
+" Because of a lack of information about the black hole spin, the observational constraint is indicated by straight lines irrespective of the value of a."," Because of a lack of information about the black hole spin, the observational constraint is indicated by straight lines irrespective of the value of $a$."
+" The variance enhancement, in general, increases with the Kerr parameter but the detailed profile depends on the assumed scaling factor of the flare lifetime."," The variance enhancement, in general, increases with the Kerr parameter but the detailed profile depends on the assumed scaling factor of the flare lifetime."
+ The variability is yet greater when the flare lifetime close to the black hole is comparable to the time spent by the flare in the region of the highest Doppler boosting., The variability is yet greater when the flare lifetime close to the black hole is comparable to the time spent by the flare in the region of the highest Doppler boosting.
+ The angular extension of this region (as a fraction of 2m) decreases as the Kerr parameter increases (Dovéiaketal. 2004a)., The angular extension of this region (as a fraction of $2 \pi$ ) decreases as the Kerr parameter increases \citep{dovciak04a}.
+". The enhanced variability is consistent with the expectations of the relativistic enhancement, within the framework of the flare model."," The enhanced variability is consistent with the expectations of the relativistic enhancement, within the framework of the flare model."
+" A flare duration of longer timescale, τρ=104 s, is indicative of a non-rotating black hole, while the flares scaled down to ro=10? s probably correspond to a spinning black hole with a>0.5."," A flare duration of longer timescale, $\tau_0 = 10^4$ s, is indicative of a non-rotating black hole, while the flares scaled down to $\tau_0 =
+10^3$ s probably correspond to a spinning black hole with $a > 0.5$."
+" However, taking into account the large error in the observed variance and the lack of a priori knowledge of the flare timescale we cannot at this stage make any firm conclusion about the black hole rotation in NGC 4258."," However, taking into account the large error in the observed variance and the lack of a priori knowledge of the flare timescale we cannot at this stage make any firm conclusion about the black hole rotation in NGC 4258."
+ In the flare model of AGN X-ray variability the change in the X-ray flux is caused by a combination of two main effects., In the flare model of AGN X-ray variability the change in the X-ray flux is caused by a combination of two main effects.
+ The first is the intrinsic variability of both a single flare and the flare distribution., The first is the intrinsic variability of both a single flare and the flare distribution.
+ The second is the variation caused by the relativistic effects of the flare orbital (Keplerian) motion., The second is the variation caused by the relativistic effects of the flare orbital (Keplerian) motion.
+ Analyzing the change in the level of variability with the inclination we can disentangle those two effects., Analyzing the change in the level of variability with the inclination we can disentangle those two effects.
+ We chose NGC 4258 as a target of our study because of its exceptional properties., We chose NGC 4258 as a target of our study because of its exceptional properties.
+" The source is highly inclined, seen almost edge-on, but is nevertheless Compton thin and the variable primary emission is still transmitted effectively through the material located at the source equatorial plane, although the X-ray emission is significantly absorbed."," The source is highly inclined, seen almost edge-on, but is nevertheless Compton thin and the variable primary emission is still transmitted effectively through the material located at the source equatorial plane, although the X-ray emission is significantly absorbed."
+" In addition, the mass of the central black hole in this object is measured accurately thanks to the spatially resolved water maser emission."," In addition, the mass of the central black hole in this object is measured accurately thanks to the spatially resolved water maser emission."
+" On the other hand, we modified our original flare model (Czernyetal.2004) by including the exponential decay of individual flares and the effect of an avalanche-type development of flares."," On the other hand, we modified our original flare model \citep{czerny04} by including the exponential decay of individual flares and the effect of an avalanche-type development of flares."
+" On the other hand, we neglected the cold disk reflection."," On the other hand, we neglected the cold disk reflection."
+" In NGC 4258, any manifestations of reflection features, such as the broad relativistic iron line and temporary spots on the disk surface, are strongly suppressed."," In NGC 4258, any manifestations of reflection features, such as the broad relativistic iron line and temporary spots on the disk surface, are strongly suppressed."
+ We thus understand the situation by presuming that the central corona is the source of enhanced variability and the place where flares originate., We thus understand the situation by presuming that the central corona is the source of enhanced variability and the place where flares originate.
+" After developing a refined version of the model, we analyzed the dependence of the X-ray variance on the inclination angle."," After developing a refined version of the model, we analyzed the dependence of the X-ray variance on the inclination angle."
+ We noticed that the rise in the variability with the inclination angle is faster if the emissivity is more, We noticed that the rise in the variability with the inclination angle is faster if the emissivity is more
+μας.,.
+" Here. the lower cutolf wavenumber is Ai,=2. the upper cutoll wavenumber zdy3 is determined by the plasma kinematic viscosity 7. and the helicity parameter 5h must satisfy the realizability condition —1</hxl1: the velocity field. is maximally helical when ||=1.7 We choose magnetic diffusivity to be jj=105. corresponding to magnetic Revnolds number Rai~100."," Here, the lower cutoff wavenumber is $k_{\rm min}=2$, the upper cutoff wavenumber $k_{\rm max}\approx2[\kappa_L(0)/\nu]^{3/4}\approx4\nu^{-3/4}$ is determined by the plasma kinematic viscosity $\nu$, and the helicity parameter $h$ must satisfy the realizability condition $-1\le h\le1$; the velocity field is maximally helical when $|h|=1$ We choose magnetic diffusivity to be $\eta=10^{-6}$, corresponding to magnetic Reynolds number ${\rm Rm}\sim 10^6$."
+" Dv changing the value of the turbulence cutoff. wavenumber Mya. in Equation (10)). we varv the Revnolds number RezzL/r&(F4,/4) and the magnetic Prandtl number Pm=Rin/Rezv/ip."," By changing the value of the turbulence cutoff wavenumber $k_{\rm max}$ in Equation \ref{POWER_V}) ), we vary the Reynolds number ${\rm Re}\approx 1/\nu\approx(k_{\rm max}/4)^{4/3}$ and the magnetic Prandtl number ${\rm Pm}={\rm Rm}/{\rm Re}\approx\nu/\eta$."
+ We study (vo cases for the Prandt] number: a case when itis large and a case when it is small., We study two cases for the Prandtl number: a case when itis large and a case when it is small.
+ Ii the first case. Pin~15029 1. which is achieved," In the first case, ${\rm Pm}\sim 150\gg 1$ , which is achieved"
+The reconstructed two-dimensional mass distribution appears to be centered on the cD galaxy and shows elliptical profiles oriented in. direction North-South.,The reconstructed two-dimensional mass distribution appears to be centered on the cD galaxy and shows elliptical profiles oriented in direction North-South.
+ No substructure on scales larger than 30” was detected., No substructure on scales larger than $30''$ was detected.
+ In order to remove the (see. e.g.. Kaiser Squires 1993). we fitted the mass profiles with non-singular isothermal sphere models: from this we obtain. for example. M(r.«lhΜρε]=5.3«101!ho} M..," In order to remove the (see, e.g., Kaiser Squires 1993), we fitted the mass profiles with non-singular isothermal sphere models; from this we obtain, for example, $M(r < 1 \, h_{50}^{-1} \mbox{ Mpc}) 
+= 5.3 \times 10^{14} \, h_{50}^{-1} \mbox{M}_\odot$ ."
+ We note that. because of the smoothing operated in the two-dimensional mass maps. the weak lensing mass profile for;<1' is bound to be underestimate.," We note that, because of the smoothing operated in the two-dimensional mass maps, the weak lensing mass profile for $r < 1'$ is bound to be underestimate."
+" On the other hand. at large radii (say. ο> 3%. an error on the removal of the mass-sheet degeneracy can in principle lead to an unreliable ""total"" weak lensing mass estimate."," On the other hand, at large radii (say, $r > 3'$ ), an error on the removal of the mass-sheet degeneracy can in principle lead to an unreliable “total” weak lensing mass estimate."
+ Note that the four profiles (from B. V. R and I optical images) agree very well to each other. which strongly support the results of our analysis.," Note that the four profiles (from B, V, R and I optical images) agree very well to each other, which strongly support the results of our analysis."
+ The differential X-ray best-fit mass model has been weighted by the relative portion of the shell observed in each ring and. then. cumulated up to the radius of 1020 kpe (the outer radius of our spatially resolved spectroscopy).," The differential X-ray best-fit mass model has been weighted by the relative portion of the shell observed in each ring and, then, cumulated up to the radius of 1020 kpc (the outer radius of our spatially resolved spectroscopy)."
+ In Fig. 5..," In Fig. \ref{fig:mass},"
+ we plot and compare the projected mass profiles of the galaxy cluster 1224 obtained from both the spatially resolved spectral analysis of the observation and the weak lensing analysis of FORSI-VLT multicolor imaging., we plot and compare the projected mass profiles of the galaxy cluster $-$ 1224 obtained from both the spatially resolved spectral analysis of the observation and the weak lensing analysis of FORS1-VLT multicolor imaging.
+ The two independently reconstructed mass profiles agree very well within 1o uncertainty both in absolute values and in the overall shape of the profile., The two independently reconstructed mass profiles agree very well within $1 \sigma$ uncertainty both in absolute values and in the overall shape of the profile.
+ Note that the mass center is fixed to the peak of the lensing map density that is consistent with the X-ray peak as discussed in Sect., Note that the mass center is fixed to the peak of the lensing map density that is consistent with the X-ray peak as discussed in Sect.
+ 2., 2.
+ This result does not change when we consider the different density and temperature profiles observed in the northern region where a significant surface brightness excess is located and a higher signal-to-noise ratio is available. arguing for the robustness of the X-ray mass estimates once gas density and temperature distributions can be properly mapped.," This result does not change when we consider the different density and temperature profiles observed in the northern region where a significant surface brightness excess is located and a higher signal-to-noise ratio is available, arguing for the robustness of the X-ray mass estimates once gas density and temperature distributions can be properly mapped."
+»ositive time lage between the 0.52 and 2δ.δ keV. bands is approximately SOO s at low Fourier. [frequencies and ~ Os at high. frequencies. with a sharp drop in time ag values at a frequency of ~10* Iz.,"positive time lag between the 0.5–2 and 2–8.8 keV bands is approximately 800 s at low Fourier frequencies and $\sim 0$ s at high frequencies, with a sharp drop in time lag values at a frequency of $\sim 10^{-4}$ Hz."
+ This stepped appearance of the lag spectrum resembles that seen in ονο X-I in ie [owπαν ancl intermediate states. where he steps in 1e lag spectrum correspond to the transition pequencies between peaks in the corresponding PSDs (Nowaketal.1999:Nowak2000).," This stepped appearance of the lag spectrum resembles that seen in Cyg X-1 in the low/hard and intermediate states, where the steps in the lag spectrum correspond to the transition frequencies between peaks in the corresponding PSDs \citep{nowak99,nowak00}."
+. The correspondence tween. lag spectrum steps and PSD peaks suggests that cach constant-lag value is associated with a given variability component. seen as a Lorentzian component in the PSD.," The correspondence between lag spectrum steps and PSD peaks suggests that each constant-lag value is associated with a given variability component, seen as a Lorentzian component in the PSD."
+ Therefore. the shape of the lag spectrum in ssugeests that the two-component interpretation of the PSD is Correct.," Therefore, the shape of the lag spectrum in suggests that the two-component interpretation of the PSD is correct."
+ To test the above hypothesis we clireethy fitted the lag spectrum with a mocel corresponding to the lags expected from two Lorentzian variability components. cach with a single distinct time lag. which is constant for all Fourier [requencies.," To test the above hypothesis we directly fitted the lag spectrum with a model corresponding to the lags expected from two Lorentzian variability components, each with a single distinct time lag, which is constant for all Fourier frequencies."
+ Phe observed lag at any particular frequency. is then a function of the overlap between the two Lorentzians. and must be caleulated in the complex Fourier. domain. by evaluating the cross-spectrum as a function of Fourier [requeney. p.," The observed lag at any particular frequency is then a function of the overlap between the two Lorentzians, and must be calculated in the complex Fourier domain, by evaluating the cross-spectrum as a function of Fourier frequency, $\nu$."
+ The cross-spectrum is given by C'(p)=SQL Go). where 4G) and Sv) ave the Fourier transforms of the hard. and soft. band light. curves and the asterisk denotes the complex. conjugate.," The cross-spectrum is given by $C(\nu)=S(\nu)H^{*}(\nu)$ , where $H(\nu)$ and $S(\nu)$ are the Fourier transforms of the hard and soft band light curves and the asterisk denotes the complex conjugate."
+ ο)=Sv)|SG) where the subscripts / and. / denote the low ane high-requency Lorentzians respectively. and similarly for the vel band. (9)=Hitv)|ff).," $S(\nu)=S_l(\nu)+S_h(\nu)$ where the subscripts $l$ and $h$ denote the low and high-frequency Lorentzians respectively, and similarly for the hard band, $H(\nu)=H_l(\nu)+H_h(\nu)$."
+" For the low and high requency Lorentzian components. the soft band light curve is correlated. with the hard band light curve except with a shase lag ój(9). o,(G) respectively."," For the low and high frequency Lorentzian components, the soft band light curve is correlated with the hard band light curve except with a phase lag $\phi_l(\nu)$, $\phi_h(\nu)$ respectively."
+ We determine the phase ag from the fixed low and. high-frequency. Lorentzian ages T and το bv using ó(p)=2zxvr.," We determine the phase lag from the fixed low and high-frequency Lorentzian lags $\tau_l$ and $\tau_h$, by using $\phi(\nu)=2\pi \nu \tau$."
+ lo avoid clutter. we now drop the frequeney-depencdence of the various xwameters and. write the real part of the cross-speetrum as follows: where Re and. £m denote real and imaginary components.," To avoid clutter, we now drop the frequency-dependence of the various parameters and write the real part of the cross-spectrum as follows: where $Re$ and $Im$ denote real and imaginary components."
+" Oy and ©, are simply random. phase values for the low and high frequency. Lorentzians respectively. measured. for each realisation of the light. curves. and. will be cilferent for cillerent realisations."," $\phi_0$ and $\phi_1$ are simply random phase values for the low and high frequency Lorentzians respectively, measured for each realisation of the light curves, and will be different for different realisations."
+" “Phey are not related to the phase lags of interest between bands. Le. ó; and ©, for the"," They are not related to the phase lags of interest between bands, i.e. $\phi_l$ and $\phi_h$ for the"
+"The isothermal sections of extrasolar gas giant atmospheres often reach temperatures close to 2000 K (Spiegeletal.2009) and in some cases, even higher (Boruckietal.2009)..","The isothermal sections of extrasolar gas giant atmospheres often reach temperatures close to 2000 K \citep{2009ApJ...699.1487S} and in some cases, even higher \citep{2009Sci...325..709B}."
+" These temperatures are not high enough to ionize H or He significantly, however, alkali metals such as Na and K will be partially ionized."," These temperatures are not high enough to ionize H or He significantly, however, alkali metals such as Na and K will be partially ionized."
+" As a result, electrical conductivity in the interior of a hot Jupiter is dominated by ionization of hydrogen, while in the outer region of the planet, electrical conductivity is primarily due to the ionization of alkali metals, with the transition between the two inoization regimes taking place at P~300 Bars."," As a result, electrical conductivity in the interior of a hot Jupiter is dominated by ionization of hydrogen, while in the outer region of the planet, electrical conductivity is primarily due to the ionization of alkali metals, with the transition between the two inoization regimes taking place at $P\sim300$ Bars."
+" Thermal ionization is governed by the Saha equation: where n; and nj are the total and positively ionized number densities of constituent 7 respectively, ne= is the total electron number density, m, is the “nFelectron mass, ky is Boltzmann’s constant, T is temperature, fA is Plank’s constant, and J; is the ionization potential of constituent j."," Thermal ionization is governed by the Saha equation: where $n_{j}$ and $n_{j}^{+}$ are the total and positively ionized number densities of constituent $j$ respectively, $n_{e} = \sum n_{j}^{+}$ is the total electron number density, $m_{e}$ is the electron mass, $k_{b}$ is Boltzmann's constant, $T$ is temperature, $\hbar$ is Plank's constant, and $I_{j}$ is the ionization potential of constituent $j$ ."
+" If the ionization is far from complete (nj« fjn), the abundances of alkali metals, f; are held constant, and the atmosphere is isothermal, it is easy to show that the electron number density takes on an exponential profile with an ionization scale-height that is twice as large as the density heght: where y’s are the RHS’s of equation (1), ro is the radial distance at some reference point (P—10 and H=kyT/ug is the density scale-height."," If the ionization is far from complete $n_j^{+}\ll f_jn$ ), the abundances of alkali metals, $f_j$ are held constant, and the atmosphere is isothermal, it is easy to show that the electron number density takes on an exponential profile with an ionization scale-height that is twice as large as the density scale-heght: where $\chi$ 's are the RHS's of equation (1), $r_0$ is the radial distance at some reference point $P=10$ Bars) and $H=k_bT/\mu g$ is the density scale-height."
+" In our Bars)ionization calculations, we considered the following alkali metals: Na, K, Li, Rb, Fe, Cs and Ca."," In our ionization calculations, we considered the following alkali metals: Na, K, Li, Rb, Fe, Cs and Ca."
+ Their abundances and ionization potentials were inferred from Lodders(1999) and Coxetal.(2000) respectively., Their abundances and ionization potentials were inferred from \cite{1999ApJ...519..793L} and \cite{2000asqu.book..499C} respectively.
+ The atmospheric temperatures above the isothermal layer differ significantly from planet to planet., The atmospheric temperatures above the isothermal layer differ significantly from planet to planet.
+" In particular, thermal inversions have been detected in the atmospheres of HD209458b (Burrowsetal.2007) and Tres-4b (Knutsonetal.2009) but not in HD189733b."," In particular, thermal inversions have been detected in the atmospheres of HD209458b \citep{2007ApJ...668L.171B} and Tres-4b \citep{2009ApJ...691..866K} but not in HD189733b."
+" In our models, we adopt atmospheric temperature profiles similar to that of Spiegeletal. for HD209458b and Tres-4b, and the 1D profile of (2009)Fortneyetal.(2010) for HD189733b."," In our models, we adopt atmospheric temperature profiles similar to that of \cite{2009ApJ...699.1487S} for HD209458b and Tres-4b, and the 1D profile of \cite{2010ApJ...709.1396F} for HD189733b."
+ The relatively cool temperatures attained above P <0.1 Bars are of significant importance to our models because they provide insulating shells which are impenetrable to radial current., The relatively cool temperatures attained above P $\lesssim0.1$ Bars are of significant importance to our models because they provide insulating shells which are impenetrable to radial current.
+" Consequently, current loops are necessarily setup through the interior, and any current flowing in the ionosphere is not relevant."," Consequently, current loops are necessarily setup through the interior, and any current flowing in the ionosphere is not relevant."
+ We place the radiative/convective boundary at P~100 Bars in all of our models., We place the radiative/convective boundary at $ P\sim100$ Bars in all of our models.
+" We did not have to explicitly compute the ionization fractions of H and He, as they are published in the equation of state (Saumonetal.1995), which we employed in our model."," We did not have to explicitly compute the ionization fractions of H and He, as they are published in the equation of state \citep{1995ApJS...99..713S}, which we employed in our model."
+" In particular, we used the ""interpolated"" version of the equation of state, where ionization occurs smoothly with pressure and temperature."," In particular, we used the ""interpolated"" version of the equation of state, where ionization occurs smoothly with pressure and temperature."
+" Although the planetary structure was core-less, we mimicked the presence of a core by changing the Helium content from Y—0.24 to Y—0.3 al.2003) in some of our models."," Although the planetary structure was core-less, we mimicked the presence of a core by changing the Helium content from $Y=0.24$ to $Y=0.3$ \citep{2003ApJ...594..545B} in some of our models."
+" Having computed the electron number density, the electrical conductivity of a gas is given by (TiplerLlewellyn2002) where n and A are the number density, and the number density weighted cross-section of everything other than electrons."," Having computed the electron number density, the electrical conductivity of a gas is given by \citep{2002moph.book.....T}
+ where $n$ and $A$ are the number density, and the number density weighted cross-section of everything other than electrons."
+" Strictly speaking, the above equation is only valid for non-degenerate gas."," Strictly speaking, the above equation is only valid for non-degenerate gas."
+" However, by the point matter becomes degenerate in our models, the resistivity is completely negligible and the details of its profile have no noticeable effect on the results."," However, by the point matter becomes degenerate in our models, the resistivity is completely negligible and the details of its profile have no noticeable effect on the results."
+" Since we are only interested in the part of the planet, interior to the atmospheric temperature minimum, we define the model radius r=R as the point of maximal conductivity in the atmosphere (P=75 mbars), and we set the outer edge of our model at the conductivity minimum, r=R++7 (P=30 mbars)."," Since we are only interested in the part of the planet, interior to the atmospheric temperature minimum, we define the model radius $r=R$ as the point of maximal conductivity in the atmosphere $P = 75$ mbars), and we set the outer edge of our model at the conductivity minimum, $r=R+\gamma$ $P=30$ mbars)."
+ We place the bottom boundary of the “weather” layer of the atmosphere at a pressure of P=10 Bars and denote it as r=R—ó.," We place the bottom boundary of the “weather"" layer of the atmosphere at a pressure of $P=10$ Bars and denote it as $r=R-\delta$."
+" Consequently, the “inert” layer of the atmosphere is between 100XP<10 Bars."," Consequently, the “inert"" layer of the atmosphere is between $100\lesssim P\lesssim10$ Bars."
+" A computed electrical conductivity profile for HD209458b is presented in figure (2), along with a simplified conductivity profile resulting from equation (2)."," A computed electrical conductivity profile for HD209458b is presented in figure (2), along with a simplified conductivity profile resulting from equation (2)."
+" Because the functional profiles (dashed curve) are in good agreement with the numerically computed profile, we utilize them in all future calculations (see appendix)."," Because the functional profiles (dashed curve) are in good agreement with the numerically computed profile, we utilize them in all future calculations (see appendix)."
+" Global circulation models (Showmanetal.2008,2009; have shown that winds on hot Jupiters, specifically HD209458b and HD189733b, can attain velocities oforder v~ 1km/s. It appears that two qualitative wind patterns are present."," Global circulation models \citep{2008ApJ...682..559S,2009ApJ...699..564S,2008ApJ...674.1106L,2009ApJ...700..887M} have shown that winds on hot Jupiters, specifically HD209458b and HD189733b, can attain velocities oforder $v\sim1$ km/s. It appears that two qualitative wind patterns are present."
+" In the upper atmosphere (P< 30mbars), wind flows from the sub-stellar point to the anti-stellar point symmetrically across the terminator."," In the upper atmosphere $P\lesssim30$ mbars), wind flows from the sub-stellar point to the anti-stellar point symmetrically across the terminator."
+" Deeper down, a strong eastward zonal jet develops."," Deeper down, a strong eastward zonal jet develops."
+"where The halo-halo power spectrum {ο(0) aud bispectrum: By(sy.45.η). are related to the linear mass power spectrum Zu,(&) by equations (10)) aud (13)).","where The halo-halo power spectrum $P_\halo(k)$ and bispectrum $B_\halo(k_1,k_2,k_3)$ are related to the linear mass power spectrum $P_\lin(k)$ by equations \ref{Pbias}) ) and \ref{Bbias}) )."
+ The expressious for the mass bispectrum above simplily cousiderably for the equilateral triaugle configuration. and where Here we have written out explicitly the bias factors 6(A) using equatious (10)) aud (15)). aud we lave neglected terms with 65(A) as discussecl in 82.3.," The expressions for the mass bispectrum above simplify considerably for the equilateral triangle configuration, and where Here we have written out explicitly the bias factors $b(M)$ using equations \ref{Pbias}) ) and \ref{Bbias2}) ), and we have neglected terms with $b_2(M)$ as discussed in 2.3."
+ Iu this section we compare tlie predictions of our analytical model described in 82. 3. and. E with results [romcosmological N-body simulations.," In this section we compare the predictions of our analytical model described in 2, 3, and 4 with results fromcosmological $N$ -body simulations."
+ We examine two cosmological models: an i=—2 scale-free 1uodel aud a low-cleusity ACDM wmocel., We examine two cosmological models: an $n=-2$ scale-free model and a low-density $\Lambda$ CDM model.
+ These are the same simulations studied in Ma Fry (2000a)., These are the same simulations studied in Ma Fry (2000a).
+ The v=—2 simulation has 256% particles aud a Plummer force softening length. of L/25120. where L is the box length.," The $n=-2$ simulation has $256^3$ particles and a Plummer force softening length of $L/5120$ , where $L$ is the box length."
+" The ACDAI model is spatially flat with matter density ,,=0.3+) and cosmological constant O4= 0.7.", The $\Lambda$ CDM model is spatially flat with matter density $\Omega_m=0.3$ and cosmological constant $\ov=0.7$ .
+ This run has 128? particles aud is performed in a (100Mpc)? comoving box with a comoviug force softening length. of 250kpc for Hubble parameter //= 0.75.," This run has $128^3$ particles and is performed in a $
+(100\,{\rm Mpc})^3$ comoving box with a comoving force softening length of $ 50\,{\rm kpc} $ for Hubble parameter $ h=0.75 $ ."
+ The baryon [raction is set to zerofor simplicity., The baryon fraction is set to zerofor simplicity.
+ The primordialpower spectrum has a spectral, The primordialpower spectrum has a spectral
+⊟≻↥⋅↑∐⊔↸∖⊺∠≺∐∖↑↸∖↥⋅∐∐∐↸∖≼∏⋝∙↖⇁↑↕∐∖↕∪↥⋅↕⊔⋜↧↕↴∖↴↸∖↥⋅↕↸∖↴∖↴≪↗⊤∶∑∖↽ . B d B ⋅∖⋜∐⋅↴⋝↕⊓⋅⋜∐⋅⋅↖↽↸⊳∪∐↴∖↴↑⋜⋯↑↴∖↴∙↑∐↸∖∶↴∙⊾↸∖∐↸∖↥⋅⋜↧↕⋀∖⊽∫⇀≋↕���↸∖↥⋅⋜∐⋅↸⊳∐⋅↖↽↸∖≺∣∏⋜↧↑↕∪∐↕↴∖↴ Iuteerabilitv of this equation is associated with the Zakharov-Shabat problem (36)) and the time evolutiou where Iu the last equation we have used that for VW=0. Cy=0.,"For time $t$ determined by the formal series $\partial_t = {\sum^\infty_{N=0}} E_N \partial_{t_{N}}$ where $E_N$ are arbitrary constants, the general NLS hierarchy equation is Integrability of this equation is associated with the Zakharov-Shabat problem \ref{ZS11}) ) and the time evolution where In the last equation we have used that for $N=0$, $C_0 = 0$."
+ Then we have The above equation (19)) eives integrable noulinear exteusion of linear Selidiinger equation with eeneral analytic dispersion cousidered iu Section 1., Then we have The above equation \ref{GNLShierarchy}) ) gives integrable nonlinear extension of linear Schrd̈iinger equation with general analytic dispersion considered in Section 1.
+ Let one considers the classical particle svstem with the enerev-niomoentum relation E(p)=Ey|Eyp|Bop? |... , Let one considers the classical particle system with the energy-momentum relation $E(p) = E_0 + E_1 p + E_2 p^2 + ...$ .
+Then the corresponding time-dependent Schréddinger wave equation is (1)) where the Ihuniltonian operator results from the standard substitution for inoiieutuni p>ihe in the dispersion.," Then the corresponding time-dependent Schröddinger wave equation is \ref{Schrodinger}) ) where the Hamiltonian operator results from the standard substitution for momentum $p \rightarrow -i \hbar
+\frac{\partial}{\partial x}$ in the dispersion."
+ Equation (1)) together with its complex conjugate can be rewritten as The momentum operator here is just the recursion operator in the linear approximation Ry=loa.," Equation \ref{Schrodinger}) ) together with its complex conjugate can be rewritten as The momentum operator here is just the recursion operator in the linear approximation ${\cal{R}}_0 = i
+\sigma_3 \frac{\partial}{\partial x}$."
+ ence. (53)) is the linear Schroddinger equation with arbitrary analytic dispersion.," Hence, \ref{MSchr}) ) is the linear Schröddinger equation with arbitrary analytic dispersion."
+" The noulinear integrable extension of this equation appears as (19)). which corresponds to the replacement Ry>R. (fh= 1). so that From this point of view. the standard substitution for classical 1onientun """" . .) . ⋅⋅ ⋅ ↕⋟∪↥⋅⋯∙∶↴∙⊾↕↖↽↸∖↴∖↴≺∣∏⋜⋯↑↕∑⋜↧↑↕∪∐↕∐↑∐↸∖↕⋟∪↥⋅⋯∪↕≯↑∐↸∖∐∐↸∖⋜∐⋅≋↸⊳↕∐⋅≺⋈⋉∐∐∐∶↴⋁↸∖↥⋅↸∖≺∣∏⋜↧↑↕∪∐∙↖↖⊽∐∐↸∖ ↴∖↴∏↴"," The nonlinear integrable extension of this equation appears as \ref{GNLShierarchy}) ), which corresponds to the replacement ${\cal{R}}_0 \rightarrow {\cal{R}}$, $\hbar = 1$ ), so that From this point of view, the standard substitution for classical momentum $p \rightarrow -i \hbar \frac{\partial}{\partial x}$ or equivalently $p \rightarrow -i \hbar
+\sigma_3\frac{\partial}{\partial x} = {\cal{R}}_0$ for the equation in spinor form, gives quantization in the form of the linear Schröddinger equation."
+⋝↴∖↴↑↕↑∏↑↕∪∐∣↗≻↗∖∣≦↴∙⊾↕↖⇁↸∖↴∖↴∎⊲∐∪∐∐∐↸∖⋜∐⋅≺∣∏⋜⋯↑↕∑⋜↧⊓∪∐∥⋜⋯≼↧↑↕∐∖∐∪∐∐∐↸∖⋜∐⋅≋↸⊳↕∐⋅≺⋈⋉∐∐∐∶↴⋁↸∖↥⋅ ↕∐↸∖↥⋅⋜∐⋅↸⊳∐⋅↖⇁↸∖≺∣∏⋜↧↑↕∪∐∙," While substitution $p \rightarrow
+{\cal{R}}$ gives ""nonlinear quantization"" andthe nonlinear Schröddinger hierarchy equation."
+the shell can escape the ealaxv entirely by eas pressure (xine2003)..,the shell can escape the galaxy entirely by gas pressure \citep{King2}.
+ Therefore. given an adecquate mass supply (such as in a merger). we get (Ixing2005) Ilere. the proportionality. constant fj&/z.G7a lies1. within. the observational. constraints.," Therefore, given an adequate mass supply (such as in a merger), we get \citep{King}
+ Here, the proportionality constant $f_g \kappa/ \pi G^2$ lies within the observational constraints."
+". To sununarize. the Mj),—0 relation is obtained with three important. assumptions: (1) isothermal gas density distribution throughout the galaxy. formation. (2) super Eddineton accretion. and (3) an aclequate mass supply."," To summarize, the $M_{BH,f}- \sigma$ relation is obtained with three important assumptions: (1) isothermal gas density distribution throughout the galaxy formation, (2) super Eddington accretion, and (3) an adequate mass supply."
+" MacMillanandHenriksen(2002) obtained a relation between Api), and o by assuming that the density and velocity distributions of matter are sel[-similar."," \citet{MacMillan} obtained a relation between $M_{BH,f}$ and $\sigma$ by assuming that the density and velocity distributions of matter are self-similar."
+ They assumed that the ealaxy is formed by (he extended. collapse of a halo composed of collisionless matter., They assumed that the galaxy is formed by the extended collapse of a halo composed of collisionless matter.
+ The central blackhole is grown proportionally to the halo as matter continues to fall in., The central blackhole is grown proportionally to the halo as matter continues to fall in.
+" The relation is given by (AlacMillanandIlenriksen2002) where 6 and o are scales in space and Gime respectively, and their ratio is related {ο the power-law index of the initial density perturbation e in the spherical infall model of halo erowth (IHenriksenandWidrow1999): The power-law index e=(7 +3)/2. where » is the index of the primordial matter power spectrum. P(h)xA""."," The relation is given by \citep{MacMillan}
+ where $\delta$ and $\alpha$ are scales in space and time respectively, and their ratio is related to the power-law index of the initial density perturbation $\epsilon$ in the spherical infall model of halo growth \citep{Henriksen}: The power-law index $\epsilon=(n+3)/2$ , where $n$ is the index of the primordial matter power spectrum $P(k) \propto k^n$."
+ Won=—2. Eq. (," If $n=-2$, Eq. ("
+8) agrees with the observation Mpyjryxol.,"8) agrees with the observation $M_{BH,f} \propto \sigma^4$."
+ This model involves a relation Eq. (, This model involves a relation Eq. (
+9) which is quite model dependent.,9) which is quite model dependent.
+" Adamsοἱal.(2001). assume the dark matter ancl barvons to be unsegregated. and (he isothermal initial mass density. distribution (CM,;(r)x r).", \citet{Adams} assume the dark matter and baryons to be unsegregated and the isothermal initial mass density distribution $M_t(r) \propto r$ ).
+ The specilic orbital energy. is conserved when(he particles fall into the small seed blackhole:, The specific orbital energy is conserved whenthe particles fall into the small seed blackhole:
+al. (,al. (
+1985) with different results. due to the revision of the isochroues as well as the calibration of abundances.,"1985) with different results, due to the revision of the isochrones as well as the calibration of abundances."
+ A more accurate AMIR was obtained by Edvardsson et al. (, A more accurate AMR was obtained by Edvardsson et al. (
+1993).,1993).
+ They have derived. abundances of 13 differcut elements. such as O. Fe. Si. Ba ete.," They have derived abundances of 13 different elements, such as O, Fe, Si, Ba etc.,"
+ as well as individual photometric ages. for 189 nearby field F aud € dwarts.," as well as individual photometric ages, for $189$ nearby field F and G dwarfs."
+ Their abundance analysis was made with theoretical LTE model atinospheres. based on the extensive hieh resolution. high S/N. spectroscopic observations of carefully selected field stars.," Their abundance analysis was made with theoretical LTE model atmospheres, based on the extensive high resolution, high S/N, spectroscopic observations of carefully selected field stars."
+ The resulted AMIR of Edvardssou et al. (, The resulted AMR of Edvardsson et al. (
+1993) was used at the present. study.,1993) was used at the present study.
+ llowever. this AMIR docs not constitute a tight constraints of the chemical model. since there is a considerable scatter.," However, this AMR does not constitute a tight constraints of the chemical model, since there is a considerable scatter."
+ Moreover. the results of the survey of Edvardssonu et al. (," Moreover, the results of the survey of Edvardsson et al. ("
+1993) concerning O vs. Fe relationships for field stars ire used in this study.,1993) concerning O vs. Fe relationships for field stars are used in this study.
+ As for metal-poor stars. the correlation between |O/Fe| aud |Fe/II] are taken from: Darbuy (1985).," As for metal-poor stars, the correlation between [O/Fe] and [Fe/H] are taken from Barbuy (1988)."
+ The radial Galactic profiles of atomic aud molecular livdrogen are discussed in Lacey Fall (1985)., The radial Galactic profiles of atomic and molecular hydrogen are discussed in Lacey Fall (1985).
+ An updated discussion is given in Dane (1993)., An updated discussion is given in Dame (1993).
+ Inside the solar circle. the molecular aud atomic gas are found in roughly equal amounts.," Inside the solar circle, the molecular and atomic gas are found in roughly equal amounts."
+" However. the surface density of atomic hydrogen. which ποσα» to be coustaut frou, I&pe to Ape. dominates the gas profiles outside the solar circle."," However, the surface density of atomic hydrogen, which seems to be constant from $4kpc$ to $15kpc$, dominates the gas profiles outside the solar circle."
+ The radial distribution of the suu of atomic aud molecular bydrogen given im Dame (1993) is adopted im this paper., The radial distribution of the sum of atomic and molecular hydrogen given in Dame (1993) is adopted in this paper.
+ The radial distribution of the preseut SER in the Galaxy are taken from Gusten Mereer (1983). Lyue cet al (1983) and ποτ ct al (1978).," The radial distribution of the present SFR in the Galaxy are taken from Gusten Merger (1983), Lyne et al (1983) and Guibert et al (1978)."
+ Data are based on several tracers of star formation: Lxiuan coutiuuun plotous from ΠΠ regions. pulsars aud supernova remnants.," Data are based on several tracers of star formation: Lyman continuum photons from HII regions, pulsars and supernova remnants."
+ It is normalized to the preseut SER in the solar neighbourhood (as in Lacey Fall 1985) since the absolute values depend ou poorly kuown couversion factors., It is normalized to the present SFR in the solar neighbourhood (as in Lacey Fall 1985) since the absolute values depend on poorly known conversion factors.
+ It is assumed that the Calactic disk is sheet-like. which originates aud grows ouly frou the infall of primordial gas.," It is assumed that the Galactic disk is sheet-like, which originates and grows only from the infall of primordial gas."
+ The disk is considered as a system of independent rings with Lhpe wide for cach., The disk is considered as a system of independent rings with $1kpc$ wide for each.
+ No radial inflows or outflows are considered aud the ceuter of cach ring locates at its median Galactoceutzic radius., No radial inflows or outflows are considered and the center of each ring locates at its median Galactocentric radius.
+" The ring ceutered at Galactocentric distauce r, = 8.5 kpe is labeled as the solar ποσαοσους.", The ring centered at Galactocentric distance $r_\odot$ = 8.5 kpc is labeled as the solar neighbourhood.
+ The age of the disk is adopted to be L3.0Carr (Rana 1991)., The age of the disk is adopted to be 13.0Gyr (Rana 1991).
+ The iustautaueous-recveliug approximation (RA) is relaxed. but instantaneous mixing of the eas with the stellar ejecta is assunied. be. the gas is characterized by a unique composition at cach epoch of tine.," The instantaneous-recycling approximation (IRA) is relaxed, but instantaneous mixing of the gas with the stellar ejecta is assumed, i.e., the gas is characterized by a unique composition at each epoch of time."
+ We solve ummerically the classical set of equations of, We solve numerically the classical set of equations of
+Foundation.,Foundation.
+ J.E.M. was supported by the CAC guest investigator program supported (his work through grant. ART-S003A. S.J.W. was supported by NASA contract NAS8-03060.2MASS..(IRAC/MIPS)..(ACIS).," J.E.M. was supported by the CXC guest investigator program supported this work through grant AR7-8003A. S.J.W. was supported by NASA contract NAS8-03060.,."
+Active Galactic Nuclei (AGN) emit enormous energies over the electro magnetic spectrum bands from radio through gamma ravs.,Active Galactic Nuclei (AGN) emit enormous energies over the electro magnetic spectrum bands from radio through gamma rays.
+ In. view of their short time variability. it has never been questioned that the power supply is primarilv eravitational.," In view of their short time variability, it has never been questioned that the power supply is primarily gravitational."
+ Phe shortest time scales variability is seen at hieh energv (Alellardy 1990. Grandi et al.," The shortest time scales variability is seen at high energy (McHardy 1990, Grandi et al."
+ 1992)., 1992).
+ Since a substantial variability cannot be observed on a time scale shorter than the light crossing time of the source (It« cot). the X-ray. variability is a powerful. probe of the innermost regions of AGN and can be used to constrain the emission mechanisms. the size of the emitting region and the elliciency of the matterfraciation conversion processes.," Since a substantial variability cannot be observed on a time scale shorter than the light crossing time of the source $<$ $\delta$ t), the X-ray variability is a powerful probe of the innermost regions of AGN and can be used to constrain the emission mechanisms, the size of the emitting region and the efficiency of the matter/radiation conversion processes."
+ Long duration. uninterrupted X-ray observations of AGN (hereafter with the term AGN we refer to QSO and Seyfert Galaxies. Le. AGN) were first carried out by EXOSAT (White and Peacock 1989). which showed that AGN are X-ray sources with lux changes over a wide range of time scales and that short termi X-ray. variability. clown toa few hundred seconds. is common (Mellardy 1990).," Long duration, uninterrupted X-ray observations of AGN (hereafter with the term AGN we refer to QSO and Seyfert Galaxies, i.e. AGN) were first carried out by EXOSAT (White and Peacock 1989), which showed that AGN are X-ray sources with flux changes over a wide range of time scales and that short term X-ray variability, down to a few hundred seconds, is common (McHardy 1990)."
+ A reanalysis of the EEXOSAT. data. (Grandi et al., A reanalysis of the EXOSAT data (Grandi et al.
+ 1992) showed that ~40% of AGN show variability on time scales less than one day., 1992) showed that $\sim$ of AGN show variability on time scales less than one day.
+ On longer time scales (typically weeks to months) of the same sample showed: variability. sugeesting than the long time variability is much more common (see also Mushotzky. Done Pouncs 1993).," On longer time scales (typically weeks to months) of the same sample showed variability, suggesting than the long time variability is much more common (see also Mushotzky, Done Pounds 1993)."
+ The EXOSAT. and. Ginga observations have also made evident that spectral shape changes can accompany unminosity variations., The EXOSAT and $Ginga$ observations have also made evident that spectral shape changes can accompany luminosity variations.
+ Although the prevailing trend appears o be a softening of the spectrum with increasing intensity. AGN do not show a unique spectral behavior (Cirandi et al.," Although the prevailing trend appears to be a softening of the spectrum with increasing intensity, AGN do not show a unique spectral behavior (Grandi et al."
+ 1992)., 1992).
+ This is in contrast to the results obtained for DL Lac objects. a particular class of AGN which show veculiaw characteristics as a featurcless optical spectrum. ugh polarization and rapid. variability.," This is in contrast to the results obtained for BL Lac objects, a particular class of AGN which show peculiar characteristics as a featureless optical spectrum, high polarization and rapid variability."
+ Using the EXNOSAT data. Ciüommi et al. (," Using the EXOSAT data, Giommi et al. ("
+1990) have found that BL Lac show a svsteniatic hardening of the spectrum. as the sources xightens.,1990) have found that BL Lac show a systematic hardening of the spectrum as the sources brightens.
+ Recent ROSA observations have confirmed that the overall picture of the spectral variability is. still. rather confused., Recent $ROSAT$ observations have confirmed that the overall picture of the spectral variability is still rather confused.
+ For example Boller et al. (, For example Boller et al. (
+1996) using a sample of narrow-line Sevfert. 1. galaxies have found. that some objects (e.g. AINN 957) show a softening of the spectrum,1996) using a sample of narrow-line Seyfert 1 galaxies have found that some objects $e.g.$ MKN 957) show a softening of the spectrum
+magnitude in panel of Fig. 2..,magnitude in panel of Fig. \ref{LVVI}.
+ The same data for Fornax cluster earlv-tvpe chwarts have been taken from Mieskect(2007) ancl plotted in the same panel., The same data for Fornax cluster early-type dwarfs have been taken from \cite{mi07} and plotted in the same panel.
+ A linear fit to the Fornax cluster My? vs. E relation is overplotted (full red line)., A linear fit to the Fornax cluster $_V$ vs. $-$ I relation is overplotted (full red line).
+ The Perseus cdwarfs obviously. adhere quite closely to the color-magnitude relation (6ΛΗΛ) defined by the Fornax chvarls., The Perseus dwarfs obviously adhere quite closely to the color-magnitude relation (CMR) defined by the Fornax dwarfs.
+ In order to investigate the physical nature οἱ the CALR. we convert the Iron abundance Fe/1]. measured by Alichiclsenetal.(2007). for the DO5 sample of cluster and group cles. into a 1 color using an empirical color-metallicity relation calibrated for old. stellar. populations. such as globular clusters (Couture.Harris.Allwright1990:ισα&Whitmore 1998).," In order to investigate the physical nature of the CMR, we convert the Iron abundance [Fe/H], measured by \cite{m07} for the D05 sample of cluster and group dEs, into a $-$ I color using an empirical color-metallicity relation calibrated for old stellar populations, such as globular clusters \citep{cha90,kw98}."
+. his places the DO5 dwarls essentially along the extension of the Mieskeetal.(2007) CALR., This places the D05 dwarfs essentially along the extension of the \cite{mi07} CMR.
+ The same exercise can be done for the Local Croup dSphs. with the metallicity taken from the compilation ov Crebel.Gallagher.Uarbeck (2003).. with the same result: they end. up following the Mieskeetal.(2007) CAL.," The same exercise can be done for the Local Group dSphs, with the metallicity taken from the compilation by \cite{ggh03}, with the same result: they end up following the \cite{mi07} CMR."
+ This already suggests that the VoL CMB of carly-vpe dwarf galaxies is. in fact. a Luminositv-metallicity relation.," This already suggests that the $-$ I CMR of early-type dwarf galaxies is, in fact, a luminosity-metallicity relation."
+" As a further test. we convert the T, colors of he Antlia dSphs/dlZs. measured. by SmithCastelli (2005). into Leolors using empirical (CLP) ) Fe/LY and D relations as an intermediate step."," As a further test, we convert the $-$ $_1$ colors of the Antlia dSphs/dEs, measured by \cite{smc08}, into $-$ I colors using empirical $-$ $_1$ $-$ [Fe/H] and $-$ $-$ I) relations as an intermediate step."
+ This places he Antlia dSphs/cllés almost exactly on the extension of the CAIR of the Fornax dSphs., This places the Antlia dSphs/dEs almost exactly on the extension of the CMR of the Fornax dSphs.
+ Thus. it appears that. the. observed VoL ολ of carly tvpe galaxies. from cwarls to giants. is a Duminositv-metallicity relation of galaxies that have stopped forming stars sulliciently long ago for there being almost no age information left (see also SmithCastellietal.(2008). and references There is considerable uniformity in the photometric sroperties of carly-twpe galaxies. from cdwarls to giants.," Thus, it appears that the observed $-$ I CMR of early type galaxies, from dwarfs to giants, is a luminosity-metallicity relation of galaxies that have stopped forming stars sufficiently long ago for there being almost no age information left (see also \cite{smc08} and references There is considerable uniformity in the photometric properties of early-type galaxies, from dwarfs to giants."
+" ""hotometric: parameters. quantifying the structure and stellar populations of earlv-tvpe galaxies. such as the half-ight radius. It. the central surface brightness fov. the Séresic exponent p. and E color all correlate with galaxy uminosity over a range of more than 6 orders of magnitude in luminositv."," Photometric parameters quantifying the structure and stellar populations of early-type galaxies, such as the half-light radius, $_{\rm e}$ the central surface brightness $\mu_{0,V}$, the Sérrsic exponent $n$, and $-$ I color all correlate with galaxy luminosity over a range of more than 6 orders of magnitude in luminosity."
+ The scaling relations involving the Sérrsic »wanmeters. Contrary to. previous claims. do not keep a constant slope over the whole luminosity range.," The scaling relations involving the Sérrsic parameters, contrary to previous claims, do not keep a constant slope over the whole luminosity range."
+" The Sérrsic exponent m varies2a with. luminosity."" L as nxLQAUu forqua. ealaxies brighter than Aly=14 mae but scatters around a constant value within the range nz0.51.0 for fainter dSphs.", The Sérrsic exponent $n$ varies with luminosity $L$ as $n \propto L^{0.25 - 0.3}$ for galaxies brighter than $_V \approx -14$ mag but scatters around a constant value within the range $n \approx 0.5 - 1.0$ for fainter dSphs.
+ This is in agreement with the fact that the surface brightness profiles of dSphs can be well approximated: by ing profiles with a concentration in the range ez3.10., This is in agreement with the fact that the surface brightness profiles of dSphs can be well approximated by King profiles with a concentration in the range $c \approx 3 - 10$.
+ Central surface. brightness increases with luminosity until the formation of the very brightest. cored ellipticals.," Central surface brightness increases with luminosity until the formation of the very brightest, cored ellipticals."
+ At Mi=14 mag. the slope of the Afio. changes abruptis.," At $_V \approx -14$ mag, the slope of the $_V -
+\mu_{0,V}$ changes abruptly."
+ We have shown that the My vs. Lis essentially a metallicitv-Iuminosity relation of old stellar populations. keeping the same slope over the whole luminosity range investigated here.," We have shown that the $_V$ vs. $-$ I is essentially a metallicity-luminosity relation of old stellar populations, keeping the same slope over the whole luminosity range investigated here."
+ Clearly. the absolute magnitude Alyzm14 mag is not just an arbitrary. divide between dSphs ancl cles.," Clearly, the absolute magnitude $_V \approx -14$ mag is not just an arbitrary divide between dSphs and dEs."
+ The rather abrupt changes in the slopes of some of the photometric scaling relations suggest that below and above this luminosityv. dillerent physical processes dominate the evolution of carly type galaxies.," The rather abrupt changes in the slopes of some of the photometric scaling relations suggest that below and above this luminosity, different physical processes dominate the evolution of early type galaxies."
+ The near-indepencdence of these sealing relations with respect. to. environment and the physical dillerences between dSphs ancl cles will x: investigated. theoretically using N-bodyv/SPIH-mocdels in another paper in this series (De Iüjeke. Valeke. Conselice. Penny. Held. in. prep.).," The near-independence of these scaling relations with respect to environment and the physical differences between dSphs and dEs will be investigated theoretically using N-body/SPH-models in another paper in this series (De Rijcke, Valcke, Conselice, Penny, Held, in prep.)."
+ In. a sense. the divide. between des and Es. which has historically been placed at. Mizm19 mag seems more arbitrary since the behaviour of he basic parameters describing the shapes of the surface xiehtness profiles as a function of Luminosity (1.6. the Sérrsic xwameters) does not change.," In a sense, the divide between dEs and Es, which has historically been placed at $_V \approx
+-19$ mag seems more arbitrary since the behaviour of the basic parameters describing the shapes of the surface brightness profiles as a function of luminosity (i.e. the Sérrsic parameters) does not change."
+ Acomparison of numerical simulations with observed scaling relations suggests that the uminosity. dependence of the Sérrsic parameters is due to act that the effects of supernova feedback. become more important as galactic gravitational potential wells become, Acomparison of numerical simulations with observed scaling relations suggests that the luminosity dependence of the Sérrsic parameters is due to fact that the effects of supernova feedback become more important as galactic gravitational potential wells become
+nreversibly lost.,irreversibly lost.
+ Then. for à given mechanical forcing. a relevant question is whether a larger thermal diffusivity reduces or euliances the kinetic energy of the flow.," Then, for a given mechanical forcing, a relevant question is whether a larger thermal diffusivity reduces or enhances the kinetic energy of the flow."
+ We are lacking quantitative results especially ii nou-linear reginües to answer such a basic question and more eenerallv to understand the effect of thermal diffusivity i a stellar context., We are lacking quantitative results especially in non-linear regimes to answer such a basic question and more generally to understand the effect of thermal diffusivity in a stellar context.
+ This situation is partlv due to the difficulty to remoduce flows with realistic Praudtl uuubers either in laboratory experiments or by uuncerical siuulations., This situation is partly due to the difficulty to reproduce flows with realistic Prandtl numbers either in laboratory experiments or by numerical simulations.
+ The Prandtl umuber P.=vir which compares the kinematic viscosity v aud the thermal diffusivity αν varies between 105 and 10Ὁ within the sun whereas it is equal to A7 in the air.," The Prandtl number $P_r =\nu /\kappa$ which compares the kinematic viscosity $\nu$ and the thermal diffusivity $\kappa$, varies between $10^{-6}$ and $10^{-9}$ within the sun whereas it is equal to $0.7$ in the air."
+ Although: some fiuid like metal liquid may have simall Praudtl iunubers in laboratory conditious (D.=0.025 for iuercury. see for example Cioni ct al. 19973)).," Although some fluid like metal liquid may have small Prandtl numbers in laboratory conditions $P_r =0.025$ for mercury, see for example Cioni et al. \cite{som}) )),"
+ these values remains far from the stellaz case., these values remains far from the stellar case.
+ The severe πιοΊσα. limitation is explained by the huge separation between the time scales of viscous dissipation aud thermal diffusion., The severe numerical limitation is explained by the huge separation between the time scales of viscous dissipation and thermal diffusion.
+ The coniputation of both processes over a few dynamical times would require a prolibitive amount of computer time., The computation of both processes over a few dynamical times would require a prohibitive amount of computer time.
+ Iu this paper. we investigate the unit where the time scale characterizing the thermal exchanecs is πιο shorter han the time scale of the motions (the ratio between both iue scales defines the Pécelet ummber).," In this paper, we investigate the limit where the time scale characterizing the thermal exchanges is much shorter than the time scale of the motions (the ratio between both time scales defines the Pécclet number)."
+ In Sect., In Sect.
+ 2. an asvinptotic form of the eoverniug equations is derived iu he contest of the Boussinesq approximation.," 2, an asymptotic form of the governing equations is derived in the context of the Boussinesq approximation."
+ Evideuces hat these asvuiptotie equatious actually approximate he Bossuesq equations for small Pécclet umubers. are resented in Sect., Evidences that these asymptotic equations actually approximate the Boussinesq equations for small Pécclet numbers are presented in Sect.
+ 3., 3.
+ Then. in Sect.," Then, in Sect."
+ L the elementary oxopertiesi of the xinall-Pécclet-inniber equations are described. emphasizius their heoretical auc practical interests.," 4, the elementary properties of the small-Pécclet-number equations are described, emphasizing their theoretical and practical interests."
+ Finally. the relevance of this approximation in a stellar context is comunented in Sect.," Finally, the relevance of this approximation in a stellar context is commented in Sect."
+ 5., 5.
+ We restrict ourselves to a fluicl laver embedded in Hi uniform vertical eravity field and bounded by two rorizoutal plates., We restrict ourselves to a fluid layer embedded in an uniform vertical gravity field and bounded by two horizontal plates.
+ A imechanical forcing is assumed to dive motions which cau be deseribed by the Boussimesq approximation., A mechanical forcing is assumed to drive motions which can be described by the Boussinesq approximation.
+ We do not need to specify the forcing x the moment. we only assune that it introduces a velocity scale C..," We do not need to specify the forcing for the moment, we only assume that it introduces a velocity scale $U_*$."
+ The temperature is fixed on both ates so that a lincar diffusive profile denoted T(:) js established initially., The temperature is fixed on both plates so that a linear diffusive profile denoted $T^{i}(z)$ is established initially.
+" The dyiuauucal effect of the stable stratification is measured by the Briuut-Vaiisalla requency, NV=ου... where g denotes hne eravity acceleration. ο) is the thera expansion coetiicicnt. AT the temperature difference setween the upper aud lower plates aud 1, the distance separating he plates."," The dynamical effect of the stable stratification is measured by the Brunt-Väiisällä frequency, $N_* = \left(\beta g {\Delta T}_*/ L_* \right)^{1/2}$, where $g$ denotes the gravity acceleration, $\beta$ is the thermal expansion coefficient, ${\Delta T}_*$ the temperature difference between the upper and lower plates and $L_*$ the distance separating the plates."
+" Tn the context of the Boussinesq approximation. the covering non-dimensional equations read:where. u=ve,|ce,we, is the velocity vector. p the pressure aud Ονμι).=Torey.τ}T'(:) the temperature deviation from the initial teurperature profile."," In the context of the Boussinesq approximation, the governing non-dimensional equations read:where, ${\bf u} = u {\bf e}_{x} + v {\bf e}_{y} + w {\bf e}_{z}$ is the velocity vector, $p$ the pressure and $\theta(x,y,z) = 
+T(x,y,z) - T^{i}(z)$ the temperature deviation from the initial temperature profile."
+ The : axis refers to the vertical direction. while the «aud y axis refer to the horizoutal ones.," The $z$ axis refers to the vertical direction, while the $x$ and $y$ axis refer to the horizontal ones."
+ In the heat equation. the third term of the left haud side corresponds to the vertical advection of teniperature against the 10211 teiiperatire eraciont dT7(:)/d:.," In the heat equation, the third term of the left hand side corresponds to the vertical advection of temperature against the mean temperature gradient $d T^{i}(z)/dz$ ."
+ This eracdicut is equal to unitv in the dineusiouless unit., This gradient is equal to unity in the dimensionless unit.
+ To nou-dimenusionalize the equations we used the velocity scale C... the leugth scale L.. the dynamical time scale fp=L./U.. the pressure scale o9TUZ aud the temperature variation+ AT...," To non-dimensionalize the equations we used the velocity scale $U_*$, the length scale $L_*$, the dynamical time scale $t_{\rm D} = L_*/U_*$ , the pressure scale $\varrho_0 U_*^2$ and the temperature variation ${\Delta T}_*$."
+" The system is then governed by the Richardson uuuber. A; the Pécclet προς, {δι aud the Revuolds munhber. A... respectively defined as The Richardson πο is the square of the ratio between the dvuiauical time scale fp aud the buovaucy time scale fj=1/N.."," The system is then governed by the Richardson number, $R_ i$, the Pécclet number, $P_e $, and the Reynolds number, $R_e$, respectively defined as The Richardson number is the square of the ratio between the dynamical time scale $t_{\rm D}$ and the buoyancy time scale $t_{\rm B}= 1/N_*$."
+" The thermal diffusivity & appears iu the Péccle uber which compares the thermal diffusion time. scale. f,=Lz/sks2 with: the dvnamical: time: scale."," The thermal diffusivity $\kappa$ appears in the Pécclet number which compares the thermal diffusion time scale, $t_{\kappa} = L_*^2/ \kappa$ with the dynamical time scale."
+ The: Revuolds umber is the ratio between the viscous time scale LZD/ aud the ανασα.. time- scale., The Reynolds number is the ratio between the viscous time scale $L_*^2/ \nu$ and the dynamical time scale.
+ Iu the nuit of μια] Pécclet nuuber. we assiuue that the solutions u aud 0 of he Doussinesq equations behave like Tavlor series: Note that in the contest of the Doussinesq equation. the pressure is an iutermecdiate variable determined bx the inconmpressibilitv condition (3)).," In the limit of small Pécclet number, we assume that the solutions ${\bf u}$ and $\theta$ of the Boussinesq equations behave like Taylor series: Note that in the context of the Boussinesq equation, the pressure is an intermediate variable determined by the incompressibility condition \ref{eq:div1}) )."
+ By inserting these asviuptotic expansions iu the heat equation. we find at the zero order in P: Since the temperature remains fixedto its iuitfial value ou bothbounding plates. temperature deviations vanish ou both plates.," By inserting these asymptotic expansions in the heat equation, we find at the zero order in $P_e $ Since the temperature remains fixedto its initial value on bothbounding plates, temperature deviations vanish on both plates."
+ Then. Eq. (6))," Then, Eq. \ref{eq:temp0_N}) )"
+ implies, implies
+found the 3jm feature of solid [sO ice in their observations of the DI ejecta plume.,"found the $3\,\micron$ feature of solid $_2$ O ice in their observations of the DI ejecta plume."
+ Sublimation of volatiles within the ejecta seems to have lasted over (he course of a few davs (Sehleicheretal., Sublimation of volatiles within the ejecta seems to have lasted over the course of a few days \citep{Schleicher}.
+2006).. We thus consider ongoing sublimation of volatiles in the (hilly ageregales and assume that a few percent of volatiles in volume still remain on the surface of the aggregates. (, We thus consider ongoing sublimation of volatiles in the fluffy aggregates and assume that a few percent of volatiles in volume still remain on the surface of the aggregates. (
+For a comparison. we also consider volatile[ree fluffy: aggregates).,"For a comparison, we also consider volatile–free fluffy aggregates)."
+ We then derive (he absorption cross sections C5 for the compact and [hilly aggregates. using Mie (heory with the optical constants deduced. from: the AIG 1992)..," We then derive the absorption cross sections $C_{\lambda}$ for the compact and fluffy aggregates, using Mie theory with the optical constants deduced from the MG \citep[e.g.,][]{Mukai}."
+ The use of the MG for aggregates has been examined based on the discrete dipole approximation (DDA) (Ixolokolovaetal. 2007)., The use of the MG for aggregates has been examined based on the discrete dipole approximation (DDA) \citep{Kolokolova}. .
+". From values of Cy αἱ wavelengths À of 8.8. 10.5. and 12.4jm where STO-Ix. observed. we caleulate the color temperature as 7.= and (he silicate-feature strength as Sy=2€859,5/(CsudC15)—1."," From values of $C_{\lambda}$ at wavelengths $\lambda$ of $8.8$, $10.5$ , and $12.4\,\micron$ where STO-K observed, we calculate the color temperature as $T_{\rm c}=C_{8.8}/C_{12.4}$ and the silicate-feature strength as $S_{\rm si}=2C_{10.5}/(C_{8.8}+C_{12.4})-1$."
+ We also caleulate 1 ab A=Q.8jmir and at a phase angle of the DI event for these aggregates.," We also calculate $P_{\rm l}$ at $\lambda = 0.8\,\micron$ and at a phase angle of the DI event for these aggregates."
+ Since {ή does nol stronglv depend on the number of monomers. we perlorm numerical calculations with Clusters of 256 monomers by using the DDA and the TAIA," Since $P_{\rm l}$ does not strongly depend on the number of monomers, we perform numerical calculations with clusters of 256 monomers by using the DDA and the TMM."
+L All the results are listed in Table l., All the results are listed in Table 1.
+" We first interpret the observational data by STO-Ix. Table 1 shows that the compact aggregate has higher T. and lower 5,4 (han the volatilecoated. [hilly aggregate.", We first interpret the observational data by STO-K. Table 1 shows that the compact aggregate has higher $T_{\rm c}$ and lower $S_{\rm si}$ than the volatile–coated fluffy aggregate.
+ Impact cratering mechanism indicates that high-velocity ejecta are thrown from nearsurface regions and low-velocily ejecta contain primarily materials excavated [rom deeper regions 1989)..," Impact cratering mechanism indicates that high-velocity ejecta are thrown from near–surface regions and low-velocity ejecta contain primarily materials excavated from deeper regions \citep[e.g.,][]{Melosh}."
+" Thus. a natural expectation in (he standard model is (hat hieh-velocity ejecta in (he DI ejecta plume would have higher 7, and lower 5,; than the low-velocity ejecta."," Thus, a natural expectation in the standard model is that high-velocity ejecta in the DI ejecta plume would have higher $T_{\rm c}$ and lower $S_{\rm si}$ than the low-velocity ejecta."
+ This agrees with the observational data by STO-K that show the high-velocity (~180 ms 1) ejecta to have higher 7; and lower $4 than the low-velocity ejecta (~100 ms 4)., This agrees with the observational data by STO-K that show the high-velocity $\sim 180$ m $^{-1}$ ) ejecta to have higher $T_{\rm c}$ and lower $S_{\rm si}$ than the low-velocity ejecta $\sim 100$ m $^{-1}$ ).
+ We next compare the polarimetric observation by Furushoetal.(2007) with £4 in Table 1., We next compare the polarimetric observation by \citet{Furusho} with $P_{\rm l}$ in Table 1.
+ Table 1 shows that the both aggregates have a positive polarization and that 774 for the compact aggregates (Lhe highvelocity ejecta) is larger than (hat for the fIuffv aggregates (the lowvelocity ejecta)., Table 1 shows that the both aggregates have a positive polarization and that $P_{\rm l}$ for the compact aggregates (the high–velocity ejecta) is larger than that for the fluffy aggregates (the low–velocity ejecta).
+ This is consistent with the observation by Furushoetal.(2Q07) who showed that the high.velocity ejecta have a slightly higher 74 compared to the lowvelocity ejecta., This is consistent with the observation by \citet{Furusho} who showed that the high–velocity ejecta have a slightly higher $P_{\rm l}$ compared to the low–velocity ejecta.
+ The STO data at time /=3.5 hrs alter the DIevent are shown by STO-S as infrared speclva al various radial distances., The STO data at time $t=3.5$ hrs after the DIevent are shown by STO-S as infrared spectra at various radial distances.
+ We also interpret the STO-5 data by estimating the radial distributions of Z7( and 55 fromthe data. where 77=fox /fioy.," We also interpret the STO-S data by estimating the radial distributions of $T^{\prime}_{\rm c}$ and $S^{\prime}_{\rm si}$ fromthe data, where $T^{\prime}_{\rm c} = I_{8.8}/I_{12.4}$ ,"
+>30 level.where σ is the flux uncertainty.,"$>3 \sigma$ level,where $\sigma$ is the flux uncertainty."
+ At jan. 59 out of 150 SWPs. and 97 out of 11v SWOPs were detected � >3o level.," At $\,\mu$ m, 59 out of 150 SWPs, and 97 out of 117 SWOPs were detected at $>3 \sigma$ level."
+" For the non-detections at jm. upper limits were calculated as 307, if the observed flux in the aperture ο was negative. and Fs)""+307, if positive (Carpenteretal.2009)."," For the non-detections at $\,\mu$ m, upper limits were calculated as $3\sigma_{70}$ if the observed flux in the aperture $F^{\rm
+ obs}_{70}$ ) was negative, and $F^{\rm obs}_{70} + 3\sigma_{70}$ if positive \citep{carpenter2009}."
+".. A star has significant excess if Az,=(Fi—Fryf.L3 (Tab. 2)).", A star has significant excess if $\chi_{70}=(F^{\rm obs}_{70}-F^{\rm pred}_{70})/\sigma_{70}>3$ (Tab. \ref{tab:data}) ).
+ In the SWPs sample. three stars have significant excess αἱ jm. and 22 stars at jam. In the SWODs sample. one star shows significant Εμ excess. while 17 stars have significant excess al 70 jm. The names of stars with excess are listed in the notes to Tab. 2..," In the SWPs sample, three stars have significant excess at $\,\mu$ m, and 22 stars at $\,\mu$ m. In the SWOPs sample, one star shows significant $\,\mu$ m excess, while 17 stars have significant excess at $\,\mu$ m. The names of stars with excess are listed in the notes to Tab. \ref{tab:data}."
+ Figure 1. shows the histogram of observed-to-photospheric [lux ratios al 24 and jan (IBo;. Rey).," Figure \ref{fig:hist} shows the histogram of observed-to-photospheric flux ratios at 24 and $\,\mu$ m $_{24}$ , $_{70}$ )."
+ At both wavelengths the distribution peaks at Rw 1. and can be approximated with a Gaussian of &=Εκ αἱ jan (the error budget is dominated by the uncertainty of the absolute flux calibration) aud o=24% at jm (the photometry is confusion noise limited at this wavelength).," At both wavelengths the distribution peaks at $\,{\approx}\,$ 1, and can be approximated with a Gaussian of $\sigma=4\%$ at $\,\mu$ m (the error budget is dominated by the uncertainty of the absolute flux calibration) and $\sigma=24\%$ at $\,\mu$ m (the photometry is confusion noise limited at this wavelength)."
+ Out of the 39 stars with excess. shown with dark shades in Figure 1.. 28 were known in the literature to harbour debris disks.," Out of the 39 stars with excess, shown with dark shades in Figure \ref{fig:hist}, , 28 were known in the literature to harbour debris disks."
+ For the remaining 11 our work shows the first indication lor a debris disk (Tah. 3))., For the remaining 11 our work shows the first indication for a debris disk (Tab. \ref{tab:new}) ).
+ At jan. all observed stars were detected at several 0 level.," At $\,\mu$ m, all observed stars were detected at several $\sigma$ level."
+ Thus. the distribution of excesses can be compared with traditional (two-sample tests. which show that SWPs and SWOPs do not differ.," Thus, the distribution of excesses can be compared with traditional two-sample tests, which show that SWPs and SWOPs do not differ."
+ This is due to the [act that the overwhelming majoritv of stars exhibit pure photospheric emission at jm. The jan data set is censoreclthere are upper limitsand survival analvsis is required to compare the two distributions (e.g..Moro-Martinetal.," This is due to the fact that the overwhelming majority of stars exhibit pure photospheric emission at $\,\mu$ m. The $\,\mu$ m data set is censored—there are upper limits—and survival analysis is required to compare the two distributions \citep[e.g.,][]{moromartin}."
+2007).. The Meier (IXM) estimator gives the cumulative distribution of a statistical variable taking into account the upper limits (e.g..Feigelson&Nelson 1985).," The Kaplan-Meier (KM) estimator gives the cumulative distribution of a statistical variable taking into account the upper limits \citep[e.g.,][]{fn}. ."
+. To use survival analvsis. the censoring should not depend on the variable itself.," To use survival analysis, the censoring should not depend on the variable itself."
+ This condition can befulfilled by using the observed-to-photospheric [ux ratio (25; ancl f) as thevariable., This condition can befulfilled by using the observed-to-photospheric flux ratio $R_{24}$ and $R_{70}$ ) as thevariable.
+ In average 5WI?s are, In average SWPs are
+their paper.,their paper.
+ E am also grateful to Dr Ixurt. van cler Ilevden for his help and to the referee for suggestions., I am also grateful to Dr Kurt van der Heyden for his help and to the referee for suggestions.
+from 90° to 907. in the representative library of Gaussian. for every (0.q) combination.,"from $-90^\circ$ to $90^\circ$, in the representative library of Gaussian, for every $(\sigma,q')$ combination."
+ The NNLS fit to the photometry will be performed with a three-dimensional grid (o.d.c) of Gaussians.," The NNLS fit to the photometry will be performed with a three-dimensional grid $(\sigma,q',\psi)$ of Gaussians."
+ The only problem of this approach is that the memory requirements will increase considerably. but a 30 grid of Gaussians can already be solved with current averages PCs.," The only problem of this approach is that the memory requirements will increase considerably, but a $\times$ $\times$ 30 grid of Gaussians can already be solved with current averages PCs."
+ The complete source code implementing the linear and nonlinear MGE fitting algorithms. described in this paper. both with and without isophote twist. together with manual and usage examples. can be found atf httpz/Awww.strw.leidenuniv.nl/--meappell/idl.," The complete source code implementing the linear and nonlinear MGE fitting algorithms, described in this paper, both with and without isophote twist, together with manual and usage examples, can be found at $\sim$ mcappell/idl."
+ The MGE method is one of the few simple parametrizations that are general enough to reproduce the surface brightness of realistic multicomponent galaxies., The MGE method is one of the few simple parametrizations that are general enough to reproduce the surface brightness of realistic multicomponent galaxies.
+ In addition many dynamical and IL1otometrie quantities can be evaluated easily and accurately when the density is expressed in MGE form., In addition many dynamical and photometric quantities can be evaluated easily and accurately when the density is expressed in MGE form.
+ In this paper we have deseribed a simple yet powerful gorithm that reduces the process of generating an MGE fit to multiple galaxy images to a simple. fast and automatic task.," In this paper we have described a simple yet powerful algorithm that reduces the process of generating an MGE fit to multiple galaxy images to a simple, fast and automatic task."
+ We have provided examples of its practical use and have tested it by accurately reproducing the photometry of a relatively large sample of galaxies. both with and without isophote twists.," We have provided examples of its practical use and have tested it by accurately reproducing the photometry of a relatively large sample of galaxies, both with and without isophote twists."
+ We have also compared some of our fits with previously obtained photometric models., We have also compared some of our fits with previously obtained photometric models.
+ We have also described an alternative algorithm that. although currently less practical. due to the larger computing power requirements. is guaranteed to converge to the minimum \E solution within the accuracy imposed by an adopted grid of xurameters in the solution space.," We have also described an alternative algorithm that, although currently less practical, due to the larger computing power requirements, is guaranteed to converge to the minimum $\chi^2$ solution within the accuracy imposed by an adopted grid of parameters in the solution space."
+ These algorithms have been implemented in an IDL program hat can produce an MGE model starting from the observed images of a galaxy. requiring the user to only input the coordinates of he galaxy centre. the PA and a characteristic flattening.," These algorithms have been implemented in an IDL program that can produce an MGE model starting from the observed images of a galaxy, requiring the user to only input the coordinates of the galaxy centre, the PA and a characteristic flattening."
+ Multiple resolution images (e.g. ground-based and HST) can easily be titted together. in one single step.," Multiple resolution images (e.g., ground-based and HST) can easily be fitted together, in one single step."
+ The complete IDL source code implementing the algorithms described in this paper is made sublicly available., The complete IDL source code implementing the algorithms described in this paper is made publicly available.
+ It is a pleasure to thank Eric Emsellem for stimulating discussions that provided the motivation to undertake this project., It is a pleasure to thank Eric Emsellem for stimulating discussions that provided the motivation to undertake this project.
+ I wish to thank Tim de Zeeuw and Ellen Verolme for careful reading of the manuscript and very useful comments and suggestions., I wish to thank Tim de Zeeuw and Ellen Verolme for careful reading of the manuscript and very useful comments and suggestions.
+ I thank the anonymous referee for constructive comments., I thank the anonymous referee for constructive comments.
+ I am grateful to José Funes and Alessandro Pizzella for obtaining the CCD image shown in Fig., I am grateful to José Funes and Alessandro Pizzella for obtaining the CCD image shown in Fig.
+ 10 and to Claudia Searlata for providing the reduced version of it., \ref{fig:ngc4698_mge} and to Claudia Scarlata for providing the reduced version of it.
+ That image was obtained with the VATT: the Alice LLennon Telescope and the Thomas BBannan Astrophysics Facility., That image was obtained with the VATT: the Alice Lennon Telescope and the Thomas Bannan Astrophysics Facility.
+where we used the expausious of the electromagnetic Lagrangian given. in Tsai Erber (1975. Eq.,"where we used the expansions of the electromagnetic Lagrangian given in Tsai Erber (1975, Eq."
+ 38a-b)., 38a-b).
+ In. these expressious. a ds the fine-structure constatἐν Ey ds the plasma. frequency auc Ns is the electron deusity.," In these expressions, $\alpha$ is the fine-structure constant, $E_{\rm pl}$ is the plasma frequency and $N_e$ is the electron density."
+ Civeu the ellipticities of the norual modes of propagation. it is Sralglitforward to calculate the abso'ption aud scattering terms 1 equation (6).," Given the ellipticities of the normal modes of propagation, it is straightforward to calculate the absorption and scattering terms in equation (6)."
+ The quantity q determines direcly the polarizatious of the normal modes aid thus all the eements of do;;/d€. which involve he inoduli of he eycli* projectlons o‘the unit polarization vectors outo the coordinate axis wit la given inagnetic ield direction.," The quantity $q$ determines directly the polarizations of the normal modes and thus all the elements of $\scatt$, which involve the moduli of the cyclic projections of the unit polarization vectors onto the coordinate axis with a given magnetic field direction."
+ In Figure 2.. we plot the [-diagoual terws Of ασια) as aluiction οἱ. elec1Ο ¢eusity for cifTerei|photon energies ancl rectious of propagatlon.," In Figure \ref{Fig:opac}, we plot the off-diagonal terms of $\scatt$ as a function of electron density for differentphoton energies and directions of propagation."
+ Figure 2 sows a Feimarkaldle οιhancement in he o-diagonal terms o‘the scattering matrix near the vacuui resonance deusity., Figure \ref{Fig:opac} shows a remarkable enhancement in the off-diagonal terms of the scattering matrix near the vacuum resonance density.
+ These terms «lescrile propagaticÀ-mode chauging and can siguilicautly alter the spectrum aud angular clistribtion of raciatioi »xopagatiug througla magnetized mecjui., These terms describe propagation-mode changing and can significantly alter the spectrum and angular distribution of radiation propagating through a magnetized medium.
+ This is because they can couve ‘ta photoiu from a mode witha simall mean-f'ee path to one with a large meau-free path and vice versa., This is because they can convert a photon from a mode with a small mean-free path to one with a large mean-free path and vice versa.
+ As the different pates of Figure 2 show. he conversion probability aud the width of the resonance cle»euds on the ploon energy and directjon ol propagation.," As the different panels of Figure 2 show, the conversion probability and the width of the resonance depends on the photon energy and direction of propagation."
+ These rapidly-chanueineOm off-diagonalDm terius are clillicη do haudle nuirerically when modeliug the transport of radiation., These rapidly-changing off-diagonal terms are difficult to handle numerically when modeling the transport of radiation.
+ A number of different approaches have been employed to date. which we discuss in the next section.," A number of different approaches have been employed to date, which we discuss in the next section."
+ As discusse lin 833. vacuui1 polarization 1itroduces nar‘ow alid energy-cepeucdent resolαλλο iu the opacities aid cross-1node juteractlou teris.," As discussed in 3, vacuum polarization introduces narrow and energy-dependent resonances in the opacities and cross-mode interaction terms."
+ Given that the racliative δειilibrium calculations require ai euergy grid that extends over four o‘cers of magnitucle aid a dept1 grid that spa15 len orders of imagniude. resolviug hne resonance ον using an arbitrarilv large uuuber of grid »oints in these two variables is computatioually prohinitive.," Given that the radiative equilibrium calculations require an energy grid that extends over four orders of magnitude and a depth grid that spans ten orders of magnitude, resolving the resonance by using an arbitrarily large number of grid points in these two variables is computationally prohibitive."
+ Lusteact. he soluion requires new numerical methods. which can sample the 'esoniauce regioi with high acewacy.," Instead, the solution requires new numerical methods, which can sample the resonance region with high accuracy."
+ The acetracy of the solLLIOLIS depeuds ou a co‘rect calculation ofthe total odtical depth uer the resonance., The accuracy of the solutions depends on a correct calculation of the total optical depth under the resonance.
+ Here we iuroduce a new algorithin to overcome this problem hat involves samplie the resonance regiol with a very large 1tunber of points on an aixiliary. grid. (denoted by prije) in order to compute accurately the tota optical depth., Here we introduce a new algorithm to overcome this problem that involves sampling the resonance region with a very large number of points on an auxiliary grid (denoted by prime) in order to compute accurately the total optical depth.
+ We sinoot1 the recistributiou matrix elenents according to, We smooth the redistribution matrix elements according to =
+OI the comparisons from the literature. the input. physics of Baralleοἱal.(1998) are (he most similar to ows.,"Of the comparisons from the literature, the input physics of \citet{bar98} are the most similar to ours."
+ The input. plivsies of Baralleetal.(1995). differ [rom our relerence calculation bv their use of the Saumon.Chabrier.&VanHorn(1995) equation of state only. adoption of a mixing length parameter. a=1.0. aud a deeper 7—100 fitting point of the outer boundary. condition.," The input physics of \citet{bar98} differ from our reference calculation by their use of the \citet{sau95} equation of state only, adoption of a mixing length parameter, $\alpha=1.0$ , and a deeper $\tau$ =100 fitting point of the outer boundary condition."
+ We do not find any impact on the LDB ages as a result of using only the Saumon.Chabrier.&VanLorn(1995). equation of state (see section ??))., We do not find any impact on the LDB ages as a result of using only the \citet{sau95} equation of state (see Section \ref{eos}) ).
+ Our adoption of a mixing length parameter. a=1.0. results in older LDD ages (see Section 3)).," Our adoption of a mixing length parameter, $\alpha=1.0$ , results in older LDB ages (see Section \ref{mixlen}) )."
+ Our adoption of the 7=100 fitàng point results in older LDB ages for the voungest elusters. and slightly vounger LDB ages for older clusters (see Section ??)).," Our adoption of the $\tau=100$ fitting point results in older LDB ages for the youngest clusters, and slightly younger LDB ages for older clusters (see Section \ref{atmos}) )."
+ The caleulations of Baralleοἱal.(1993). deviate from our reference calculation in a manner that qualitatively resemblesour calculation adopting the 7100 fitting point.," The calculations of \citet{bar98}
+ deviate from our reference calculation in a manner that qualitatively resemblesour calculation adopting the $\tau=100$ fitting point."
+ However. the LDB ages of Daraffeetal.(1998) at log(L/ Li. )—-2.9 are 5-0 vounger (han our expectations based on the internal error-budget calculation (Figure 8)).," However, the LDB ages of \citet{bar98} at $\log$ $_{\odot}$ )=-2.9 are $\sigma$ younger than our expectations based on the internal error-budget calculation (Figure \ref{errfig}) )."
+ The input plivsies ofthe D'Antona&Mazzitelli(1994)models differ [rom our reference model in (heir use of the MIID equation of state (Dappenοἱal.—1988) lor densities. p-0.01  and the Magni&Mazzitelli(1979) equation of state otherwise. and a grav atmosphere.," The input physics of the \citet{dan94} models differ from our reference model in their use of the MHD equation of state \citep{dap88} for densities, $\rho <0.01$ $^{-3}$ and the \citet{mag79} equation of state otherwise, and a gray atmosphere."
+ The [our sets of models compared in Figure 7 are combinations of using either ihe Alexander&Ferguson.(1994) or Ixurucz.(1998). low temperature opacities and either {he mixing length convection theory or the Canuto&Mazzitelli(1991) convection theory., The four sets of models compared in Figure \ref{comp} are combinations of using either the \citet{ale94} or \citet{kur98} low temperature opacities and either the mixing length convection theory or the \citet{can91} convection theory.
+ Overall. the four models result in vounger LDB ages.," Overall, the four models result in younger LDB ages."
+ The long-dash-dot line shows the comparison (hat is closest to our reference caleulation input physics (Alexander&(1994) opacity and mixing length convection theory).," The long-dash-dot line shows the comparison that is closest to our reference calculation input physics \citet{ale94}
+ opacity and mixing length convection theory)."
+ Comparing (his model to our grav atmosphere caleulation (solid line - Figure 5)) reveals a remaining difference ofLOY%., Comparing this model to our gray atmosphere calculation (solid line - Figure \ref{atmosfig}) ) reveals a remaining difference of.
+. Ienoring implementation differences. we infer the impact of using (he Magni&Mazzitelli(1979) equation of state is al most1055.," Ignoring implementation differences, we infer the impact of using the \citet{mag79} equation of state is at most."
+.. The lone-dash-small-dash line shows the comparison where the only change with the previous moclel is (he Kuruez(1998). low temperature opacity., The long-dash-small-dash line shows the comparison where the only change with the previous model is the \citet{kur98} low temperature opacity.
+ The largest age difference occurs al log( L/L. )=-2.1. and (his mimics our result when using (he Απο(1998) opacity (dashed line - Figure 4)).," The largest age difference occurs at $\log$ $_{\odot}$ )=-2.1, and this mimics our result when using the \citet{kur98} opacity (dashed line - Figure \ref{opacfig}) )."
+ The long-clash line is ihe comparison lor the D'Antona&Mazzitelli(1994) model that uses the AlexanderFerguson(1994) opacity and Canuto&Mazzitelli(1991) convection theory., The long-dash line is the comparison for the \citet{dan94} model that uses the \citet{ale94} opacity and \citet{can91} convection theory.
+ The line shows the difference when using the lxurucz(1998) opacity and CantoMazzitelli(1991) convection theory., The short-dash-dot line shows the difference when using the \citet{kur98} opacity and \citet{can91} convection theory.
+ The published data for the previous set of input parametersend al 34= 0.1...," The published data for the previous set of input parametersend at $M=0.1
+M_{\odot}$ ."
+ The only change to the input. physics relevant tothe LDB age technique between D'Antona&Mazzitelli(1997). and D'Antona&Mazzitelli(1994) is ihe convection treatment (Venturaetal. 1998)..., The only change to the input physics relevant tothe LDB age technique between \citet{dan97} and \citet{dan94} is the convection treatment \citep{ven98}. .
+ As we see in the comparison to D'Antona (open circles). the alternative convection (treatment has Little impact on the LDB ages.," As we see in the comparison to \citet{dan97} (open circles), the alternative convection treatment has little impact on the LDB ages."
+à sort of. separate solution of the transport equation. for subrelativistic and relativistic particles. so that the two spectra must be somehow connected at pmeposteriori.,"a sort of separate solution of the transport equation for subrelativistic and relativistic particles, so that the two spectra must be somehow connected at $p\sim mc$."
+" In (Berezhko. Yelshin Wksenofontoy 1994: Berezhko. Ixsenofontov Yelshin 1995: Berezhko 1996) the elfects of the non-linear reaction of accelerated: particles on the maximum achievable ομοον were investigated, together with the ellects of gcometry."," In (Berezhko, Yelshin Ksenofontov 1994; Berezhko, Ksenofontov Yelshin 1995; Berezhko 1996) the effects of the non-linear reaction of accelerated particles on the maximum achievable energy were investigated, together with the effects of geometry."
+ The maximum energy. of the particles accelerated in supernova remnants in the presence of large acceleration efficiencies was also studied by Ptuskin Zirakashvili (2003a.b).," The maximum energy of the particles accelerated in supernova remnants in the presence of large acceleration efficiencies was also studied by Ptuskin Zirakashvili (2003a,b)."
+ The need for aprachead solution of the acceleration oblem in the non-linear regime was recognized by Derezhko Ellison (1999). where a simple analytical broken-power-law approximation of the non-linear spectra was »resentecd.," The need for a solution of the acceleration problem in the non-linear regime was recognized by Berezhko Ellison (1999), where a simple analytical broken-power-law approximation of the non-linear spectra was presented."
+ ltecentlv. some promising analvtical solutions of the xoblem of non-linear shock acceleration have appeared in he literature (Malkov 1997: Malkov. Diamond Voll 2000: Blasi 2002. 2004).," Recently, some promising analytical solutions of the problem of non-linear shock acceleration have appeared in the literature (Malkov 1997; Malkov, Diamond Völlk 2000; Blasi 2002, 2004)."
+ Blasi (2004) considered for the first time he important οσοι of seed. pre-existing particles in the acceleration region (the linear theory of this phenomenon was first stucied by Bell (1978))., Blasi (2004) considered for the first time the important effect of seed pre-existing particles in the acceleration region (the linear theory of this phenomenon was first studied by Bell (1978)).
+ In a recent work by Ixang Jones (2005) the seed particles were included in numerical simulations of the acceleration. process., In a recent work by Kang Jones (2005) the seed particles were included in numerical simulations of the acceleration process.
+ Sumerical simulations have been instrumental to identify the dramatic elfects of the particles reaction: they showed that even when the fraction of particles injected from he thermal gas is relatively small. the energy. channelled into these few particles can be an appreciable part of the kinetic energv of the unshocked fluid. making the test xwticle approach unsuitable.," Numerical simulations have been instrumental to identify the dramatic effects of the particles reaction: they showed that even when the fraction of particles injected from the thermal gas is relatively small, the energy channelled into these few particles can be an appreciable part of the kinetic energy of the unshocked fluid, making the test particle approach unsuitable."
+ Phe most visible effects of the reaction of the accelerated: particles on the shock appear in the spectrum of the accelerated. particles. which shows a »eculiar. hardeningὃν at the highest> energies.," The most visible effects of the reaction of the accelerated particles on the shock appear in the spectrum of the accelerated particles, which shows a peculiar hardening at the highest energies."
+C The analytical approaches reproduce well the basic features arising from nonlinear elfects in shock acceleration., The analytical approaches reproduce well the basic features arising from nonlinear effects in shock acceleration.
+ There is an important. point which is still lacking in the calculations of the non-linear particle acceleration at shock waves. namely the possible amplification of the backgrounc magnetic field. found in the numerical simulations by Lucek Bell (2000. 2000a) and Bell Lucek (2001) and recently described by Bell (2004).," There is an important point which is still lacking in the calculations of the non-linear particle acceleration at shock waves, namely the possible amplification of the background magnetic field, found in the numerical simulations by Lucek Bell (2000, 2000a) and Bell Lucek (2001) and recently described by Bell (2004)."
+ Phis effect is still ignored in al calculations of the reaction of cosmic ravs on the shock structure., This effect is still ignored in all calculations of the reaction of cosmic rays on the shock structure.
+ We will not include this effect. in the presen )per., We will not include this effect in the present paper.
+ Nonlinear effects in shock acceleration of therma xwticles result. in the appearance of multiple solutions in certain regions of the parameter space., Nonlinear effects in shock acceleration of thermal particles result in the appearance of multiple solutions in certain regions of the parameter space.
+ This phenomenon is very general and was found in both the two-Luicl mocdels (Drury Vollk 1980. 1981) and in the kinetic moclels (Malkov. 1997: Malkoy et al.," This phenomenon is very general and was found in both the two-fluid models (Drury Völlk 1980, 1981) and in the kinetic models (Malkov 1997; Malkov et al."
+ 2001: Blasi 2004)., 2001; Blasi 2004).
+ Monte Carlo approaches do not show multiple solutionga, Monte Carlo approaches do not show multiple solutions.
+ This behaviour resembles that of critical svstems. with a bifureation occurring when some threshold is reached in a given order parameter.," This behaviour resembles that of critical systems, with a bifurcation occurring when some threshold is reached in a given order parameter."
+ In the case of shock acceleration. it is not easy to find a wav of discriminating among the multiple solutions when they appear.," In the case of shock acceleration, it is not easy to find a way of discriminating among the multiple solutions when they appear."
+ ln (Mond Drury 1998). a two Iluid approach has been used to demonstrate hat when three solutions appear. the one with intermediate ellicieney for particle acceleration is unstable to corrugations in the shock structure ancl emission. of acoustic waves.," In (Mond Drury 1998), a two fluid approach has been used to demonstrate that when three solutions appear, the one with intermediate efficiency for particle acceleration is unstable to corrugations in the shock structure and emission of acoustic waves."
+ P[ausibilitvy arguments may be put forward. to justify that he system mace of the shock plus the accelerated particles may sit at the critical point (see for instance the paper » AMalkov. Diamond Χο (2000)). but we are not aware of any real proof that this is what happens.," Plausibility arguments may be put forward to justify that the system made of the shock plus the accelerated particles may sit at the critical point (see for instance the paper by Malkov, Diamond Völlk (2000)), but we are not aware of any real proof that this is what happens."
+ The physical parameters that plav a role in this approach to criticality are the maximum momentum achievable by the xwiicles in the acceleration process. the Alach number of 16 shock. and the injection elliciencv.. namely the raction of thermal particles crossing the shock that are accelerated. to nonthermal energies.," The physical parameters that play a role in this approach to criticality are the maximum momentum achievable by the particles in the acceleration process, the Mach number of the shock, and the injection efficiency, namely the fraction of thermal particles crossing the shock that are accelerated to nonthermal energies."
+ The last of them. the injection ellicienev. hides a crucial physics problem by itselt. and plavs an important role in establishing the level of shock modification.," The last of them, the injection efficiency, hides a crucial physics problem by itself, and plays an important role in establishing the level of shock modification."
+ This cllicicney parameter in reality is determined by the microphysics of the shock and. shoulc not be a free parameter of the problem., This efficiency parameter in reality is determined by the microphysics of the shock and should not be a free parameter of the problem.
+ Unfortunately. our poor knowledge of such microphysics. in particular for collisionless shocks. does not allow us to establish a clear zux universal connection between the injection ellicieney auc he macroscopic shock properties.," Unfortunately, our poor knowledge of such microphysics, in particular for collisionless shocks, does not allow us to establish a clear and universal connection between the injection efficiency and the macroscopic shock properties."
+ Put aside the possibility o have a fully self-consistent. picture of this phenomenon. one can try to achieve a phenomenological description of it.," Put aside the possibility to have a fully self-consistent picture of this phenomenon, one can try to achieve a phenomenological description of it."
+ kang. Jones CGieseler (2002) introduced a sort of weight unction to determine a return. probability of. particles in he downstream fluid to the upstream fluid. as a function of xwticle momentum.," Kang, Jones Gieseler (2002) introduced a sort of weight function to determine a return probability of particles in the downstream fluid to the upstream fluid, as a function of particle momentum."
+ Only sulliciently suprathermal particles can jump back to the upstream region and therefore take xw in the acceleration process., Only sufficiently suprathermal particles can jump back to the upstream region and therefore take part in the acceleration process.
+ Here we adopt an injection recipe which is similar to the model of Ixang et al. (, Here we adopt an injection recipe which is similar to the model of Kang et al. (
+2002) (see also previous papers bx. Malkov. (1998) and by Cieseler et al. (,2002) (see also previous papers by Malkov (1998) and by Gieseler et al. (
+20000) and implement it in the semi-analvtical approach of Blasi (2002. 2004).,"2000)) and implement it in the semi-analytical approach of Blasi (2002, 2004)."
+ We investigate ren the phenomenon of multiple solutions and show that 16 injection recipe dramatically reduces the appearance of jese situations., We investigate then the phenomenon of multiple solutions and show that the injection recipe dramatically reduces the appearance of these situations.
+ We also study in some detail the ellicienev or particle acceleration as a function of the Mach number of 1e shock and the maximum momentum of the accelerated articles., We also study in some detail the efficiency for particle acceleration as a function of the Mach number of the shock and the maximum momentum of the accelerated particles.
+ The paper is structured. as follows: in section 2. we ον describe the method. proposed. bv. Blasi (2002) for rw calculation of the spectum ancl pressure. of particles vecelerated at a modified shock., The paper is structured as follows: in section \ref{sec:nonlin} we briefly describe the method proposed by Blasi (2002) for the calculation of the spectum and pressure of particles accelerated at a modified shock.
+ We describe the appearance X multiple solutions in section 3.. and the comparison with yw method. of Alalkov (1997). in section 4..," We describe the appearance of multiple solutions in section \ref{sec:multiplesolutions}, and the comparison with the method of Malkov (1997) in section \ref{sec:comparisonmalkov}."
+ In. section 5 we introduce a recipe for the injection of particles [rom 1e thermal pool., In section \ref{sec:injectionrecipe} we introduce a recipe for the injection of particles from the thermal pool.
+ Vhis recipe is then used in section 6. to garow how the regions of parameter space where multiple solutions appear aredrastically reduced by the ποτοσαος injection., This recipe is then used in section \ref{sec:singlesolutions} to show how the regions of parameter space where multiple solutions appear are drastically reduced by the self-regulated injection.
+ In section 7 we discuss the ellicieney. of particle acceleration. at. modified shocks. anc stress the role of escape of particles from upstream infinity.," In section \ref{sec:escape}
+ we discuss the efficiency of particle acceleration at modified shocks, and stress the role of escape of particles from upstream infinity."
+ The consequences, The consequences
+Due to the turbulent energy and its dissipation in the shocked region. after (he radiation. the electrons are re-accelerated by the stochastic acceleration and emit the early X-ray fluxes.,"Due to the turbulent energy and its dissipation in the shocked region, after the gamma-ray radiation, the electrons are re-accelerated by the stochastic acceleration and emit the early X-ray fluxes."
+ Therefore. (he energy injection in the central engine is not required to reproduce (he emission in the shallow decay. phase.," Therefore, the energy injection in the central engine is not required to reproduce the emission in the shallow decay phase."
+ Our results manifest that the temporal feature in (he shallow decav phase represents the process of turbulent. dissipation., Our results manifest that the temporal feature in the shallow decay phase represents the process of turbulent dissipation.
+" The microphvsical parameters. £,. £g and z,. are varied with the (me in the shocked region."," The microphysical parameters, $\varepsilon_e$, $\varepsilon_B$ and $\varepsilon_t$, are varied with the time in the shocked region."
+" Given (he special values of p=2.0 and α=0.5 in the adiabatic case. our calculation can represent the (vpical observational flux as Fx|""."," Given the special values of $p=2.0$ and $a=0.5$ in the adiabatic case, our calculation can represent the typical observational flux as $F\propto t^{-0.5}$."
+ NIoreover. it is noted that the same shallow decav phase is also detected in the optical band as well (e.g..Masonοἱal.2006).. our simple interpretation is that the svnclirotron emission of X-ray band. and optical band may be original from the same shocked region with turbulent energv.," Moreover, it is noted that the same shallow decay phase is also detected in the optical band as well \citep{mason06}, our simple interpretation is that the synchrotron emission of X-ray band and optical band may be original from the same shocked region with turbulent energy."
+ Another advantage of our model is that the crisis of radiative ellicieney (Llovd-Ionning&Zhang2004:lokaοἱal.2006:Fan&PiranZhangetal.2007) is therefore dispelled without. anv additional assumption of ejection/ejecta from the central engine.," Another advantage of our model is that the crisis of radiative efficiency \citep{lloyd04,ioka06,fan06,zhang07} is therefore dispelled without any additional assumption of ejection/ejecta from the central engine."
+ From this point of view. we support (he internal shock pattern of standard fireball model.," From this point of view, we support the internal shock pattern of standard fireball model."
+ In this work. we assume /(5)x5? where p=2.2 as the typical value.," In this work, we assume $f(\gamma)\propto \gamma^{-p}$ where $p=2.2$ as the typical value."
+ The detailed caleulation of the energy distribution f(5) is not needed. because in our scenario (he timescale I acceleration is much smaller than Chat of turbulent energy dissipation.," The detailed calculation of the energy distribution $f(\gamma)$ is not needed, because in our scenario the timescale of acceleration is much smaller than that of turbulent energy dissipation."
+ Since the timescale XI acceleration is also smaller (han that of shock hyvdrodynanmices. therefore. during the phase of shallow decay. the index of spectrum does not change (Liang.Zhang&2007).," Since the timescale of acceleration is also smaller than that of shock hydrodynamics, therefore, during the phase of shallow decay, the index of spectrum does not change \citep{liang07}."
+". Generally. the varied. values of 5, and oy lead to the different hydrodvnamic evolution and emissions in the entire alterglow (Mao&Wang2001a.b)."," Generally, the varied values of $\varepsilon_e$ and $\varepsilon_B$ lead to the different hydrodynamic evolution and emissions in the entire afterglow \citep{mao01a,mao01b}."
+. In this Letter. the early Xara [luxes in (he shallow decay. phase are estimated either in (he adiabatic or radiative case (rough the turbulent process.," In this Letter, the early X-ray fluxes in the shallow decay phase are estimated either in the adiabatic or radiative case through the turbulent process."
+ The average slope of radiative case is deeper than that of adiabatie case., The average slope of radiative case is deeper than that of adiabatic case.
+ Compared with the ART observation. the hvdrocdvnamice evolution of adiabatie case could be the realized regime during the shallow decay. phase.," Compared with the XRT observation, the hydrodynamic evolution of adiabatic case could be the realized regime during the shallow decay phase."
+" Although it is hard to obtain the values of 5, ancl 25 in our scenario. we put the constraints of z,«1 and ep«1."," Although it is hard to obtain the values of $\varepsilon_e$ and $\varepsilon_B$ in our scenario, we put the constraints of $\varepsilon_e\ll 1$ and $\varepsilon_B\ll 1$."
+ More observational samples are particularly requested for the further investigations., More observational samples are particularly requested for the further investigations.
+ We thank the staff of Swift group in Astronomico di Brera (Merate). G. Ghisellini and Z. G. Dai lor the general discussion.," We thank the staff of Swift group in Astronomico di Brera (Merate), G. Ghisellini and Z. G. Dai for the general discussion."
+ This work is financially supported by the Chinese National science Fund 10673028., This work is financially supported by the Chinese National Science Fund 10673028.
+Type II supernovae (SNe) are classifiel by the presence of IH in their spectra. and sub-classes are based on (their observed light curves.,Type II supernovae (SNe) are classified by the presence of H in their spectra and sub-classes are based on their observed light curves.
+ Thev are most likely caused [rom, They are most likely caused from
+able to reduce the FWHM of the output image to half of its input value and to increase the peak intensity by a factor ~6.,able to reduce the FWHM of the output image to half of its input value and to increase the peak intensity by a factor $\sim6$.
+ Figure 13 illustrates that correction., Figure 13 illustrates that correction.
+" Since the lucky imaging technique is feasible with an input seeing FWHM of ~1” (Hormuthetal. this experiment demonstrates how our instrument 2008),,can improve the performance of a lucky imaging device."," Since the lucky imaging technique is feasible with an input seeing FWHM of $\sim$ 1"" \cite{hor2008}, this experiment demonstrates how our instrument can improve the performance of a lucky imaging device."
+" During the campaign of September 2008, the system was tested under better seeing conditions."," During the campaign of September 2008, the system was tested under better seeing conditions."
+" In the left panel of Figure 14, an image of the double star WDS01095+4795 (with a visual magnitude of 4.59 and 5.61) is shown, taken with SAOLIM, without applying any corrections."," In the left panel of Figure 14, an image of the double star WDS01095+4795 (with a visual magnitude of 4.59 and 5.61) is shown, taken with SAOLIM, without applying any corrections."
+ The natural seeing FWHM was 1.1” during the observations.," The natural seeing FWHM was 1.1"" during the observations."
+ The separation of the star is 0.4” hence it was unresolved and appeared as a single spot in the image.," The separation of the star is 0.4"" hence it was unresolved and appeared as a single spot in the image."
+ The right panel shows the same object observed with the real-time closed-loop active., The right panel shows the same object observed with the real-time closed-loop active.
+" The corrected FWHM was 0.32” and as a result of this, the double star is clearly resolved."," The corrected FWHM was 0.32"" and as a result of this, the double star is clearly resolved."
+ The frame rate of the loop was 300 Hz., The frame rate of the loop was 300 Hz.
+ This result is quite promising if the prototype is attached to a lucky imaging camera., This result is quite promising if the prototype is attached to a lucky imaging camera.
+" When combined with an AO system, Lucky Imaging selects the periods when the turbulence that the adaptive optics system must correct is reduced."," When combined with an AO system, Lucky Imaging selects the periods when the turbulence that the adaptive optics system must correct is reduced."
+" In these periods, lasting a small fraction of a second, the correction given by the AO system is sufficient to give excellent resolution with visible light."," In these periods, lasting a small fraction of a second, the correction given by the AO system is sufficient to give excellent resolution with visible light."
+ The Lucky Imaging system sums the images taken during the excellent periods to produce a final image with much higher resolution than is possible with a conventional long-exposure AO camera (Lawetal.2006).., The Lucky Imaging system sums the images taken during the excellent periods to produce a final image with much higher resolution than is possible with a conventional long-exposure AO camera \cite{law2006}.
+ (Schiuidt1959)., \citep{sch59}.
+. Te Xugu~JM N=ντε0.15 (Ixeunicutt-Schiidt (Lig).," $_2$ $\Sigma_{\rm SFR} \sim \Sigma_{\rm
+gas}^{N}$ $N = 1.4\pm 0.15$ \citep[Kennicutt-Schmidt law, hereafter KS
+law;][]{ken98b}."
+ Cao&Solomon(20," $\lir$ \\citep{gao04a,gao04b,nar05}."
+0la.) Aly) that if Πον (1-0). cuisson faithfully traces dense molecular core mass (above a cutoff luminosity of ζω>104 E.). then constant SFR per unit mass is a result of a deuse molecular clump comprising a “fiudamenutal unit” of star formation.," \citet{gao04a,gao04b} $_{\rm H2}$ \citet{wu05} that if HCN (1-0) emission faithfully traces dense molecular core mass (above a cutoff luminosity of $L_{\rm bol} > 10^{4.5} \: L_\sun$ ), then constant SFR per unit mass is a result of a dense molecular clump comprising a “fundamental unit” of star formation."
+ The observed extragalactic linear correlation ix a natural extension of the constant SER per unif mass observed toward dense molecular chumps in the Milkv Way (Plumeetal.1992:Shirley2003. 2007).," The observed extragalactic linear correlation is a natural extension of the constant SFR per unit mass observed toward dense molecular clumps in the Milky Way \citep{plu92,shi03,shi07}."
+. Iu this picture. ultraliuinous IR galaxies such as Arp 220 which lie5 ou the linear SER-IICN (1-0) relation simply coutain more cluster forming waits. aud a higher fraction of dense molecular gas.," In this picture, ultraluminous IR galaxies such as Arp 220 which lie on the linear SFR-HCN (1-0) relation simply contain more cluster forming units, and a higher fraction of dense molecular gas."
+ The interpretation outlined im Wuetal.(2005) predicts a linear relation between SER aud tracers of even higher critical (πο) deusitv than HCN (1-0)., The interpretation outlined in \citet{wu05} predicts a linear relation between SFR and tracers of even higher critical $n_{\rm crit}$ ) density than HCN (1-0).
+ However. recent theoretical models from Naravananetal.(2007.hereafterNOT) and drunholz&Thonp-son(2007.hereafterITO07). predict that the power-law index between the SER aud higher critical density tracers such as TICN (3-2) should in fact be below unity.," However, recent theoretical models from \citet[][hereafter
+N07]{nar07} and \citet[][hereafter KT07]{kru07} predict that the power-law index between the SFR and higher critical density tracers such as HCN (3-2) should in fact be below unity."
+ The physical explanation for this behavior is that iu svstenis with predominantly low density gas. enuüssion from biel critical deusitv lines originates in the extreme tails of the density distribution resulting iu a Ly07» relation with a slope ereater than unity.," The physical explanation for this behavior is that in systems with predominantly low density gas, emission from high critical density lines originates in the extreme tails of the density distribution resulting in a $_{\rm mol}$ $<$$n$$>$ relation with a slope greater than unity."
+" This effect drives a SER-L,,, relation with a slope below nity for tracers with higher vai than that of Πο (J=1-0).", This effect drives a $_{\rm mol}$ relation with a slope below unity for tracers with higher $n_{\rm crit}$ than that of HCN (J=1-0).
+ To test this prediction. we have measured the ICN (3-2) line Iununositv πιNor on! froin a sample of 30 ealaxics and compared our results with recently published. data from Caacia-Carpioctal.(2008.hereafterGCO7).," To test this prediction, we have measured the HCN (3-2) line luminosity $\lhcnjthree$ ) from a sample of 30 galaxies and compared our results with recently published data from \citet[][hereafter GC07]{graciacarpio2007}."
+. Both datasets show a — sslope that is siguificautly below uuity. in agreement witli the model predictions.," Both datasets show a - slope that is significantly below unity, in agreement with the model predictions."
+ Observations of ICN (3-2) (14=265.886131 Cz) were obtained from 2007 February through 2007 June, Observations of HCN (3-2) $\nu_{rest} = 265.886431 \: {\rm GHz}$ ) were obtained from 2007 February through 2007 June
+The Ixelvin-Lelmholtz (stl) instability can occur anywhere that has a velocity shear and. as à result. is an important instability. in almost any system involving Uuics.,"The Kelvin-Helmholtz (KH) instability can occur anywhere that has a velocity shear and, as a result, is an important instability in almost any system involving fluids."
+ The instability has been studied. in a variety. of astrophysical systems. from solar wines (Amerstorleret.al.2007:Bet-arinietal.2006:Llasceawact2004) and. pulsar winds (Buceiantini&DelZanna2006) to thermal Hares (Venter&Meintjes 2006).," The instability has been studied in a variety of astrophysical systems, from solar winds \citep{amerstorfer07, bettarini06, hasegawa04} and pulsar winds \citep{bucciantini06} to thermal flares \citep{venter06}."
+. Due to its ability to drive mixing and urbulence. the IKI instability has been considered. relevant in protoplanetary disks (Johansenetal.2006:GomezOstriker 2005).. accretion disks and magnetospheres uravan 2004).. and other jets and outllows (Baty&Ixep-»ens 2006).," Due to its ability to drive mixing and turbulence, the KH instability has been considered relevant in protoplanetary disks \citep{johansen06, gomez05}, accretion disks and magnetospheres \citep{li04}, and other jets and outflows \citep{baty06}."
+.. One environment in which the presence of turbulent energv is of particular interest is molecular clouds., One environment in which the presence of turbulent energy is of particular interest is molecular clouds.
+ One possible source of this turbulent energy. is the interaction of protostcllar jets with the surrounding cloud. as first proposed. by Norman&Silk(19080).," One possible source of this turbulent energy is the interaction of protostellar jets with the surrounding cloud, as first proposed by \citet{norman80}."
+" The transfer. of momentunm from a jet to the cloud. is. most likely to occur through so-called: ""prompt entrainment"" (c.g.Dyson 1984).", The transfer of momentum from a jet to the cloud is most likely to occur through so-called “prompt entrainment” \citep[e.g.][]{dyson84}.
+. As the supersonic protostellar jet. propagates into the surrounding cloud. it forms a bow shock.," As the supersonic protostellar jet propagates into the surrounding cloud, it forms a bow shock."
+ This shock accelerates molecular cloucl material. imparting momentunm to it.," This shock accelerates molecular cloud material, imparting momentum to it."
+ D is worth noting that the momentum. imparted. is relatively well ordered: a further. process. must occur to convert this momentum into turbulent motions., It is worth noting that the momentum imparted is relatively well ordered: a further process must occur to convert this momentum into turbulent motions.
+ Early stuclies of the KID instability in. protostellar jets were carried out using perturbative linear analysis., Early studies of the KH instability in protostellar jets were carried out using perturbative linear analysis.
+ These initial stuclies (asreviewbyBirkinshaw1991) ONAL the [uic equations uncer set-ups of various combinations of magnetic field and shear lavers. etc.," These initial studies \citep[as reviewed by][]{birkinshaw91} examine the fluid equations under set-ups of various combinations of magnetic field and shear layers, etc."
+ )v approaching the problem. using a simple mathematical treatment. the cllect," By approaching the problem using a simple mathematical treatment, the effect"
+"association between the SGR and the star cluster. which would immediately rule out an association with SNR G42.8+0.6,","association between the SGR and the star cluster, which would immediately rule out an association with SNR G42.8+0.6."
+ We conclude that there Is a significant likelihood of random alignment between SGR 1900-14. and G42.8+0.6. and that there is litthe convincing evidence in favor of a physical association given the complex nature of this part of the sky.," We conclude that there is a significant likelihood of random alignment between SGR 1900+14 and G42.8+0.6, and that there is little convincing evidence in favor of a physical association given the complex nature of this part of the sky."
+ SGR 1627-41 has an X-ray counterpart which shows it to be embedded in the complicated radio region CTB 33 (9., SGR 1627–41 has an X-ray counterpart which shows it to be embedded in the complicated radio region CTB 33 \cite{wkv+99}) ).
+ CTB 33 was originally classified as an ccomplex (). an interpretation confirmed by recent high-resolution observations which show this source to be a large region in which a compact SNR. G337.0-0.1. is embedded ().," CTB 33 was originally classified as an complex \cite{sg70c}) ), an interpretation confirmed by recent high-resolution observations which show this source to be a large region in which a compact SNR, G337.0–0.1, is embedded \cite{sggf97}) )."
+ The X-ray localization for SGR 1627-41 puts it ~30” outside the rim of the SNR ()., The X-ray localization for SGR 1627–41 puts it $\sim30''$ outside the rim of the SNR \cite{hsk+00}) ).
+ Smith (1999)) estimate the probability of a random alignment between the SNR and SGR to be ~S%.., Smith \nocite{sbl99}) ) estimate the probability of a random alignment between the SNR and SGR to be $\sim$.
+ A common distance of 11.0 kpe has been determined independently for SNR G337.0-0.1 0) and SGR 1627-41 (9., A common distance of 11.0 kpc has been determined independently for SNR G337.0–0.1 \cite{fgr+96}) ) and SGR 1627–41 \cite{ccdd99}) ).
+ However. this simply indicates that the SNR and SGR are probably part of the same star-forming complex which includes CTB 33; it does not make a strong case that they correspond to the same supernova explosion.," However, this simply indicates that the SNR and SGR are probably part of the same star-forming complex which includes CTB 33; it does not make a strong case that they correspond to the same supernova explosion."
+ A possible fifth SGR. 1801-23. has been reported by Cline (2000)).," A possible fifth SGR, 1801–23, has been reported by Cline \nocite{cfg+00}) )."
+ The error box for this source is very elongated. extending across almost four degrees of the Galactic Plane in a complex region near the Galactic Center.," The error box for this source is very elongated, extending across almost four degrees of the Galactic Plane in a complex region near the Galactic Center."
+ Although this error circle passes through or near ~7 SNRs. any such error circle drawn randomly on the sky in this region would do so.," Although this error circle passes through or near $\sim$ 7 SNRs, any such error circle drawn randomly on the sky in this region would do so."
+ No serious case can be made for an association with any particular SNR until this SGR is confirmed and its position refined., No serious case can be made for an association with any particular SNR until this SGR is confirmed and its position refined.
+ In the second half of Table ?2.. we have listed the three SGRs which lie near SNRs.," In the second half of Table \ref{tab_snrs}, , we have listed the three SGRs which lie near SNRs."
+ As in Section 4.4. for the AXPs. for each system we have listed the age. distance and radius of the SNR. the offset of the SGR with respect to the SNR's center. the normalized offset. ./. and the implied transverse velocity. Vj.," As in Section \ref{sec_axp_overall} for the AXPs, for each system we have listed the age, distance and radius of the SNR, the offset of the SGR with respect to the SNR's center, the normalized offset, $\beta$, and the implied transverse velocity, $V_T$."
+" In stark distinction to the AXPs. for all three SGRs we find that ped, "," In stark distinction to the AXPs, for all three SGRs we find that $\beta \ge 1$."
+No estimate of an age or distance is available for but for SNRs N49 and G337.0-0.1. transverse velocities which are at the upper end of. or even beyond. the velocity distribution of the radio pulsar population (2: )) are inferred.," No estimate of an age or distance is available for but for SNRs N49 and G337.0–0.1, transverse velocities which are at the upper end of, or even beyond, the velocity distribution of the radio pulsar population \cite{lbh97}; \cite{cc98}) ) are inferred."
+ These velocities appear inconsistent with the low transverse velocities inferred for the AXPs., These velocities appear inconsistent with the low transverse velocities inferred for the AXPs.
+ The small statistics available thus suggest that if SGRs are associated with their nearby SNRs. then they must be a higher velocity population than are AXPs or radio pulsars.," The small statistics available thus suggest that if SGRs are associated with their nearby SNRs, then they must be a higher velocity population than are AXPs or radio pulsars."
+ Indeed it has been proposed that magnetars should have unusually high space velocities (>1000 s7!)) as a result of anisotropic neutrino emission immediately after core-collapse (:: ))., Indeed it has been proposed that magnetars should have unusually high space velocities $\ga1000$ ) as a result of anisotropic neutrino emission immediately after core-collapse \cite{dt92a}; \cite{td93a}) ).
+ However. these associations with SNRs rule. out any relationship or evolutionary connection with the AXPs.," However, these associations with SNRs rule out any relationship or evolutionary connection with the AXPs."
+ If SGRs are a younger population than the AXPs. as has been proposed by Kouveliotou ((1998)). then SGRs should generally have smaller values of ./ than do AXPs. which ts clearly not the case.," If SGRs are a younger population than the AXPs, as has been proposed by Kouveliotou \nocite{kds+98}) ), then SGRs should generally have smaller values of $\beta$ than do AXPs, which is clearly not the case."
+ On the other hand. if SGRs are an older incarnation of the AXPs (€: )). then the SGRs should have values of ./ larger than for the AXPs (as is observed). but the velocity distributions of the two populations should be Thus if all the associations in Table ??. are genuine. then AXPs and SGRs must represent two discrete populations of object.," On the other hand, if SGRs are an older incarnation of the AXPs \cite{ggv99}; \cite{gvd99}) ), then the SGRs should have values of $\beta$ larger than for the AXPs (as is observed), but the velocity distributions of the two populations should be Thus if all the associations in Table \ref{tab_snrs} are genuine, then AXPs and SGRs must represent two discrete populations of object."
+ This could result either if SGRs are magnetars and AXPs are accreting systems. or if there are two types of magnetar — apopulation with extreme (~10 G) magnetic fields and high velocities (the SGRs). and a separate group with lower magnetic fields (~1011 G) and lower velocities (the AXPs).," This could result either if SGRs are magnetars and AXPs are accreting systems, or if there are two types of magnetar — apopulation with extreme $\sim10^{15}$ G) magnetic fields and high velocities (the SGRs), and a separate group with lower magnetic fields $\sim10^{14}$ G) and lower velocities (the AXPs)."
+ Disregarding their values οἱ.) and Vj. AXPs and SGRs have remarkably similar properties eitekkm+00:: ο).," Disregarding their values of $\beta$ and $V_T$, AXPs and SGRs have remarkably similar properties \\cite{kkm+00}; \cite{kgc+01}) )."
+ There is thus considerable reluctance to conclude that the two populations are not related in some way., There is thus considerable reluctance to conclude that the two populations are not related in some way.
+ The only way this AXP/SGR connection can be maintained i5 if one abandons the associations between SGRs and SNRs. and therefore removes the discrepancy in the velocity distributions of the AXP and SGR populations.," The only way this AXP/SGR connection can be maintained is if one abandons the associations between SGRs and SNRs, and therefore removes the discrepancy in the velocity distributions of the AXP and SGR populations."
+ As argued earlier. none of the SGR/SNR associations are particularly compelling.," As argued earlier, none of the SGR/SNR associations are particularly compelling."
+ It is worth noting that of the claimed associations between radio pulsars and SNRs for which ./>|. almost all have been subsequently argued to be spurious citegj95en iu).," It is worth noting that of the claimed associations between radio pulsars and SNRs for which $\beta \ge 1$, almost all have been subsequently argued to be spurious \\cite{gj95c}; \cite{njk96}; \cite{kcm+98}; \cite{sgj99}) )."
+ If SGRs are neutron stars. but have ages of 50-100 kyr and have low space velocities. then we would expect their associated SNRs to have faded. but for these sources to still be near regions of supernova and star-forming activity. as is observed.," If SGRs are neutron stars, but have ages of 50–100 kyr and have low space velocities, then we would expect their associated SNRs to have faded, but for these sources to still be near regions of supernova and star-forming activity, as is observed."
+ In this case the the data are consistent with the hypothesis that SGRs and AXPs are related sources. but imply that the SGRs represent an older or longer-lived population whose SNRs have dissipated.," In this case the the data are consistent with the hypothesis that SGRs and AXPs are related sources, but imply that the SGRs represent an older or longer-lived population whose SNRs have dissipated."
+ Given this result. it is tempting to argue that AXPs evolve into SGRs.," Given this result, it is tempting to argue that AXPs evolve into SGRs."
+ However. this scenario ts problematic in that if AXPs are magnetars. then their youth and narrow range of spin-periods argue that their magnetic fields decay rapidly to strengths below ~[Οἱ G 0).," However, this scenario is problematic in that if AXPs are magnetars, then their youth and narrow range of spin-periods argue that their magnetic fields decay rapidly to strengths below $\sim10^{13}$ G \cite{cgp00}) )."
+ It is then hard to see how AXPs can evolve into SGRs. the latter of which are believed to have magnetic field strengths >10? G. Furthermore. the AXPs and SGRs have similar period distributions: 5-8 sec for SGRs. and 6-12 sec for AXPs.," It is then hard to see how AXPs can evolve into SGRs, the latter of which are believed to have magnetic field strengths $\ga10^{15}$ G. Furthermore, the AXPs and SGRs have similar period distributions: 5–8 sec for SGRs, and 6–12 sec for AXPs."
+ If SGRs are 5-10 times older than AXPs as proposed here. and if one extrapolates the steady spin-dow1 seen in several AXPs (:: )) to such ages. we would then expect the SGRs to have periods 510 s. which is not observed.," If SGRs are 5–10 times older than AXPs as proposed here, and if one extrapolates the steady spin-down seen in several AXPs \cite{gvd99}; \cite{kcs99}) ) to such ages, we would then expect the SGRs to have periods $\gg$ 10 s, which is not observed."
+ We note that some AXPs and SGRs do not show smooth spin-dowrt (G; )). while there is evidence in the case of 1E 22594586 for a short period of spin-up ()).," We note that some AXPs and SGRs do not show smooth spin-down \cite{kgc+01}; \cite{wkv+99b}) ), while there is evidence in the case of 1E 2259+586 for a short period of \cite{bs96}) )."
+ It has indeed been arguedthat magnetars can undergo periods of very low spin-down or evel spin-up depending on the level of internal activity ())., It has indeed been arguedthat magnetars can undergo periods of very low spin-down or even spin-up depending on the level of internal activity \cite{tdw+00}) ).
+ These effects must be significant if AXPs are to evolve into SGRs., These effects must be significant if AXPs are to evolve into SGRs.
+"to vary with the distance from the apex of the jet rà as Bor)2Bitriri)"" and nels.)2mírir)""s,7DU respectively. and rj=1pe.","to vary with the distance from the apex of the jet $r$ as $B(r)={B_1}(r/{r_1})^{-m}$ and }$ respectively, and ${r_1}=1~{\rm pc}$."
+ Wo the bulk motion velocity of the Jet is ve (corresponding to a Lorentz factor ? ) with an opening half-angle ó. and the axis of the jet makes an angle @ with the direction of the observer. the projection of the distance from the origin of the jet. A4. at which the optical depth to the svnchrotron self-absorption at the observing [requeney £. in the rest frame of the sources equals unity. is given by equation (3) in Wonniel (1981) as where ce(a) is the constant in 10. svnchrotron absorption coefficient. 6 is the Doppler factor. and 2n|m(2a13)2]/(2a 5).," If the bulk motion velocity of the jet is ${\beta }c$ (corresponding to a Lorentz factor $\gamma $ ) with an opening half-angle $\phi $, and the axis of the jet makes an angle $\theta $ with the direction of the observer, the projection of the distance from the origin of the jet, $l_{\rm core}$, at which the optical depth to the synchrotron self-absorption at the observing frequency $\nu_{\rm e}$ in the rest frame of the sources equals unity, is given by equation (3) in Könnigl (1981) as , where $c_{2}(\alpha)$ is the constant in the synchrotron absorption coefficient, $\delta$ is the Doppler factor, and $k_{m}=[2n+m(2\alpha+3)-2]
+/(2\alpha+5)$ ."
+ We assume that the mass loss rate of the jet is where ÀJ is the accretion rate., We assume that the mass loss rate of the jet is where $\dot M$ is the accretion rate.
+ We further assume that the magnetic field. pressure in the base of the jet can be sealed with the radiation pressure of the disc at radius rq: —o— We assume m=1. and mass conservation in the jet. i.c. n=2.," We further assume that the magnetic field pressure in the base of the jet can be scaled with the radiation pressure of the disc at radius $r_{\rm d}$: = We assume $m=1$, and mass conservation in the jet, i.e., $n=2$."
+ Substituting Eqs. (, Substituting Eqs. (
+4) and (5) into (3). we haveOL. ,"4) and (5) into (3), we have. . ] } } ),"
+where m is the cimensionless accretion rate in unit of rate.⊲⋠ and ry =20raion:3 this work.," where $\dot m$ is the dimensionless accretion rate in unit of Eddington rate, and ${\tilde r}_{\rm d}=r_{\rm d}/{\frac {2GM}{c^2}}$ ."
+"I5ddington we have adopted Ay,=1."," In this work, we have adopted $k_m=1$."
+ Ht is found that JiucxALIS for a=1., It is found that $l_{\rm core}\propto M^{0.64}$ for $\alpha=1$.
+ Ve note that the index is insensitive to the value ofa., We note that the index is insensitive to the value of $\alpha$.
+ Lhe statistic results in Figs., The statistic results in Figs.
+ 1 and 6 can therefore be interpreted approximately by the inhomogeneous jet. model., 1 and 6 can therefore be interpreted approximately by the inhomogeneous jet model.
+ In. fact. the lincar length. of the core fee is a function of several dillerent. physical quantities. which may be the reason that LT source?Plotted in Eig.," In fact, the linear length of the core $l_{\rm core}$ is a function of several different physical quantities, which may be the reason that the sources plotted in Fig."
+ 1 are dispersed over a large range., 1 are dispersed over a large range.
+" We find almost no correlation between the black hole mass Ady and redshift z in the restricted redshift range (0.4«2< 1). while a significant correlation is still present between the linear length of the core foo. and the black hole mass Mi, for this subsample."," We find almost no correlation between the black hole mass $M_{\rm bh}$ and redshift $z$ in the restricted redshift range $0.4<z<1$ ), while a significant correlation is still present between the linear length of the core $l_{\rm core}$ and the black hole mass $M_{\rm bh}$ for this subsample."
+ Phe possibility that the correlation we found is caused by the common redshift dependence can be ruled out., The possibility that the correlation we found is caused by the common redshift dependence can be ruled out.
+ Phe Spearman partial correlation analysis confirms this conclusion., The Spearman partial correlation analysis confirms this conclusion.
+ It is therefore an intrinsic correlation., It is therefore an intrinsic correlation.
+ We have to point out that the correlation between leave and. AM might be allected by the angular resolution imit of VLBI observations.," We have to point out that the correlation between $l_{\rm
+core}$ and $M_{\rm bh}$ might be affected by the angular resolution limit of VLBI observations."
+ The linear size of the core is obviously a function of redshift. which is similar to the situation for luminosity.," The linear size of the core is obviously a function of redshift, which is similar to the situation for luminosity."
+ I£ this is the case. this elfect will » important in the correlation analysis due to the common redshiftdependence.," If this is the case, this effect will be important in the correlation analysis due to the common redshiftdependence."
+ An analysis on the sourcesin a narrow restrictedredshift range can reduce this ellect. since similar incar resolutions would be available for all sources.," An analysis on the sourcesin a narrow restrictedredshift range can reduce this effect, since similar linear resolutions would be available for all sources."
+ Our, Our
+2001).,2001).
+" Some results indicate that the formation of lower mass PMS stars continues after the most massive stars have formed, although with varying efficiency, while other studies conclude that the normal sequence of formation is in fact from lower masses to higher masses (Ojha et al."," Some results indicate that the formation of lower mass PMS stars continues after the most massive stars have formed, although with varying efficiency, while other studies conclude that the normal sequence of formation is in fact from lower masses to higher masses (Ojha et al."
+ 2010)., 2010).
+ Some mentions are made in this context to the possibility that the opposite case would result in inhibition of the formation of low mass stars caused by radiative feedback from the more massive members (Price Bate 2009; Zinnecker Beuther 2008)., Some mentions are made in this context to the possibility that the opposite case would result in inhibition of the formation of low mass stars caused by radiative feedback from the more massive members (Price Bate 2009; Zinnecker Beuther 2008).
+" The determination of age spreads in real clusters are in any case influenced by the undetected presence of binaries, which can lead to a large overestimate of the spread (Weidner et al."," The determination of age spreads in real clusters are in any case influenced by the undetected presence of binaries, which can lead to a large overestimate of the spread (Weidner et al."
+" 2009) With our procedure we obtain age values with uncertainties for all assigned members, as explained in Sect."," 2009) With our procedure we obtain age values with uncertainties for all assigned members, as explained in Sect."
+ 2.2., 2.2.
+ The age of the PMS sequence in each cluster is then calculated as the average of all ages for the candidate members., The age of the PMS sequence in each cluster is then calculated as the average of all ages for the candidate members.
+" The rms deviation of this mean value, 2, is in principle indicative of the possible age spread."," The rms deviation of this mean value, $\Sigma$, is in principle indicative of the possible age spread."
+" The age values obtained from the four models used here (D97,P99,S00 and YO1) are included in the on-line catalogue."," The age values obtained from the four models used here (D97,P99,S00 and Y01) are included in the on-line catalogue."
+" These values are in good agreement with previous results on different YOCs (Park Sung 2002, Mamajek et al."," These values are in good agreement with previous results on different YOCs (Park Sung 2002, Mamajek et al."
+ 2002)., 2002).
+" These works indicate that ages from D97 are distinctly lower than those measured using the other three models, by amounts reaching several Myr."," These works indicate that ages from D97 are distinctly lower than those measured using the other three models, by amounts reaching several Myr."
+" On the other hand, differences between P99, SOO and YO1 models are at most of the order of 0.5 to 1 Myr but still systematic and significant."," On the other hand, differences between P99, S00 and Y01 models are at most of the order of 0.5 to 1 Myr but still systematic and significant."
+ The P99 models predict lower values than both S00 and YOI models., The P99 models predict lower values than both S00 and Y01 models.
+ These last two produce similar results., These last two produce similar results.
+" The average PMS ages of each cluster, obtained from S00 models, are listed in Table 2.."," The average PMS ages of each cluster, obtained from S00 models, are listed in Table \ref{t3}."
+" In reference to the age spreads, the use of independent membership determination from spectroscopy is especially needed to add the necessary constraints to the results from photometry and models."," In reference to the age spreads, the use of independent membership determination from spectroscopy is especially needed to add the necessary constraints to the results from photometry and models."
+" Without this membership confirmation, the standard deviation of the mean age, X, cannot be associated with a real age spread."," Without this membership confirmation, the standard deviation of the mean age, $\Sigma$, cannot be associated with a real age spread."
+" In spite of this shortcoming, a remarkable relation is found between the calculated standard deviation and the typical clustering scale, A, calculated as the average value of the distances between all pairs of members."," In spite of this shortcoming, a remarkable relation is found between the calculated standard deviation and the typical clustering scale, $\Delta$, calculated as the average value of the distances between all pairs of members."
+ This defines the spatial concentration of the population considered (Kaas et al., This defines the spatial concentration of the population considered (Kaas et al.
+ 2004)., 2004).
+ In Fig., In Fig.
+" 7 X is plotted versus this clustering scale, calculated in each cluster for all pairs of PMS members with mass above the approximate completeness limit defined in Sect."," \ref{AdvsC} $\Sigma$ is plotted versus this clustering scale, calculated in each cluster for all pairs of PMS members with mass above the approximate completeness limit defined in Sect."
+ 3.1.2., 3.1.2.
+" In the figure we overplot a line of slope 0.33, of the scaling relation between age dispersion and size proposed by Efremov Elmegreen (1998) for clusters and associations in a wider range of scales, both below and well above the spatial scale involved in our clusters."," In the figure we overplot a line of slope 0.33, of the scaling relation between age dispersion and size proposed by Efremov Elmegreen (1998) for clusters and associations in a wider range of scales, both below and well above the spatial scale involved in our clusters."
+" Three clusters do not follow the relation, due probably to undetected errors in the membership assignment."," Three clusters do not follow the relation, due probably to undetected errors in the membership assignment."
+" For the most outstanding case,18,, DAYO7 found signs of two clusters superimposed in the line of sight, of different ages and at well separated distances."," For the most outstanding case, DAY07 found signs of two clusters superimposed in the line of sight, of different ages and at well separated distances."
+ It is one of the clusters with the largest age uncertainty in the sample., It is one of the clusters with the largest age uncertainty in the sample.
+ These properties probably affect the calculated spreads., These properties probably affect the calculated spreads.
+" Excluding the three deviating clusters, a linear fit to the data provides a slope of 0.31, with a correlation coefficient of 0.8."," Excluding the three deviating clusters, a linear fit to the data provides a slope of 0.31, with a correlation coefficient of 0.8."
+ We consider that the agreement observed in Fig., We consider that the agreement observed in Fig.
+ 7 provides additional support to the ability of our procedure to detect PMS members in YOCs and to measure their physical properties., \ref{AdvsC} provides additional support to the ability of our procedure to detect PMS members in YOCs and to measure their physical properties.
+ The spatial distribution of cluster members is a classical matter of debate., The spatial distribution of cluster members is a classical matter of debate.
+" A recent general review on this topic is contained in Sánnchez Alfaro (2009), centred on the distribution of PMS members in YOCs as a primary source of information on the physical properties of the star formation process, both caused from the initial parameters of the parental cloud and from the influence of various environmental conditions."," A recent general review on this topic is contained in Sánnchez Alfaro (2009), centred on the distribution of PMS members in YOCs as a primary source of information on the physical properties of the star formation process, both caused from the initial parameters of the parental cloud and from the influence of various environmental conditions."
+ The issue is not settled., The issue is not settled.
+ Here we recall some related observational results., Here we recall some related observational results.
+" In a study of6383,, Fitzgerald et al. ("," In a study of, Fitzgerald et al. ("
+1978) found a ratio of 21 MS to 8 PMS members in the cluster core.,1978) found a ratio of 21 MS to 8 PMS members in the cluster core.
+" Baade (1983) found a weak concentration of PMS members towards cluster center in457,,7380,, and 1805,, together with the presence of extended halos, which could otherwise be just weak members, not necessarily PMS stars."," Baade (1983) found a weak concentration of PMS members towards cluster center in, and , together with the presence of extended halos, which could otherwise be just weak members, not necessarily PMS stars."
+" In a study ofIC348,, Lada Lada (1995) found a radial decay of both star density and subclustering, without specification of differences between MS and PMS stars."," In a study of, Lada Lada (1995) found a radial decay of both star density and subclustering, without specification of differences between MS and PMS stars."
+" In this same region, Herbig (1998) found a tendency of Weak-lined TTauri (WTT) stars to be more concentrated than Classical"," In this same region, Herbig (1998) found a tendency of Weak-lined TTauri (WTT) stars to be more concentrated than Classical"
+which used (he same zodiacal light model considered here (Wright1993).,which used the same zodiacal light model considered here \citep{wr98}.
+. In the J-band. we have weak limit on the 1.25 jim CIRD of 8.88 4 6.26 kJv +t.," In the J-band, we have weak limit on the 1.25 $\mu$ m CIRB of 8.88 $\pm$ 6.26 kJy $^{-1}$."
+ This is also consistent with the 1.25 sam values in Table 4.. which were reported in and Wright&Johnson(2001).. also using the same zodiacal light model used here.," This is also consistent with the 1.25 $\mu$ m values in Table \ref{oldvals}, , which were reported in \cite{elw01} and \citet{wrj01}, also using the same zodiacal light model used here."
+ The Kelsalletal.(1993) zodiacal light model gives à zodiacal intensity al the ecliptic pole 3.9. 9.2 and 4.0 kJv | lower at IX. J and L. Allowing for this difference in the models. these results are also consistent. with those reported in Canmbresyοἱal.(2001). and. (2000) αἱ 1.25 and 2.2 yan ancl Matsumotoetal.(2005) at 2.2 and 3.5 jun. also shown in Table 4..," The \citet{kel98} zodiacal light model gives a zodiacal intensity at the ecliptic pole 3.9, 9.2 and 4.0 kJy $^{-1}$ lower at K, J and L. Allowing for this difference in the models, these results are also consistent with those reported in \citet{cam01} and \citet{mat00} at 1.25 and 2.2 $\mu$ m and \citet{mat05} at 2.2 and 3.5 $\mu$ m, also shown in Table \ref{oldvals}."
+ The cumulative light from galaxies is a strict lower limit on the CIRD., The cumulative light from galaxies is a strict lower limit on the CIRB.
+ A determination ol the total contribution of resolved galaxies to the CIRD was determined via galaxy number counts by Fazioetal.(2004)., A determination of the total contribution of resolved galaxies to the CIRB was determined via galaxy number counts by \citet{faz04}.
+. Using the InfraRed Array Camera (IRAC) on the Spitzer Space Telescope (Eisenhardtetal.2004).. surveys were done of the Doóttes region. the Extended Groth Strip and a deep image surrounding the QSO IIS 17002-6416.," Using the InfraRed Array Camera (IRAC) on the Spitzer Space Telescope \citep{eis04}, surveys were done of the Boöttes region, the Extended Groth Strip and a deep image surrounding the QSO HS 1700+6416."
+ Alter integrating (he light [rom galaxies [rom the 10th to the 21st magnitudes. a total integrated. intensity of 5.4 nW 7? «sr1 (6.5 kJv 1) at 3.6 jim is reported. which is less than half of the CIRD determined at that wavelength by this and other similar studies.," After integrating the light from galaxies from the 10th to the 21st magnitudes, a total integrated intensity of 5.4 nW $^{-2}$ $^{-1}$ (6.5 kJy $^{-1}$ ) at 3.6 $\mu$ m is reported, which is less than half of the CIRB determined at that wavelength by this and other similar studies."
+ This 3.6 jam value is the final entry in Table 4.., This 3.6 $\mu$ m value is the final entry in Table \ref{oldvals}.
+ This discrepancy may be partially resolved by improving the photometry of these and other Spitzer survevs., This discrepancy may be partially resolved by improving the photometry of these and other Spitzer surveys.
+ Bul. the dominant source of error in (hese DIRDE minus PATASS measurements comes from the estimation of the zodiacal light. and based on the strong ecliptic dependence of our J-band results. it will likely take an improvement of the zodiacal light models to make sienilicant progress in determining whether we can resolve this discrepancy. with improvements in data analvsis. or if we truly require some exotic diffuse source as suggested by Cambresyetal.(2001).," But, the dominant source of error in these DIRBE minus 2MASS measurements comes from the estimation of the zodiacal light, and based on the strong ecliptic dependence of our J-band results, it will likely take an improvement of the zodiacal light models to make significant progress in determining whether we can resolve this discrepancy with improvements in data analysis, or if we truly require some exotic diffuse source as suggested by \citet{cam01}."
+. ]t is also interesting to note another indirect constraint on the CIRB from the attenuation ol TeV 5-ravs bv οe pair production through collisions with CIRB photons., It is also interesting to note another indirect constraint on the CIRB from the attenuation of TeV $\gamma$ -rays by $e^+e^-$ pair production through collisions with CIRB photons.
+ A recent attempt by Mapellietal.(2004). to fit the TeV spectrum of the blazar 1114264-428 found that the spectrum can not be fit using only (he integrated lisht [rom galaxies., A recent attempt by \citet{map04} to fit the TeV spectrum of the blazar H1426+428 found that the spectrum can not be fit using only the integrated light from galaxies.
+ Their best fil uses the Wright (2001) determination of the CIRB favoring the higher values determined from direct observations of the total sky brightness., Their best fit uses the Wright (2001) determination of the CIRB favoring the higher values determined from direct observations of the total sky brightness.
+ However. ILELS.S. observations of the blazars 112356-309 and LES 1101-232 by Aharonianοἱal.(2006) suggest that the CIRD can not be much higher than the lower limits [rom galaxy counts.," However, H.E.S.S. observations of the blazars H2356-309 and 1ES 1101-232 by \citet{aha06} suggest that the CIRB can not be much higher than the lower limits from galaxy counts."
+ The Wright(2001) 1.25 and 2.2 yan values and the Dwek&Arendt(19983).3.5 jam value were used by in fitting a model of the Extragalactie Dackground Light (EBL) needed for estimating -crav attenuation. resulüine in their P1.0 EBL model which is the top curve in Figure 10..," The \citet{elw01} 1.25 and 2.2 $\mu$ m values and the \citet{dwe98} 3.5 $\mu$ m value were used by \citet{aha06} in fitting a model of the Extragalactic Background Light (EBL) needed for estimating $\gamma$ -ray attenuation, resulting in their P1.0 EBL model which is the top curve in Figure \ref{p1.0}. ."
+ This modelgives (11.18.16) kJv | a( 1.25. 2.2 and 3.5 yan and assuming a power law," This modelgives (11,18,16) kJy $^{-1}$ at 1.25, 2.2 and 3.5 $\mu$ m and assuming a power law"
+iindex for their data set of o£1.17.,index for their data set of $\alpha \approx 1.17$.
+" Comparing this to Figure 1,, we see that these values correspond to an underlying Schmidt index of N— 1.5."," Comparing this to Figure \ref{figure:sfr_lmol}, we see that these values correspond to an underlying Schmidt index of $N=1.5$ ."
+ Figures 1 and 2 therefore provide evidence that an Schmidt index Nz1.5 holds for high-redshift galaxies., Figures \ref{figure:sfr_lmol} and \ref{figure:ks_molslope} therefore provide evidence that an Schmidt index $N \approx 1.5$ holds for high-redshift galaxies.
+" It is important to note, however, that there is an uncertainty of ~25% in these models."," It is important to note, however, that there is an uncertainty of $\sim 25\%$ in these models."
+ Observations of numerous CO transitions will help to narrow the uncertainty in the derived rrelation by providing additional constraints in the N—o space probed in Figure 2.., Observations of numerous CO transitions will help to narrow the uncertainty in the derived relation by providing additional constraints in the $N-\alpha$ space probed in Figure \ref{figure:ks_molslope}.
+" In Table 1,, we provide the mapping between Schmidt indices, N and molecular Kennicutt-Schmidt indices, o for four CO transitions."," In Table \ref{table:indices}, we provide the mapping between Schmidt indices, $N$ and molecular Kennicutt-Schmidt indices, $\alpha$ for four CO transitions."
+" These numbers provide enough information to create N—o plots similar to Figure 2 for transitions other than CO (J=3-2), and test the concept that a Schmidt index of N=1.5 describes star formation in ggalaxies."," These numbers provide enough information to create $N-\alpha$ plots similar to Figure \ref{figure:ks_molslope} for transitions other than CO (J=3-2), and test the concept that a Schmidt index of $N \approx 1.5$ describes star formation in galaxies."
+" In Figure 3,, we show the simulated pplot for our model with a Schmidt index of N—1.5, and denote the galaxies which fall within our selection criteria."," In Figure \ref{figure:ksplot}, we show the simulated plot for our model with a Schmidt index of $N=1.5$, and denote the galaxies which fall within our selection criteria."
+ The trends in Figure 1 can depend on how large of a dynamic range the mock observations span., The trends in Figure \ref{figure:sfr_lmol} can depend on how large of a dynamic range the mock observations span.
+" The trends in Figure 1 depend on the degree of differential excitation, which depends on the dynamic range of SFR surface densities (or gas surface densities) observed."," The trends in Figure \ref{figure:sfr_lmol} depend on the degree of differential excitation, which depends on the dynamic range of SFR surface densities (or gas surface densities) observed."
+" For example, consider the case where observations only probe a limited range of extremely high SFR galaxies (here, we take this to mean galaxies with gas fraction 0.4«f, 0.8)."," For example, consider the case where observations only probe a limited range of extremely high SFR galaxies (here, we take this to mean galaxies with gas fraction $0.4 < f_g < 0.8$ )."
+" In the high gas fraction/high SFR regime, the excitation conditions vary extremely rapidly."," In the high gas fraction/high SFR regime, the excitation conditions vary extremely rapidly."
+" The most gas-rich, densest systems have much of their gas thermalised, whereas lower gas fraction galaxies in this range (fg&0.4— 0.6) begin to contain subthermal gas."," The most gas-rich, densest systems have much of their gas thermalised, whereas lower gas fraction galaxies in this range $f_g \approx 0.4-0.6$ ) begin to contain subthermal gas."
+" While it may seem counterintuitive that some high gas-fraction galaxies may not be fully thermalised in higher CO lines, it is important to remember that the simulations (as do observations) consider global emission, from both low and high-density regimes."," While it may seem counterintuitive that some high gas-fraction galaxies may not be fully thermalised in higher CO lines, it is important to remember that the simulations (as do observations) consider emission, from both low and high-density regimes."
+" Subthermal CO level populations have been noted both in observations of local galaxies (Narayananetal.2005,2008a;Iono2009),, as well as in observations of even the densest, most heavily star-forming submillimetre-galaxies at high-redshift (Carillietal.2010;Har-risetal. 2010)."," Subthermal CO level populations have been noted both in observations of local galaxies \citep{nar05,nar08d,ion09}, as well as in observations of even the densest, most heavily star-forming submillimetre-galaxies at high-redshift \citep{car10, har10}."
+. The rapidly dropping CO excitation (from the higher J states to the ground states) at the highest SFRs causes the observed molecular KS relation to steepen from the underlying Schmidt relation., The rapidly dropping CO excitation (from the higher J states to the ground states) at the highest SFRs causes the observed molecular KS relation to steepen from the underlying Schmidt relation.
+" In the fiducial case of the N=1.5 galaxy, the observed iindex can range from ~1.5 to ~2.7 as we consider increasing numbers of galaxies with gas fractions beyond f,>0.4 in the fits."," In the fiducial case of the $N=1.5$ galaxy, the observed index can range from $\sim$ 1.5 to $\sim$ 2.7 as we consider increasing numbers of galaxies with gas fractions beyond $f_g > 0.4$ in the fits."
+" The molecular KS relation at the other end of the spectrum, in gas-poor galaxies, is less clear."," The molecular KS relation at the other end of the spectrum, in gas-poor galaxies, is less clear."
+" On one hand, we can expect that it would steepen owing to the bulk of the gas in the galaxy being below the SF threshold (see,e.g.Bigieletal.2008)."," On one hand, we can expect that it would steepen owing to the bulk of the gas in the galaxy being below the SF threshold \citep[see, e.g. ][]{big08}."
+". However, our models don't consider the possible destruction of molecular gas in low density environments (e.g.Krumholzetal.2009a)."," However, our models don't consider the possible destruction of molecular gas in low density environments \citep[e.g. ][]{kru09a}."
+". At low densities, the neutral ISM is a mixture of HI andH»."," At low densities, the neutral ISM is a mixture of HI and."
+. This is not captured by our models as we are forced by lack of resolution to consider the aas a fixed fraction of the neutral ISM mass., This is not captured by our models as we are forced by lack of resolution to consider the as a fixed fraction of the neutral ISM mass.
+" In reality, if the aand/or CO mass also drops at low densities, the observed KS relation may not become as steep as Figure 3 would suggest."," In reality, if the and/or CO mass also drops at low densities, the observed KS relation may not become as steep as Figure \ref{figure:ksplot} would suggest."
+" The trends predicted in Figure 1 are relatively robust within the physical parameter range chosen for the bulk of this paper (0.2«fj<0.4, with a typical dynamic range in SFR of order 10)."," The trends predicted in Figure \ref{figure:sfr_lmol} are relatively robust within the physical parameter range chosen for the bulk of this paper $0.2 < f_g < 0.4$, with a typical dynamic range in SFR of order $\sim$ 10)."
+" As was just shown, large deviations from galaxies of this sort may change the observed mapping from a Schmidt relation to a KS relation."," As was just shown, large deviations from galaxies of this sort may change the observed mapping from a Schmidt relation to a KS relation."
+" Thatsaid, our assumed range of gas fractions may accuratelyrepresent real 2 galaxies."," Thatsaid, our assumed range of gas fractions may accuratelyrepresent real 2 galaxies."
+Within the cold dark matter (CDM) paradigm for structure formation. first light in the Universe is usually assumed to have come from stars which collapsed early at the centres of rare anc unusually massive clark matter haloes associate with high peaks of the initial gaussian overcensity field.,"Within the cold dark matter (CDM) paradigm for structure formation, first light in the Universe is usually assumed to have come from stars which collapsed early at the centres of rare and unusually massive dark matter haloes associated with high peaks of the initial gaussian overdensity field."
+ Soon after their birth. these stars began to inlluence the structure. the thermodynamics and the chemical conten of surrounding gas.," Soon after their birth, these stars began to influence the structure, the thermodynamics and the chemical content of surrounding gas."
+ The epoch when this occurred. ai the precise details of how it happened. are not vet wel understood., The epoch when this occurred and the precise details of how it happened are not yet well understood.
+ ltecent simulations including gravitational. chemica and. radiative processes have suggested that metal-free eas in dark matter haloes of virial temperature νι~2000K and mass AL~LOPAL. cooled. etlicientlv. by. emission from molecular hydrogen and collapsed to form a star at recshilts Is«z30. (," Recent simulations including gravitational, chemical and radiative processes have suggested that metal-free gas in dark matter haloes of virial temperature ${\rm T_{vir} \sim 2000 K}$ and mass ${\rm M} \sim 10^6{\rm M_{\odot}}$ cooled efficiently by emission from molecular hydrogen and collapsed to form a star at redshifts $18<z<30$. ("
+Abel ct al.,Abel et al.
+ 1998. 2002: Fuller Couchman 2000. Bromim et al.," 1998, 2002; Fuller Couchman 2000, Bromm et al."
+ 2002: Yoshida et al., 2002; Yoshida et al.
+ 2003: for reviews see Dromnm Larson 2004 and Clover 2004)., 2003; for reviews see Bromm Larson 2004 and Glover 2004).
+ To achieve the required resolution. these caleulations followed evolution within very small cubic subvolumes of the Universe assuming periodic boundary conditions.," To achieve the required resolution, these calculations followed evolution within very small cubic subvolumes of the Universe assuming periodic boundary conditions."
+ Such boundary conditions suppress all Huctuations with wavelength longer than the side of the simulated region., Such boundary conditions suppress all fluctuations with wavelength longer than the side of the simulated region.
+ As we shall see. this can have drastic elfects. particularly on the morphology of large-scale structure and on the abundance anc formation history of objects with mass more than roughly 10.5 of the total mass in the simulated region.," As we shall see, this can have drastic effects, particularly on the morphology of large-scale structure and on the abundance and formation history of objects with mass more than roughly $10^{-4}$ of the total mass in the simulated region."
+" 1n fact. the first haloes of any given mass (and thus quite probably the first. stars) co not. form in. “typical” regions at all. but rather in ""protocluster"" regions (White Sprinecl 2000: Barkana Loch 2004)."," In fact, the first haloes of any given mass (and thus quite probably the first stars) do not form in “typical” regions at all, but rather in “protocluster” regions (White Springel 2000; Barkana Loeb 2004)."
+ Phis is an effect of he relatively large contribution from: Alpe wavelengths to he fluctuation amplitude on even the smallest mass scales. ancl is rellected. for example. in the plots presented by Mo White (2002) for the spatial clustering strength of haloes or redshifts out to z20.," This is an effect of the relatively large contribution from Mpc wavelengths to the fluctuation amplitude on even the smallest mass scales, and is reflected, for example, in the plots presented by Mo White (2002) for the spatial clustering strength of haloes for redshifts out to $z\sim 20$."
+" Conclusions about early objects reached from simulations of small ancl 7""tvpical regions may hus be misleading: in particular. the formation redshi of the first haloes of any given mass or temperature will xf systematically underestimated."," Conclusions about early objects reached from simulations of small and “typical” regions may thus be misleading; in particular, the formation redshift of the first haloes of any given mass or temperature will be systematically underestimated."
+ An alternative strategy is to simulate regions which are constrained. to have high overdensity or to contain a high amplitude density. peak, An alternative strategy is to simulate regions which are constrained to have high overdensity or to contain a high amplitude density peak
+"for demonstration purposes a simple unodel. with core radius rà, of aand 9=2/3. parameters that approximately deserihe the structure of the 22199 cluster of ealaxies.","for demonstration purposes a simple $\beta$ -model, with core radius $r_c$ of and $\beta=2/3$, parameters that approximately describe the structure of the 2199 cluster of galaxies."
+ This image is then convolved with the monochromatic iustruucutal PSF of the LECS., This image is then convolved with the monochromatic instrumental PSF of the LECS.
+ A sufficiently. accurate parameterization for the present purpose. based ou the publicly available calibration files. is giveu by where Fre) is the encircled energy fraction of the instimment.," A sufficiently accurate parameterization for the present purpose, based on the publicly available calibration files, is given by where $F(r)$ is the encircled energy fraction of the instrument."
+ For energies of 0.19. 0.277 and 0.93 keV the scale height & is 16.27. 11.70 and5.305. aud b dis 10.57. 13.32 aud 5.12. respectively.," For energies of 0.19, 0.277 and 0.93 keV the scale height $a$ is 16.27, 14.70 and, and $b$ is 10.57, 13.32 and 5.12, respectively."
+" Finally. we then couvolve the hieh-energv image with a Gaussian for which F(r)=1expt12/207), with e=L12'.. corresponding to the FWIIM of aas used by BB."," Finally, we then convolve the high-energy image with a Gaussian for which $F(r) = 1 - \exp
+(-r^2/2\sigma^2)$, with $\sigma=4.12$, corresponding to the FWHM of as used by BB."
+ Iu Figure 5.. we show the results of our analysis.," In Figure \ref{fig1}, we show the results of our analysis."
+ Tn all but oue case. we have chosen an effective cherey of 0.19 keV for the low energy baud (around he center of the soft 0.10.3 keV band of DB). and 1.93 keV for the high enuergv baud (around the effective center of the 0.5.2.2 keV baud of DD).," In all but one case, we have chosen an effective energy of 0.19 keV for the low energy band (around the center of the soft 0.1–0.3 keV band of BB), and 0.93 keV for the high energy band (around the effective center of the 0.5–2.2 keV band of BB)."
+ Iu all cases. we plotted the ratio of the soft to lard uid. for ai equal number of photous.," In all cases, we plotted the ratio of the soft to hard band, for an equal number of photons."
+ Case A in Figure 5. shows the profile ratio or an isothermal cluster. without convolving the Hel-cnerey band with a Gaussian. thus reflecting purely the instrmucutal PSF.," Case A in Figure \ref{fig1} shows the profile ratio for an isothermal cluster, without convolving the high-energy band with a Gaussian, thus reflecting purely the instrumental PSF."
+ The strong drop iu he ceuter demonstrates the effect of the broader PSF at low cuereics., The strong drop in the center demonstrates the effect of the broader PSF at low energies.
+ It is evident that even for a cluster without spectral variations. the radial oxofile ratio varies strongly aud is nowhere equal o unity.," It is evident that even for a cluster without spectral variations, the radial profile ratio varies strongly and is nowhere equal to unity."
+ Case D in Figure 5. is similar to case A. but row we convolved the high cncrev baud with a Caussian as BB did.," Case B in Figure \ref{fig1} is similar to case A, but now we convolved the high energy band with a Gaussian as BB did."
+ Again. for an isothermal cluster the ratio varies strongly as a function of r. mt now with opposite signatures compared with he previous case.," Again, for an isothermal cluster the ratio varies strongly as a function of $r$, but now with opposite signatures compared with the previous case."
+ Apparently. the lard bid has )een suoothed too much.," Apparently, the hard band has been smoothed too much."
+ Adoptiug the approach of BB. a soft flux deficit around wwould iudeed lave been interred. despite the fact hat the cluster is isothermal.," Adopting the approach of BB, a soft flux deficit around would indeed have been inferred, despite the fact that the cluster is isothermal."
+ Taking one step wither. if BB would have corrected their profile with a curve like case D. they would have found he soft excess reported by us!," Taking one step further, if BB would have corrected their profile with a curve like case B, they would have found the soft excess reported by us!"
+ Case C iu Figure 5. is similar to case B (ic. an isothermal cluster with the BB approach of sanootling the ligh-enerey baud). but now for a low-energv baud of 0.277 keV instead of 0.19 keV. The oversmoothius of the hiel-energy baa is more evident here.," Case C in Figure \ref{fig1} is similar to case B (i.e., an isothermal cluster with the BB approach of smoothing the high-energy band), but now for a low-energy band of 0.277 keV instead of 0.19 keV. The oversmoothing of the high-energy band is more evident here."
+ The comparison of curves D aud € shows that if radial profiles are use for this kind of analvsis. the spectral/spatia enerev distribution should be kuown aud iiodele appropriately.," The comparison of curves B and C shows that if radial profiles are used for this kind of analysis, the spectral/spatial energy distribution should be known and modeled appropriately."
+ This has not becu done by BB., This has not been done by BB.
+ Case D in Figure 5 simulates the effect of he presence of a cooling flow or an abunudauce eracdieut., Case D in Figure \ref{fig1} simulates the effect of the presence of a cooling flow or an abundance gradient.
+ It shows the theoretical harducss ratio in case the core radius of the lard component is V8 times smaller than that of the soft component., It shows the theoretical hardness ratio in case the core radius of the hard component is 0.8 times smaller than that of the soft component.
+ This mnimics. e.g..Mw] the case when the Fe-L complex (ün the hard baud) is ceutrally concentrated. jecause of either an abundance gradient or a cooling flow.," This mimics, e.g., the case when the Fe-L complex (in the hard band) is centrally concentrated, because of either an abundance gradient or a cooling flow."
+ The profiles have uot been convolved with the iustruuenut or the Gaussian., The profiles have not been convolved with the instrument or the Gaussian.
+" It also has an ~20 excess bevoud ~6"".. thereby iinuickius approxinatelv the soft excess as found by us iu A 2199."," It also has an $\sim$ 20 excess beyond $\sim$, thereby mimicking approximately the soft excess as found by us in A 2199."
+ Finally. case E in Figure 5. correspouds to the case presented in case D. but now the profiles are convolved with the mstruneut aud the lard baud convolved with the Gaussian.," Finally, case E in Figure \ref{fig1} corresponds to the case presented in case D, but now the profiles are convolved with the instrument and the hard band convolved with the Gaussian."
+ A comparison of curve E with D shows that the convolved profile is completely differcut from the original profile. and comparing curve E with cuve D shows that he differences due to a different cluster model are partially washed out. thus confirmune that the nethod of BB teuds to destroy ΠοιαΊο by oversinootline.," A comparison of curve E with D shows that the convolved profile is completely different from the original profile, and comparing curve E with curve B shows that the differences due to a different cluster model are partially washed out, thus confirming that the method of BB tends to destroy information by oversmoothing."
+ All these effects are completely neglected: by DD., All these effects are completely neglected by BB.
+ Thev simply compare their results with the expected scaled softuess ratio of 1 rather than he appropriate curve., They simply compare their results with the expected – scaled – softness ratio of 1 rather than the appropriate curve.
+ That curve can only be obtained by a full spatial/spectral analysis as. ¢.e.. done by our team.," That curve can only be obtained by a full spatial/spectral analysis as, e.g., done by our team."
+ But then the procedure of BB is obsolete. since our spectral modeling already gave the correct answers.," But then the procedure of BB is obsolete, since our spectral modeling already gave the correct answers."
+ We have shown in this Letter that the analysis method of BB for assessing the soft excess in A 2199 is inadequate and leads to erroneous conclusions regarding the presence of a soft excess., We have shown in this Letter that the analysis method of BB for assessing the soft excess in A 2199 is inadequate and leads to erroneous conclusions regarding the presence of a soft excess.
+ Iu fact. the method of DD is oversimplifiee," In fact, the method of BB is oversimplified"
+"planets reported in the literature todate’;; transmission spectra for HD 189733 (Swainetal.2008),, GJ-436 (Pontetal. 2009),, and XO-1 (Tinettietal.2010),, and dayside emission spectra for HD 189733 (Swainetal.2009b) and HD 209458 (Swainetal.2009a).","planets reported in the literature to; transmission spectra for HD 189733 \citep{Swain_2008}, GJ-436 \citep{Pont_2009}, , and XO-1 \citep{Tinetti_2010}, and dayside emission spectra for HD 189733 \citep{Swain_2009b} and HD 209458 \citep{Swain_2009a}."
+" The emission spectra are obtained in much the same way as the transmission spectra, by observing an eclipse of the planet, and using the same linear decorrelation methods to correct for the baseline function."," The emission spectra are obtained in much the same way as the transmission spectra, by observing an eclipse of the planet, and using the same linear decorrelation methods to correct for the baseline function."
+" It is clear that all five spectra show approximately the same peak-to-peak amplitude, and have wavelength dependance on similar scales."," It is clear that all five spectra show approximately the same peak-to-peak amplitude, and have wavelength dependance on similar scales."
+" This is perhaps surprising, given the planets have vastly different masses, radii and irradiation levels."," This is perhaps surprising, given the planets have vastly different masses, radii and irradiation levels."
+ Transmission and emission spectra also show the same amplitudes., Transmission and emission spectra also show the same amplitudes.
+" Whilst this is by no means evidence that the spectra are the result of systematic noise in the detector, and the similarity could be coincidental, it is nonetheless wise to explore this possibility."," Whilst this is by no means evidence that the spectra are the result of systematic noise in the detector, and the similarity could be coincidental, it is nonetheless wise to explore this possibility."
+" It is therefore important that the results reported from NICMOS transmission spectroscopy be re-analysed to confirm previous detections of molecules, in particular because the results are dependent on a complex decorrelation procedure, with an arguably model."," It is therefore important that the results reported from NICMOS transmission spectroscopy be re-analysed to confirm previous detections of molecules, in particular because the results are dependent on a complex decorrelation procedure, with an arguably model."
+" Here, we present a re-analysis of the three NICMOS transmission spectra shown in Fig. 1.."," Here, we present a re-analysis of the three NICMOS transmission spectra shown in Fig. \ref{fig:NICMOS_data}."
+" Sects. 2,,"," Sects. \ref{sect:HD189733},"
+" 3 and 4 describe the data reduction, analyses and results of the HD 189733, GJ-436, and XO-1 data, respectively."," \ref{sect:GJ436} and \ref{sect:XO1} describe the data reduction, analyses and results of the HD 189733, GJ-436, and XO-1 data, respectively."
+" Finally, Sect."," Finally, Sect."
+ summarises and presents our conclusions., \ref{sect:summary} summarises and presents our conclusions.
+" A transit of HD 189733 was monitored on 25 May 2007 with HST/NICMOS, using the G206 grism covering the wavelength range 1.4 — 2.5um."," A transit of HD 189733 was monitored on 25 May 2007 with HST/NICMOS, using the G206 grism covering the wavelength range 1.4 – 2.5."
+". Analyses of these data were first reported in Swainetal.(2008,hereafterS08)..", Analyses of these data were first reported in \citet[][hereafter S08]{Swain_2008}.
+" As HD 189733 is not in the continuous viewing zone of the HST, the transit was observed over five half-orbits (~48 minutes each), consisting of 638 spectra in total, all with exposure times of 1.624 seconds."," As HD 189733 is not in the continuous viewing zone of the HST, the transit was observed over five half-orbits $\sim48$ minutes each), consisting of 638 spectra in total, all with exposure times of 1.624 seconds."
+" The first, second, fourth and fifth orbits cover the out-of-transit part of the light curve and consist of 119, 128, 131 and 130 observations, respectively."," The first, second, fourth and fifth orbits cover the out-of-transit part of the light curve and consist of 119, 128, 131 and 130 observations, respectively."
+ These data are required to determine the photometric baseline., These data are required to determine the photometric baseline.
+" Only the third orbit was taken in-transit, and consists of the remaining 130 images."," Only the third orbit was taken in-transit, and consists of the remaining 130 images."
+" In addition to the spectra, some exposures were taken at the beginning of the first orbit to enable wavelength calibration."," In addition to the spectra, some exposures were taken at the beginning of the first orbit to enable wavelength calibration."
+" The calibrated images, one of which is displayed in Fig. 2,,"," The calibrated images, one of which is displayed in Fig. \ref{fig:HD189733_image},"
+ include all basic calibrations except for flat-fielding., include all basic calibrations except for flat-fielding.
+" Correctly flat fielding spectroscopic data requires taking into account the wavelength dispersion and position of the source, and therefore a different flat-field correction would be required at each point along the spectral trace of the target."," Correctly flat fielding spectroscopic data requires taking into account the wavelength dispersion and position of the source, and therefore a different flat-field correction would be required at each point along the spectral trace of the target."
+" S08 argue that flat-fielding is not required for differential spectroscopy, as each light curve will be normalised before measurement of parameters from the light curves."," S08 argue that flat-fielding is not required for differential spectroscopy, as each light curve will be normalised before measurement of parameters from the light curves."
+" However, strictly speaking flat-fielding is required to accurately determine the background value in each wavelength channel."," However, strictly speaking flat-fielding is required to accurately determine the background value in each wavelength channel."
+" This is particularly important, as underestimating or overestimating the background results in a variable transit depth, which is exactly what we are trying to measure."," This is particularly important, as underestimating or overestimating the background results in a variable transit depth, which is exactly what we are trying to measure."
+" Furthermore, the G206 grism has a relatively high background compared with NICMOS's other two grisms (about 800 electrons per pixel per exposure), due to thermal background emission from HST."," Furthermore, the G206 grism has a relatively high background compared with NICMOS's other two grisms (about $800$ electrons per pixel per exposure), due to thermal background emission from HST."
+" Whilst the background is very stable temporally, unfortunately, it does significantly vary spatially over the detector by as much as 200 electrons per pixel per exposure."," Whilst the background is very stable temporally, unfortunately, it does significantly vary spatially over the detector by as much as 200 electrons per pixel per exposure."
+" This will particularly effect the light curves extracted from the edge of the spectrum, where the overall counts are much lower."," This will particularly effect the light curves extracted from the edge of the spectrum, where the overall counts are much lower."
+ We did try flat-fielding each image using a flat-field taken with the G206 grism in place., We did try flat-fielding each image using a flat-field taken with the G206 grism in place.
+ This corrects for diffuse light over the whole detector., This corrects for diffuse light over the whole detector.
+ This will not effect the depth of the final normalised light curves except through the background value., This will not effect the depth of the final normalised light curves except through the background value.
+" This process does not provide a satisfactory background correction, and the analyses described below were attempted with and without flat-field corrections."," This process does not provide a satisfactory background correction, and the analyses described below were attempted with and without flat-field corrections."
+" We also note that the background varies quite smoothly, and therefore is unlikely to be responsible for any narrow 'features' seen in the transmission spectrum, but can certainly affect the overall depth and shape."," We also note that the background varies quite smoothly, and therefore is unlikely to be responsible for any narrow `features' seen in the transmission spectrum, but can certainly affect the overall depth and shape."
+" Prior to extracting the spectra, we identified the defective pixels flaggedby the calibration pipeline, and additional significant outliers, and corrected for them by"," Prior to extracting the spectra, we identified the defective pixels flaggedby the calibration pipeline, and additional significant outliers, and corrected for them by"
+loop. aud was then split. presumably due to the small coronal loop.,"loop, and was then split, presumably due to the small coronal loop."
+ Ou the east side. a weak front became detached aliead of the main strip. as respectively iudicated by the yellow aud red arrows in the third panel of Figure 1..," On the east side, a weak front became detached ahead of the main strip, as respectively indicated by the yellow and red arrows in the third panel of Figure \ref{fig1}."
+ The weak front propagated much faster than the main strip. aud the two [routs were widely separated at 09:00:55 UT. as illustrated by the right pauel of Figure 1..," The weak front propagated much faster than the main strip, and the two fronts were widely separated at 09:00:55 UT, as illustrated by the right panel of Figure \ref{fig1}."
+ On the west side. ouly a bright strip is disceruable by eves. which propagated quickly.," On the west side, only a bright strip is discernable by eyes, which propagated quickly."
+" Iu order to investigate the dynamics of the wave [routs more quantitatively. we analyze the time evolution of the brightuess distributions along two slices. A aud B. which are separated by the stnall coronal loop at(7207... —200"") as indicated in Figure 1.."," In order to investigate the dynamics of the wave fronts more quantitatively, we analyze the time evolution of the brightness distributions along two slices, A and B, which are separated by the small coronal loop at, $-200\arcsec$ ) as indicated in Figure \ref{fig1}."
+ The two slices are both great. circles along the solar surface linking the flare site., The two slices are both great circles along the solar surface linking the flare site.
+ The results are displayed in Figure 2. as a space-time plot. where any oblique patter nieauns a propagating front.," The results are displayed in Figure \ref{fig2} as a space-time plot, where any oblique pattern means a propagating front."
+ The top panel of Figure 2. displays the evolution of the wave front. propagation aloug slice A. We can identify several fronts: Fl. F2. aud S1.," The top panel of Figure \ref{fig2} displays the evolution of the wave front propagation along slice A. We can identify several fronts: F1, F2, and S1."
+ Front F1 appeared first. whieh was propagating with a high velocity of 560 uuntil 08:52:LE UT.," Front F1 appeared first, which was propagating with a high velocity of 560 until 08:52:44 UT."
+ We interpret it as a fast-mode wave., We interpret it as a fast-mode wave.
+ A separate bright front 51 was propagating with an initial speed of 190 ssince 05:52:00 UT., A separate bright front S1 was propagating with an initial speed of 190 since 08:52:00 UT.
+ Its speed was decreasing., Its speed was decreasing.
+" After 09:02:28 UT. front 51 stopped at a distance of ~9253"" [rom the flare site."," After 09:02:28 UT, front S1 stopped at a distance of $\sim 253\arcsec$ from the flare site."
+ We identify it as the classical “EIT wave.," We identify it as the classical “EIT wave""."
+ Front S1 might be already existing duriug the period from 08:19:04 UT to 08:52:00 UT. but it is almost uuclistineuishable [rom the fast front. Fl in Figure 2..," Front S1 might be already existing during the period from 08:49:04 UT to 08:52:00 UT, but it is almost undistinguishable from the fast front F1 in Figure \ref{fig2}."
+ It dis interesting to see that au additional wave patteru. front ΕΞ. emanated from (rout Sl at the distauce of ~220” from the [Lare site at 08:5E04 UT.," It is interesting to see that an additional wave pattern, front F2, emanated from front S1 at the distance of $\sim 220\arcsec$ from the flare site at 08:54:54 UT."
+ This weak front. propagated with a speed o£ 3101., This weak front propagated with a speed of 310.
+.. We interpret it as a fast-inode wave., We interpret it as a fast-mode wave.
+" An important reason is that after a distance of ~270"" the [ast front FI began to decelerate. and propagated with alinost the same speed as front. F2."," An important reason is that after a distance of $\sim 270\arcsec$ the fast front F1 began to decelerate, and propagated with almost the same speed as front F2."
+ The bottom panel of Figure 2. illustrates the wave [rout propagation aloug slice B. The wave pattern here is very simple., The bottom panel of Figure \ref{fig2} illustrates the wave front propagation along slice B. The wave pattern here is very simple.
+ Besides a fast-moving bright frout. F3. which had a velocity of 170‘.. there is a weak front. 82. propagating outward with a velocity of 170'.," Besides a fast-moving bright front, F3, which had a velocity of 470, there is a weak front, S2, propagating outward with a velocity of 170."
+". We identify the faster front F3 as a fast-inode wave aid the slower front 52 as an “EIT wave""."," We identify the faster front F3 as a fast-mode wave and the slower front S2 as an “EIT wave""."
+ Since the discovery. “EIT waves” were widely explained as fast-uiode iuagnetoacoustie waves. or coroual Moreton waves. as reviewed iu relsect iuntro..," Since the discovery, “EIT waves"" were widely explained as fast-mode magnetoacoustic waves, or coronal Moreton waves, as reviewed in \\ref{sect:intro}. ."
+ However. inauy of their characteristics caunot be explaiued in the [ramework of waves. which stimulated to doubt the wave nature ofthe “EIT waves”.," However, many of their characteristics cannot be explained in the framework of fast-mode waves, which stimulated to doubt the wave nature of the “EIT waves""."
+ Based ou MHD simulatious. proposed a maguetic fielcline stretcliiug model. i.e.. “EIT waves” are not fast-inocle waves. they are apparent propagation ofthe brighteniugs from," Based on MHD simulations, proposed a magnetic fieldline stretching model, i.e., “EIT waves"" are not fast-mode waves, they are apparent propagation ofthe brightenings from"
+richness and the weak lensing mass approximately linearly.,richness and the weak lensing mass approximately linearly.
+ The inferred X-ray temperature on the other hand will be close to that of the more massive system. whereas the X-ray lummosity will be much lower than expected. because it is proportional to the square of the electron density.," The inferred X-ray temperature on the other hand will be close to that of the more massive system, whereas the X-ray luminosity will be much lower than expected, because it is proportional to the square of the electron density."
+ We note that both RXJO847.1+3449 and RXJ1354.2—022] show evidence of strong lensing., We note that both $0847.1+3449$ and $1354.2-0221$ show evidence of strong lensing.
+ Interestingly. the X-ray image of RXJO847.1+43449 in ? shows evidence for a nearby cluster candidate.," Interestingly, the X-ray image of $0847.1+3449$ in \cite{Lumb04}
+ shows evidence for a nearby cluster candidate."
+ The case is less clear for RXJ1354.2-—0221. but the X-ray image shows a complex morphology.," The case is less clear for $1354.2-0221$, but the X-ray image shows a complex morphology."
+ Unfortunately we lack the dynamical data to confirm whether RXJO847.1+3449 and RXJI354.2—022] are projected systems., Unfortunately we lack the dynamical data to confirm whether $0847.1+3449$ and $1354.2-0221$ are projected systems.
+ The cluster RXJ1117.4+0743 was studied in detail by 2.. who find that this cluster is also a projection of two structures.," The cluster $1117.4+0743$ was studied in detail by \cite{Carrasco07}, who find that this cluster is also a projection of two structures."
+ For the main component ? infer a galaxy velocity dispersion. of 592+82 knys. whereas the second structure is less massive with σι=39]+85 knys. Based on our lensing analysis we obtained a velocity dispersion of 639*5 km/s. in good agreement with the dynamical results for the main cluster.," For the main component \cite{Carrasco07} infer a galaxy velocity dispersion of $592\pm82$ km/s, whereas the second structure is less massive with $\sigma_v=391\pm85$ km/s. Based on our lensing analysis we obtained a velocity dispersion of $639^{+76}_{-86}$ km/s, in good agreement with the dynamical results for the main cluster."
+ ? also performed a weak lensing analysis based on ground based imaging data and obtained a velocity dispersion of 7=778+89 km/s (where we took the average of their results for the σ΄ and 7’ band). implying a mass larger than our estimate.," \cite{Carrasco07} also performed a weak lensing analysis based on ground based imaging data and obtained a velocity dispersion of $\sigma=778\pm89$ km/s (where we took the average of their results for the $g'$ and $r'$ band), implying a mass larger than our estimate."
+ We are not able identify an obvious cause for this difference. but note that PSF-related systematics have a larger impact on ground based results.," We are not able identify an obvious cause for this difference, but note that PSF-related systematics have a larger impact on ground based results."
+ In this section we examine the correlation between M»soo and the X-ray luminosity. in. particular the normalization and the power law slope.," In this section we examine the correlation between $M_{2500}$ and the X-ray luminosity, in particular the normalization and the power law slope."
+ We fit a power law model to the combined sample to maximize the leverage in X-ray luminosity If we naively fit this model to the measurements we find that the value for 4 of the best fit is too high (47=106 with 43 degrees of freedom)., We fit a power law model to the combined sample to maximize the leverage in X-ray luminosity If we naively fit this model to the measurements we find that the value for $\chi^2$ of the best fit is too high $\chi^2=106$ with 43 degrees of freedom).
+ This indicates that there ts intrinsic scatter m the relation. which is also apparent from Figure 6..," This indicates that there is intrinsic scatter in the relation, which is also apparent from Figure \ref{mlx}."
+ We need to account for the intrinsic scatter in the fitting procedure. because ignoring it will generally bias the best fit parameters.," We need to account for the intrinsic scatter in the fitting procedure, because ignoring it will generally bias the best fit parameters."
+ We fit the model to our measurements. which have errors that follow a normal distribution.," We fit the model to our measurements, which have errors that follow a normal distribution."
+ The intrinsic scatter. however. can be described by a log-normal distribution (seee.g..Fig.13mn?)..," The intrinsic scatter, however, can be described by a log-normal distribution \citep[see e.g., Fig. 13
+ in][]{Vikhlinin09a}."
+ We will assume that the intrinsic scatter can be approximated with a normal distribution with a dispersion o~In(l0)logQ (we use the log with base 10)., We will assume that the intrinsic scatter can be approximated with a normal distribution with a dispersion $\sigma_Q\approx \ln(10)Q\log Q$ (we use the log with base 10).
+ To fit the Mosoo—Ly relation we follow a maximum likelihood approach., To fit the $M_{2500}-L_X$ relation we follow a maximum likelihood approach.
+ For a model f with parameters a. the predicted values are y;=f(vj:a).," For a model $f$ with parameters ${\bf a}$, the predicted values are $y_i=f(x_i;{\bf a})$."
+ The uncertainties inv; and v; are given by σι; and σε.," The uncertainties in $x_i$ and $y_i$ are given by $\sigma_{x,i}$ and $\sigma_{y,i}$."
+ If we assume a Gaussian intrinsic scatter 70; in the v; coordinate. the likelihood £ is given by where w; accounts for the scatter: If we consider the logarithm of the likelihood it becomes clear why including the intrinsic scatter differs from standard least squares minimization: where the second term corresponds to the usual X.," If we assume a Gaussian intrinsic scatter $\sigma_{Q,i}$ in the $y_i$ coordinate, the likelihood ${\cal L}$ is given by where $w_i$ accounts for the scatter: If we consider the logarithm of the likelihood it becomes clear why including the intrinsic scatter differs from standard least squares minimization: where the second term corresponds to the usual $\chi^2$."
+ For a linear relation with no intrinsic scatter. the first term is à constant for a given data set and the likelihood is maximimized by minimizing 4.," For a linear relation with no intrinsic scatter, the first term is a constant for a given data set and the likelihood is maximimized by minimizing $\chi^2$."
+ However. if intrinsic scatter is included as a free parameter. the first term acts as a penalty function. and cannot be ignored.," However, if intrinsic scatter is included as a free parameter, the first term acts as a penalty function, and cannot be ignored."
+ The presence of intrinsic scatter. also. exacerbates the Malmquist bias for a flux limited sample. such as the 1608D," The presence of intrinsic scatter also exacerbates the Malmquist bias for a flux limited sample, such as the 160SD"
+2) the Hubble tvpe (0— ο ορ. 1= S0. 2= σα SO/Sa. 3= Sa. 4d= Sab.5= Sb.6 = She. 7= Se (ο). S= Scd. 9= Sd. 102 Sdm Self/Sm. 11= Sm. 122 Im (lm/s). 13= generic irregulars of Coma and. peculiar galaxies. 14= S/DCD (dS/DBCD. dS0/BCD. ας). 15= Sm/DCD. 16= ln/DCD): 3) the major axis: 4) the minor axis: 5) the characteristic velocity. (i6... either the central velocity. dispersion a or the maximum rotational velocity Varese according to morphology): 6) the effective radius r5: 7) the elfective surface brightness ya. (the effective surface intensity in Ll-banc solar units is given by fo=IO(use 25015. 10) 5&3.,"2) the Hubble type $0 =$ E – E/S0, $1 =$ S0, $2 =$ S0a – S0/Sa, $3 =$ Sa, $4 =$ Sab,$5 =$ Sb, $6 =$ Sbc, $7 =$ Sc (dSc), $8 =$ Scd, $9 =$ Sd, $10 =$ Sdm – Sd/Sm, $11 =$ Sm, $12 =$ Im (Im/S), $13 =$ generic irregulars of Coma and peculiar galaxies, $14 =$ S/BCD (dS/BCD, dS0/BCD, Sd/BCD), $15 =$ Sm/BCD, $16 =$ Im/BCD); 3) the major axis; 4) the minor axis; 5) the characteristic velocity (i.e., either the central velocity dispersion $\sigma_0$ or the maximum rotational velocity $V_{Max}$ according to morphology); 6) the effective radius $r_e$; 7) the effective surface brightness $\mu_e$ (the effective surface intensity in H-band solar units is given by $I_e = 10^{-0.4 (\mu_e - 24.91)}$ ); 8) $\kappa_1$; 9) $\kappa_2$; 10) $\kappa_3$."
+" The values of à, ancl αν listed in Tab.", The values of $r_e$ and $\mu_e$ listed in Tab.
+ 1: are corrected for extinction as described in the next section., 1 are corrected for extinction as described in the next section.
+ Dust extinction. alfects. the observed. elobal metric and whotometrie properties of potentially any galaxy. somehow. since dust. production is a result. of star formation.," Dust extinction affects the observed global metric and photometric properties of potentially any galaxy somehow, since dust production is a result of star formation."
+ The magnitude of these elects changes with the pass-band. of he observations and depends on the physical ancl chemical oxoperties of the dust grains. the topology of the dust. the distribution of the cust with respect to the dilferent stellar populations and the attenuation optical cepth along the inc-of-sight in à non trivial wav.," The magnitude of these effects changes with the pass-band of the observations and depends on the physical and chemical properties of the dust grains, the topology of the dust, the distribution of the dust with respect to the different stellar populations and the attenuation optical depth along the line-of-sight in a non trivial way."
+ A careful correction of the whotometric properties for dust attenuation. in. particular. requires modeling. of the stellar populations and. of the radiative transfer within cach galaxy. on one hand. and a multi-wavelength data-set for any individual galaxy. on the other.," A careful correction of the photometric properties for dust attenuation, in particular, requires modeling of the stellar populations and of the radiative transfer within each galaxy, on one hand, and a multi-wavelength data-set for any individual galaxy, on the other."
+ This task is prohibitive for a large sample like ours., This task is prohibitive for a large sample like ours.
+ Therefore. we have to make a set of assumptions and adopt ἃ set of corrections based. on statistics.," Therefore, we have to make a set of assumptions and adopt a set of corrections based on statistics."
+ The relatively low exctinction in the L-band (Cardelli. Clavton Mathis 1989 and references therein) gives some insurance against large systematic errors. though the magnitude of the error may vary on an incividual basis.," The relatively low exctinction in the H-band (Cardelli, Clayton Mathis 1989 and references therein) gives some insurance against large systematic errors, though the magnitude of the error may vary on an individual basis."
+ Elliptical ancl lenticular galaxies may contain dust. either in the form. of lanes ancl patehes or of a diffuse component of the interstellar medium (Bertola Calletta 1978: Lawarden ct al.," Elliptical and lenticular galaxies may contain dust, either in the form of lanes and patches or of a diffuse component of the interstellar medium (Bertola Galletta 1978; Hawarden et al."
+ 1981: Jura 1986)., 1981; Jura 1986).
+ The wavelength-dependent clleets of dust. absorption and scattering on their photometric properties have been studied by Witt. Thronson Capuano (1992) under the assumption of a cliffuse dust. component.," The wavelength-dependent effects of dust absorption and scattering on their photometric properties have been studied by Witt, Thronson Capuano (1992) under the assumption of a diffuse dust component."
+ As shown by these authors. attenuation optical depths of 0.11 in the ραπ (A=0.55yam) max. explain some of the photometric properties of these stellar svtems. e.g. their color gradients.," As shown by these authors, attenuation optical depths of 0.1–1 in the V-band $\rm \lambda =0.55~\mu m$ ) may explain some of the photometric properties of these stellar sytems, e.g. their color gradients."
+ Starlight attenuation is primarily the result of absorption in the near-Ik. where the absorption optical depths are lower than in the optical (Witt. Phronson Capuano).," Starlight attenuation is primarily the result of absorption in the near-IR, where the absorption optical depths are lower than in the optical (Witt, Thronson Capuano)."
+ Dust is found capable to alfect the projection of the kinematical quantities of these stellar svstems as well (λαο. Dejonghe 2000: Daes. Dejonghe De Rijeke 2000).," Dust is found capable to affect the projection of the kinematical quantities of these stellar systems as well (Baes Dejonghe 2000; Baes, Dejonghe De Rijcke 2000)."
+ Nonetheless. under the assumption that only the ciffuse dust component is present in elliptical. and lenticular svstems. we may reasonably assume that these stellar svstems are basically dust-free in the L-bancl. in agreement with Baes Dejonghe.," Nonetheless, under the assumption that only the diffuse dust component is present in elliptical and lenticular systems, we may reasonably assume that these stellar systems are basically dust-free in the H-band, in agreement with Baes Dejonghe."
+ This assumption is questionable for spiral and irregular ealaxies., This assumption is questionable for spiral and irregular galaxies.
+ The debate on the opacity of the disk-component of spiral ancl irregular galaxies is still open and its discussion is bevond our scope (see Disney et al., The debate on the opacity of the disk-component of spiral and irregular galaxies is still open and its discussion is beyond our scope (see Disney et al.
+ 1989 for a review)., 1989 for a review).
+ Lf clust is cistributed only in the disk. disks are more absorbed at high inclinations. because of the larger optical depth of the dust in the plane of the galaxy.," If dust is distributed only in the disk, disks are more absorbed at high inclinations, because of the larger optical depth of the dust in the plane of the galaxy."
+ By contrast. bulges are more attenuated at low inclinations. as in this case basically. only the foreground light of the svstem is observed. whereas at high inclinations. almost the entire bulge comes into view. if the clust thickness is not substantial with respect to the bulece cllective radius (Ferrara ct al.," By contrast, bulges are more attenuated at low inclinations, as in this case basically only the foreground light of the system is observed, whereas at high inclinations, almost the entire bulge comes into view, if the dust thickness is not substantial with respect to the bulge effective radius (Ferrara et al."
+ L999)., 1999).
+" As a consequence. in. bulge|disk systems. the dust cllects on rr, ancl ye, will depend not only on the opacity along the line-of-sight but also on the bulge-to-clisk ratio at a given pass-hancl."," As a consequence, in bulge+disk systems, the dust effects on $r_e$ and $\mu_e$ will depend not only on the opacity along the line-of-sight but also on the bulge-to-disk ratio at a given pass-band."
+ ‘To make things more complex. bulges may be clusty too oletier et al.," To make things more complex, bulges may be dusty too (Peletier et al."
+ 1999)., 1999).
+ For bulge|disk svstenis. recipes for the correction of n. aud fmi. for dust attenuation. motivated either ov radiative transfer models or by observational stuclies. still lack.," For bulge+disk systems, recipes for the correction of $r_e$ and $\mu_e$ for dust attenuation, motivated either by radiative transfer models or by observational studies, still lack."
+ By contrast. dust attenuation in exponential disks its driven more attention both from theory (e.g. Bianchi et al.," By contrast, dust attenuation in exponential disks has driven more attention both from theory (e.g. Bianchi et al."
+ 2000) and observations (e.g. Moriondo. Ciovanelli Lavnes 1999).," 2000) and observations (e.g. Moriondo, Giovanelli Haynes 1999)."
+" For the purposes of our study. we make the ollowing assumptions: i) dust is distributed only in the clisk-component of spiral and irregular galaxies: ii) the inclination corrections for r. and yr, are the same as for the exponential disk scale-leneth () ancl central surface brightness (ο). respectively. whatever the late IIubble type of the galaxy is."," For the purposes of our study, we make the following assumptions: i) dust is distributed only in the disk-component of spiral and irregular galaxies; ii) the inclination corrections for $r_e$ and $\mu_e$ are the same as for the exponential disk scale-length $h$ ) and central surface brightness $\mu_0$ ), respectively, whatever the late Hubble type of the galaxy is."
+ Given the previous discussion. we aknowledge that the second assumption over-simplifies the problem.," Given the previous discussion, we aknowledge that the second assumption over-simplifies the problem."
+ However. at least this assumption does not lead to an inconsistency when we compare the effective Buminosities obtained through I5qu.," However, at least this assumption does not lead to an inconsistency when we compare the effective luminosities obtained through Equ."
+ 10. after applving the extinction correction of Morioncdo. Giovanelli Llavnes (1999) tor. and jn. and the total luminosities obtained from our data analysis and corrected for internal extinction (Cialactic extinction is neelieible in the L-bancl) according to the recipes of Gayazzi Boselli (1996).," 10, after applying the extinction correction of Moriondo, Giovanelli Haynes (1999) to $r_e$ and $\mu_e$ and the total luminosities obtained from our data analysis and corrected for internal extinction (Galactic extinction is negligible in the H-band) according to the recipes of Gavazzi Boselli (1996)."
+" The corrections for inclination applied to +, and 44 are:", The corrections for inclination applied to $r_e$ and $\mu_e$ are:
+"Then. the X-ray surface brightness was calculated as and the ταν emissivitv-welghted surface temperature was calculated as where 7, is the temperature of grid zone A. Here. d.(2,) is the comoving distance up to τε. Which is given as The contribution from clusters was subtracted from calculated RAL based on different criteria using X-ray brightness and temperature.","Then, the X-ray surface brightness was calculated as and the X-ray emissivity-weighted surface temperature was calculated as where $T_{Xk}$ is the temperature of grid zone $k$ Here, $d_c(z_k)$ is the comoving distance up to $z_k$, which is given as The contribution from clusters was subtracted from calculated RM based on different criteria using X-ray brightness and temperature."
+ The simplest approach woulcl be to exclude the volume associated with clusters in (he cosmic space., The simplest approach would be to exclude the volume associated with clusters in the cosmic space.
+ We emploved the following two viteria., We employed the following two criteria.
+" In the first model. labeled “Ελ, all the hot gas with T>10* IN was excluded."," In the first model, labeled “TM7”, all the hot gas with $T > 10^7$ K was excluded."
+ In ve second model. labeled “CLS”. the spherical region of comoving radius Lf+ \Ipe around je Clusters and groups will Ly>2 keV was excluded.," In the second model, labeled “CLS”, the spherical region of comoving radius $1 h^{-1}$ Mpc around the clusters and groups with $T_X > 2$ keV was excluded."
+ For simulations of RAL surveys. however. a better approach would be the exclusion of je. pixels (hat include clusters in the projected skv map.," For simulations of RM surveys, however, a better approach would be the exclusion of the pixels that include clusters in the projected sky map."
+ We would like to subtract the ontribution from all clusters. but it is still hard to detect faint X-ray clusters.," We would like to subtract the contribution from all clusters, but it is still hard to detect faint X-ray clusters."
+ So we tried the following criteria., So we tried the following criteria.
+ In the model labeled “ως. all the pixels with 7$>LO’ Iv and SX>10 Sergsbem7sv| were excluded: 10 “ergs!em7sr Lis close to the detection limit of current. N-ray facilities 2010).," In the model labeled “TS8”, all the pixels with $T_X^* > 10^7$ K and $S_X^* > 10^{-8}$ ${\rm erg~s^{-1}~cm^{-2}~sr^{-1}}$ were excluded; $10^{-8}$ ${\rm erg~s^{-1}~cm^{-2}~sr^{-1}}$ is close to the detection limit of current X-ray facilities ."
+". And in another model labeled ""ESQ"".:"" the pixels. with. 73"">10* TqI and ονTEN>10.I erg.sFem72sr! were excluded: 101! eresFem?se! is intended to mimic the improved detection limit of future X-ray facilities."," And in another model labeled “TS0”, the pixels with $T_X^* > 10^7$ K and $S_X^* > 10^{-10}$ ${\rm erg~s^{-1}~cm^{-2}~sr^{-1}}$ were excluded; $10^{-10}$ ${\rm erg~s^{-1}~cm^{-2}~sr^{-1}}$ is intended to mimic the improved detection limit of future X-ray facilities."
+ Ilere. ον and 7$. up to z=5 were used.," Here, $S_X^*$ and $T_X^*$ up to $z=5$ were used."
+"Finally. we also considered the case where no volume or pixel wasexcluded. labeled ""ALL. for comparison.","Finally, we also considered the case where no volume or pixel wasexcluded, labeled “ALL"", for comparison."
+Supergranules are cellular flow structures observed in the solar photosphere with typical diameters of |about 30 Mh. lifetimes of |about one day. ettaud flow velocities of 300 mis+ (Bieutord&Rincon 2010).,"Supergranules are cellular flow structures observed in the solar photosphere with typical diameters of about 30 Mm, lifetimes of about one day, and flow velocities of 300 $\rm m\ s^{-1}$ \citep{RieutordRincon10}."
+". They cover the cutive surface of the Sum except for the imuuediate surrondings of snuspots,", They cover the entire surface of the Sun except for the immediate surroundings of sunspots.
+ Supergranules. were discovered. by, Supergranules were discovered by
+arbitrary phase 0.55.,arbitrary phase 0.55.
+ Phe ephemeris was then used to isolate time intervals when X-ray outbursts were expected., The ephemeris was then used to isolate time intervals when X-ray outbursts were expected.
+ The duration of the time interval examined. either side of long period minima was LOO days to cover the data. coincident with the optical lighteurve “clip” regions.," The duration of the time interval examined either side of long period minima was 100 days to cover the data coincident with the optical lightcurve ""dip"" regions."
+ According to Skinner (1982)... X0538-668 was in the Ποιά of view of the Einstein satellite on four occasions.," According to Skinner \shortcite{sk}, A0538-668 was in the field of view of the Einstein satellite on four occasions."
+ According το the ephemeris described. here. detections reported on 16 Dec 1980 and 3 Feb 1981 lie in the optical dip. and non detections on LO Apr 1979 and 30 May. 1980 Πο firmlv in the centre of the bright phase of the optical liehteurve.," According to the ephemeris described here, detections reported on 16 Dec 1980 and 3 Feb 1981 lie in the optical dip, and non detections on 10 Apr 1979 and 30 May 1980 lie firmly in the centre of the bright phase of the optical lightcurve."
+ ὃν creating test datasets with the same mean anc variance as the BATSE data and containing only simulated: signals. we determined how strong a periodic signal must be in order to give a 3o detection.," By creating test datasets with the same mean and variance as the BATSE data and containing only simulated signals, we determined how strong a periodic signal must be in order to give a $\sigma$ detection."
+ This approach could also be used to model other components of a power spectrum. particularly where interfering sources and sampling structure lead to spurious frequencies.," This approach could also be used to model other components of a power spectrum, particularly where interfering sources and sampling structure lead to spurious frequencies."
+ Ixnown interfering signals in the X0538-668 data include LMC γε (which. is only 0.67 away and. dominates. the periodic Dux in this dataset) and the spacecraft precessiona moculation., Known interfering signals in the A0538-668 data include LMC X-4 (which is only $\degr$ away and dominates the periodic flux in this dataset) and the spacecraft precessional modulation.
+ These components were modelled by measuring their amplitudes in folded lighteurves at 35.5d. and. 5X respectively., These components were modelled by measuring their amplitudes in folded lightcurves at 35.5d and 52d respectively.
+ Using the ephemoeris described above BATSE data were selected which were coincident with the four long-perioc optical minima observed during the ALACLLO observations of Alcock et al €1999).., Using the ephemeris described above BATSE data were selected which were coincident with the four long-period optical minima observed during the MACHO observations of Alcock et al \shortcite{al}.
+ This resulted in four blocks of BATSE one-day averages cach of duration LOO clays spanning the time interval PJD 9124.52 - TID 109995 Lavine Lailec to detect any statistically significant. moculation consistent with 16.65 days. this subset of the BATSE observations of 0538-668. was then repeatedly randomized. to destroy any periodic content ancl combined. with an incrementalvy increasing sinusoidal signal plus fixed sinewave components to simulate the interfering sources in the real data.," This resulted in four blocks of BATSE one-day averages each of duration $~$ 100 days spanning the time interval TJD 9124.52 - TJD 10999.5 Having failed to detect any statistically significant modulation consistent with 16.65 days, this subset of the BATSE observations of A0538-668 was then repeatedly randomized to destroy any periodic content and combined with an incrementaly increasing sinusoidal signal plus fixed sinewave components to simulate the interfering sources in the real data."
+ After each step. the Lomb-Scarele periodogram was computed are the highest peak in a pre-determined frequency range was compared to the 3 o confidence level previously caleulatec for that [requeney. range as described in section 2.3.," After each step, the Lomb-Scargle periodogram was computed and the highest peak in a pre-determined frequency range was compared to the 3 $\sigma$ confidence level previously calculated for that frequency range as described in section 2.3."
+ Once this threshold was exceeded. the amplitude of the sine wave signal was recorded.," Once this threshold was exceeded, the amplitude of the sine wave signal was recorded."
+ The detectability of such a weak signa is somewhat inlluenced. by the noise present., The detectability of such a weak signal is somewhat influenced by the noise present.
+ At the very margins of detectability the power in a given frequency can Iluctuate up and down despite a gradual increase in signa amplitude., At the very margins of detectability the power in a given frequency can fluctuate up and down despite a gradual increase in signal amplitude.
+ Hence we used the criterion that the power in the selected. frequency. range must remain above the thresho value irrespective of chance structures in. noise in order qualify., Hence we used the criterion that the power in the selected frequency range must remain above the threshold value irrespective of chance structures in noise in order to qualify.
+ The derived. upper limit was independently testec using many cdillerent noise sets to verily this criterion. was met. Le. That a signal of the amplitude: quoted wou definitely have been detected by the methods described her had it been present.," The derived upper limit was independently tested using many different noise sets to verify this criterion was met, i.e. That a signal of the amplitude quoted would definitely have been detected by the methods described here, had it been present."
+ Weaker signals would not consistentA rise above the threshold although detections of marginal significance are possible., Weaker signals would not consistently rise above the threshold although detections of marginal significance are possible.
+ This object has been subject. of several recent. studies., This object has been subject of several recent studies.
+ Reig et al (1998) presented results of multiwavelength observations ancl mocdeling of the accretion rate.," Reig et al \shortcite{rei}
+ presented results of multi-wavelength observations and modeling of the accretion rate."
+ Wilson οἱ al (2002) report on a decade of X-ray observations combined with IR and Ho measurements which show a decline in the density of the circumstellar disk around the Be star., Wilson et al (2002) report on a decade of X-ray observations combined with IR and $\alpha$ measurements which show a decline in the density of the circumstellar disk around the Be star.
+ This decline was followed. by a sudden drop in the X-ray [lux and a turn-over from a spin-up trend to spin-down in the frequeney history., This decline was followed by a sudden drop in the X-ray flux and a turn-over from a spin-up trend to spin-down in the frequency history.
+ This is the first Be/N-ray. binary to show an extended interval. about 2.5 vears. where the global trend is spin-down. but the outbursts continue.," This is the first Be/X-ray binary to show an extended interval, about 2.5 years, where the global trend is spin-down, but the outbursts continue."
+ In 1995 the orbital phase of EXO 2030|375's outbursts shifted. from peaking about 6 davs after periastron to peaking before periastron., In 1995 the orbital phase of EXO 2030+375's outbursts shifted from peaking about 6 days after periastron to peaking before periastron.
+ The outburst phase slowly recovered to peaking at about 2.5 davs after periastron., The outburst phase slowly recovered to peaking at about 2.5 days after periastron.
+ Wilson et al (2002) interpret this shift in orbital phase followed by a slow recovery as evidence for a elobal one-armed oscillation propagating in the De disk., Wilson et al (2002) interpret this shift in orbital phase followed by a slow recovery as evidence for a global one-armed oscillation propagating in the Be disk.
+ This was further supported by changes in the shape of the llo profile which are commonly. believed. to be. produced by à reconfiguration of the Be disk., This was further supported by changes in the shape of the $\alpha$ profile which are commonly believed to be produced by a reconfiguration of the Be disk.
+ The truncated: viscous decretion disk model provided an explanation for the long series of normal outbursts and the evidence for an aceretion disk in the brighter normal outbursts., The truncated viscous decretion disk model provided an explanation for the long series of normal outbursts and the evidence for an accretion disk in the brighter normal outbursts.
+ In the present analvsis we find convincing evidence or elevated. flux levels at. apastron in the earlier epoch (sce Figure. 1)., In the present analysis we find convincing evidence for elevated flux levels at apastron in the earlier epoch (see Figure \ref{fig:lightcurves}) ).
+ The RNTIS ASM lighteurve for. epoch 2 was studied by Reig et al (.1998)... we also analysed his data in order to compare the lighteurve and power spectrum with the BAPSE results.," The RXTE ASM lightcurve for epoch 2 was studied by Reig et al \shortcite{rei}, we also analysed this data in order to compare the lightcurve and power spectrum with the BATSE results."
+" “Phe power spectra rom both instruments appear to share the same features in epochs 1 2. peaks at Pas, Paof2 with powers at similar ratios."," The power spectra from both instruments appear to share the same features in epochs 1 2, peaks at $P_{orb}$ $P_{orb}/2$ with powers at similar ratios."
+ Direct comparison of the BATSE and INTE folded lighteurves suggests that X-ray emission from 12XO2030|375 may be softer during the apastron outbursts han during the normal Type E outbursts at periastron., Direct comparison of the BATSE and RXTE folded lightcurves suggests that X-ray emission from EXO2030+375 may be softer during the apastron outbursts than during the normal Type I outbursts at periastron.
+ This ivpothesis is drawn from the relative amplitudes of the two »aks in the lighteurve and the dilfering spectral responses of the two instruments., This hypothesis is drawn from the relative amplitudes of the two peaks in the lightcurve and the differing spectral responses of the two instruments.
+ Comparison of the FWIIM of the orbital period peak at 46c with that derived. from. simulated data reveals that although the phase of peak outburst intensity shifts between the two epochs. its recurrence is otherwise periodic.," Comparison of the FWHM of the orbital period peak at 46d with that derived from simulated data reveals that although the phase of peak outburst intensity shifts between the two epochs, its recurrence is otherwise periodic."
+ There is no evidence for any gradual shift in phase in either epoch., There is no evidence for any gradual shift in phase in either epoch.
+ We find that during the period coincident with the NLACTIO monitoring of Aleock et al (1999)... A0538-668 cick not exhibit hard X-ray outbursts at the 16.65 αν orbital period during either the optical minima or maxima.," We find that during the period coincident with the MACHO monitoring of Alcock et al \shortcite{al}, A0538-668 did not exhibit hard X-ray outbursts at the 16.65 day orbital period during either the optical minima or maxima."
+ The upper limit, The upper limit
+was estimated based on the RMS of the pixels around each emission line.,was estimated based on the RMS of the pixels around each emission line.
+ The detected emission lines are well spectrally resolved and instrumental broadening is negligible in. most cases., The detected emission lines are well spectrally resolved and instrumental broadening is negligible in most cases.
+ Observed emission-line widths without correction for instrumental broadening are also presented in Table 2., Observed emission-line widths without correction for instrumental broadening are also presented in Table 2.
+ We also provide for each HzRG the redshift calculated from the line., We also provide for each HzRG the redshift calculated from the line.
+ For the composite spectrum the flux ratios ofii.Cur|/Crv. and are 1.618. 0.628. and 0.584. respectively.," For the composite spectrum the flux ratios of, and are 1.618, 0.628, and 0.584, respectively."
+ All the emission lines ofiv. and i] (required for the metallicity diagnostic method proposed by Nagao et al.," All the emission lines of, and ] (required for the metallicity diagnostic method proposed by Nagao et al."
+ 2006a) are detected in the spectra of all 9 HzRGs observed., 2006a) are detected in the spectra of all 9 HzRGs observed.
+ In addition. a significant emission 15 also detected in 5 out of 9 objects.," In addition, a significant emission is also detected in 5 out of 9 objects."
+ We provide a 3c upper limit on the emission for objects without a significant detection. by assuming an emission-line width of ~10A in rest frame.," We provide a $\sigma$ upper limit on the emission for objects without a significant detection, by assuming an emission-line width of $\sim 10 \ \AA$ in rest frame."
+ As shown in Figure |. various fainter emission lines are also detected including Simwtl265. O1Simtl305. Cirtl33s. [11402. 11662. Siit1808. [Neiv[02422..," As shown in Figure 1, various fainter emission lines are also detected including $\lambda$ 1265, $\lambda$ 1305, $\lambda$ 1335, $\lambda$ 1402, $\lambda$ 1663, $\lambda$ 1808, $\lambda2422$."
+ We will discuss the properties of these faint emission lines in a forthcoming paper (Matsuoka et al., We will discuss the properties of these faint emission lines in a forthcoming paper (Matsuoka et al.
+ in prep.)., in prep.).
+"following fit for the amplitude, with n an empirical exponent.","following fit for the amplitude, with $n$ an empirical exponent."
+" When Equation (26)) is applied to the results of Figure 5((a), we obtain that the exponent n—1 provides a good fit for the amplitude (see the dotted line in Figure 5((a))."," When Equation \ref{eq:fit}) ) is applied to the results of Figure \ref{fig:flow1}( (a), we obtain that the exponent $n = 1$ provides a good fit for the amplitude (see the dotted line in Figure \ref{fig:flow1}( (a))."
+" On the other hand, we perform in Figure 5((b) a wavelet power spectrum (Torrence&Compo1998) of the signal displayed in Figure 5((a)."," On the other hand, we perform in Figure \ref{fig:flow1}( (b) a wavelet power spectrum \citep{wavelet} of the signal displayed in Figure \ref{fig:flow1}( (a)."
+ We find that the period decreases in time as the thread moves towards the footpoint of the magnetic structure., We find that the period decreases in time as the thread moves towards the footpoint of the magnetic structure.
+ The WKB approximation for the period given by Equation (23)) is in excellent agreement with the position of the maximum of the wavelet spectrum (see dashed line in Figure 5((b))., The WKB approximation for the period given by Equation \ref{eq:period}) ) is in excellent agreement with the position of the maximum of the wavelet spectrum (see dashed line in Figure \ref{fig:flow1}( (b)).
+" As discussed in Section 3.2,, the WKB approximation for the period is much more accurate than the WKB approximation for the amplitude."," As discussed in Section \ref{sec:amplitude}, the WKB approximation for the period is much more accurate than the WKB approximation for the amplitude."
+" In addition, we see that the rate at which the period changes with respect to the value at t=0 is not constant."," In addition, we see that the rate at which the period changes with respect to the value at $t=0$ is not constant."
+" Using Equation (23)) and considering the parameters of this particular simulation, the period of the oscillation when the thread is at z=L/4 has decreased of about with the respect to its initial value at z=0, whereas the period decreases of about when the thread finally reaches the footpoint of the magnetic tube at z=L."," Using Equation \ref{eq:period}) ) and considering the parameters of this particular simulation, the period of the oscillation when the thread is at $z=L/4$ has decreased of about with the respect to its initial value at $z=0$, whereas the period decreases of about when the thread finally reaches the footpoint of the magnetic tube at $z=L$."
+" We repeat the simulation for L,/L=0.2 and a faster flow, vo/Vap=0.1, and take zo/L=—0.25 to consider the thread initially displaced from the center of the flux tube."," We repeat the simulation for $\lp/L = 0.2$ and a faster flow, $v_0 / \vap = 0.1$, and take $z_0 / L = -0.25$ to consider the thread initially displaced from the center of the flux tube."
+" The radial velocity perturbation at z=0 is plotted in Figure 5((c), and the corresponding wavelet power spectrum is shown in Figure 5((d)."," The radial velocity perturbation at $z=0$ is plotted in Figure\ref{fig:flow1}( (c), and the corresponding wavelet power spectrum is shown in Figure \ref{fig:flow1}( (d)."
+" These results are equivalent to those of Figures 5((a)-(b), i.e., both the amplitude and the period of the oscillation depend on the position of the prominence thread within the flux tube, taking both of them their maximum value when the thread is centered."," These results are equivalent to those of Figures \ref{fig:flow1}( (a)–(b), i.e., both the amplitude and the period of the oscillation depend on the position of the prominence thread within the flux tube, taking both of them their maximum value when the thread is centered."
+" As before, Equation (23)) is a very good approximation to the period."," As before, Equation \ref{eq:period}) ) is a very good approximation to the period."
+" In the previous Section, we have used an initial condition for v, that corresponds to the fundamental mode eigenfunction, so only this mode is excited."," In the previous Section, we have used an initial condition for $v_r$ that corresponds to the fundamental mode eigenfunction, so only this mode is excited."
+" However, it is expected that the energy from an arbitrary disturbance of the flux tube is deposited in many normal modes (seeadiscussiononthisis- 2007).."," However, it is expected that the energy from an arbitrary disturbance of the flux tube is deposited in many normal modes \citep[see a discussion on this issue in, e.g.,][]{terradasexc}. ."
+ To represent an arbitrary, To represent an arbitrary
+for the coalescing population compared to mocel A is lower than 0.25.,for the coalescing population compared to model A is lower than 0.25.
+" We couclude that. given the current size of the sample of measured DNS eccentricities. it is not possible to distinguish between binary evolution models: the exeat majorityDOE of. the oues considered. here appear quautitatively consistent with the ecceutricitv measurements, with. the sole exception. of; L2."," We conclude that, given the current size of the sample of measured DNS eccentricities, it is not possible to distinguish between binary evolution models; the great majority of the ones considered here appear quantitatively consistent with the eccentricity measurements, with the sole exception of L2."
+ Last we examine whether the origin of the low eccentricitics is related to unexpectedly low or zero kicks inparted to NS at birth., Last we examine whether the origin of the low eccentricities is related to unexpectedly low or zero kicks imparted to NS at birth.
+ In Table 2. we list the likelihood values for all ten models aud for models with zero or ⋅ ↕∪↖↖⇁↨↘↽↕↸⊳↨↘↽≓⋯⋜↧∶↴⋁↕∐↑⋯∐∖≼∐↴∖↴⊓⋅↕↴⋝∏↑↕, In Table \ref{table:lambda} we list the likelihood values for all ten models and for models with zero or low kick-magnitude distributions at present.
+∪∐↴∖��⋜↧↑, In Fig.
+↻↥⋅↸∖↴∖↴↸∖∐∙↕∐⊡∶↴⋁∙∣⋅⋅⋅ Dl the model odds ratios for the zero kicks vs. the a kick distribution., \ref{fig:oddszero} shows the model odds ratios for the zero kicks vs. the Arzoumanian kick distribution.
+ se the case0asRO of wideS idSDNS the∖ ratios 01 : ot: ye0 models are within or verv close to a factor mostof 2 del unity.. with. the. exception of. two. models (61 andmodels (3) ," In the case of wide DNS the ratios for most of the models are within or very close to a factor of 2 from unity, with the exception of two models (G1 and G3)."
+Dr This could teutativelv be used as an arguineut in favor of the hwpothesis of verv small NS kicks., This could tentatively be used as an argument in favor of the hypothesis of very small NS kicks.
+ However. it is most interesting to examine the case of coalescing binaries: the likelihood. values arezero. for cight models with zero kicks. for all en models with kicks drawn from ao— dOkhkniss| Maxwellian. aud for four models with kicks drawn from ad=5ükkuss + Maxwellian.," However, it is most interesting to examine the case of coalescing binaries: the likelihood values are, for eight models with zero kicks, for all ten models with kicks drawn from a $\sigma=10$ $^{-1}$ Maxwellian, and for four models with kicks drawn from a $\sigma=50$ $^{-1}$ Maxwellian."
+ As a result all these nodels are renderedunlikely., As a result all these models are rendered.
+ The reasou for these values being identically zero can be traced. back to the xoperties of the Ihulse-Tavlor binary PSR D1915|16: it has an eccentricity of 0.617 and all the models above ave distributions such that the probability for ο20.5 is ideutically zero (sce Fig. 1))., The reason for these values being identically zero can be traced back to the properties of the Hulse-Taylor binary PSR B1913+16: it has an eccentricity of 0.617 and all the models above have distributions such that the probability for $e\gtrsim 0.5$ is identically zero (see Fig. \ref{fig:ecoal}) ).
+ Therefore. these mocels cannot allow for the existence of PSR 1015116. and consequently their likelihood becomes equal to zero.," Therefore, these models cannot allow for the existence of PSR B1913+16, and consequently their likelihood becomes equal to zero."
+ Based on the results for coalescing DNS. we couclude that models with vanishinely or moclerately snall kicks (oX0 l3 aye inconsistent with the current observed sample if our assumptions about the masses of exploding stars are correct.," Based on the results for coalescing DNS, we conclude that models with vanishingly or moderately small kicks $\sigma\lesssim50$ $^{-1}$ ) are inconsistent with the current observed sample if our assumptions about the masses of exploding stars are correct."
+ We address the question of the origin of the low ecceutricities measured for observed DNS binaries (close and wide) that has attracted considerable attention iu the recent literature (vaudenHeuvel2001:Chaurasia&Bailes 2005).," We address the question of the origin of the low eccentricities measured for observed DNS binaries (close and wide) that has attracted considerable attention in the recent literature \citep{vdH,Chaur}."
+. We use DNS population models aud associated predictions for eccentrieity distributions aloug with a Bayesian statistical analysis auc quantitative comparison with the observed sample., We use DNS population models and associated predictions for eccentricity distributions along with a Bayesian statistical analysis and quantitative comparison with the observed sample.
+ Asstuing that all model DNS binaries can be detected as binary pulsars. we conclude that for the case of close (coalescing) DNS binaries. models with vanishinely or moderatelv μια kicks (0Xo50καν: 1j are inconsistent with the current observed sample (aud specifically the ITulse- binary pulsar).," Assuming that all model DNS binaries can be detected as binary pulsars, we conclude that for the case of close (coalescing) DNS binaries, models with vanishingly or moderately small kicks $\sigma\lesssim50$ $^{-1}$ ) are inconsistent with the current observed sample (and specifically the Hulse-Taylor binary pulsar)."
+ We also cxamine the iuflueuce of orbital evolution due to eravitational radiation frou, We also examine the influence of orbital evolution due to gravitational radiation from
+similarly intermediate [NTI[/[OTT] ratios aud [OTT]/[OIT| hat are consistent with the short CRB host region.,similarly intermediate [NII]/[OII] ratios and [OIII]/[OII] that are consistent with the short GRB host region.
+ The metallicity-scusitive |NII|/|OIT| ratio suggests that he GRB 060505 burst site has a marginally lower uctallicity than the other regions of its host galaxy. while the [OITI]/|OII| ratio shows that the ionization xuanmeter of the GRB 060505 burst site is consistent with the short CRB region of the diagram.," The metallicity-sensitive [NII]/[OII] ratio suggests that the GRB 060505 burst site has a marginally lower metallicity than the other regions of its host galaxy, while the [OIII]/[OII] ratio shows that the ionization parameter of the GRB 060505 burst site is consistent with the short GRB region of the diagram."
+ Figures 1 and 2 indicate that the ISM properties of both the CRB 160505 burst site aud the other regious of the ealaxy are consistent with the ISM properties of short-duratiou GRBs associated with compact-object iiergers.," Figures \ref{fig:N2Ha} and \ref{fig:N2O2}
+ indicate that the ISM properties of both the GRB 060505 burst site and the other regions of the galaxy are consistent with the ISM properties of short-duration GRBs associated with compact-object mergers."
+ We calculate imetalliities using the Reo; diagnostic originally described in Ἱνονίον Dopita (2002) aud ater refined and quantified in EKobuluickv Isewley (2001)., We calculate metallicities using the $_{23}$ diagnostic originally described in Kewley Dopita (2002) and later refined and quantified in Kobulnicky Kewley (2004).
+ The Ros diagnostic is double-valued. so we use he [NI] criterion of Kewley Dopita (2002) to distinguish ΟΠbetween the upper (og([NIT]/[OII]) 2 -1.2 and lower (logt[NIT|/|OTI]) « -1.2) Ros diagnostics.," The $_{23}$ diagnostic is double-valued, so we use the [NII]/[OII] criterion of Kewley Dopita (2002) to distinguish between the upper (log([NII]/[OII]) $>$ -1.2 and lower (log([NII]/[OII]) $<$ -1.2) $_{23}$ diagnostics."
+ The NII|/[OII] ratio is available for five of our long CRB host ealaxies as well as all five regious of the CRB 060505 host ealaxy aud the eutire BCC sample., The [NII]/[OII] ratio is available for five of our long GRB host galaxies as well as all five regions of the GRB 060505 host galaxy and the entire BCG sample.
+ Iu the abseuce of the NIJ A6581 flux (GRD 010921 and GRB 990712). we distinguish between the branches using the ratio (Nagao ct al.," In the absence of the [NII] $\lambda$ 6584 flux (GRB 010921 and GRB 990712), we distinguish between the branches using the $\alpha$ ratio (Nagao et al."
+ 2006)., 2006).
+ When neither [NI| nor [NOII|/Ta[NOTI] is available (CRB 051121 aud CRB 050116). we calculate he metallicities for both Ros brauches.," When neither [NII] nor [NeIII] is available (GRB 051121 and GRB 050416), we calculate the metallicities for both $_{23}$ branches."
+ For comparison. we also calculate the metallicities using the Pettini Pagel (2001) /Ue-metallicity relation. for our five ong GRB hosts with |NIIJAG583 line fluxes.," For comparison, we also calculate the metallicities using the Pettini Pagel (2004) $\alpha$ -metallicity relation for our five long GRB hosts with $\lambda$ 6583 line fluxes."
+ We determine the ionization parameter using the IKkewley Dopita (2002) relation., We determine the ionization parameter using the Kewley Dopita (2002) $q$ relation.
+ For he five regious of GRD 060505's host |OITI|/|JOTI|-4galaxy. we also calculate the age of the vouug (< 10 Mx) stellar »»pulatiou using the calibration of ITJ equivalent width with age by Schaerer Vacca (1998).," For the five regions of GRB 060505's host galaxy, we also calculate the age of the young $<$ 10 Myr) stellar population using the calibration of $\beta$ equivalent width with age by Schaerer Vacca (1998)."
+ We caleulate the Ilo bhunuimosities and corresponding star formation rates (SFRs) uxiug the relation of IXeunicutt (1998) for the GRB 060505 host regions aud our long GRB hosts (IIa fluxes were unavailable for the short GRB hosts}., We calculate the $\alpha$ luminosities and corresponding star formation rates (SFRs) using the relation of Kennicutt (1998) for the GRB 060505 host regions and our long GRB hosts $\alpha$ fluxes were unavailable for the short GRB hosts).
+ The uctallicities. ionization parameters. Ho Iunuinosities. aud SERs are given in Table 1..," The metallicities, ionization parameters, $\alpha$ luminosities, and SFRs are given in Table \ref{tab:params}."
+ We find that the GRB 060505 site has a similar uctallicity to the other regious of its host salaxv: the Ros diagnostic assigus the GRB 060505 burst site an intermediate metallicity of los(O/II) | 12 = S457 + VOL as compared to the rest of the host ealaxy. with au average log(O/II) | 12 = 8.62 + 0.2.," We find that the GRB 060505 site has a similar metallicity to the other regions of its host galaxy; the $_{23}$ diagnostic assigns the GRB 060505 burst site an intermediate metallicity of log(O/H) + 12 = 8.57 $\pm$ 0.01 as compared to the rest of the host galaxy, with an average log(O/H) + 12 = 8.62 $\pm$ 0.2."
+ The Pettini Pageaa (2001). or PPOL. relation gives the GRB 060505 site a lower metallicity of los(O/II) | 12 = 8.28 + 0.01 than the average metallicity of the other regions (log(O/II) | 12 28.15 + 0.06). but the difference is within the 0.15 dex errors of the PPOL method.," The Pettini Pagel (2004), or PP04, relation gives the GRB 060505 site a lower metallicity of log(O/H) + 12 = 8.28 $\pm$ 0.01 than the average metallicity of the other regions (log(O/H) + 12 = 8.45 $\pm$ 0.06), but the difference is within the 0.15 dex errors of the PP04 method."
+ The relative metallicity of the CRB 060505 site to the other regions of the galaxy is dependent on the diagnostic that is used: when the errors of these diagnostics are considered (0.1-0.15 dex). we find no statistically significant difference m metallicity )otween these five regions.," The relative metallicity of the GRB 060505 site to the other regions of the galaxy is dependent on the diagnostic that is used; when the errors of these diagnostics are considered (0.1-0.15 dex), we find no statistically significant difference in metallicity between these five regions."
+ The ionization parameter of the GRD 060505 site. oe(g) = 7.19 cum/s. is uuusuallv low as compared witli average ionization parameter of the long GRD lost ealaxies (log(q) = 7.95 + 0.28 οπή. The ionization xuanmeter of the CRB 060505 site is consistent with he short GRD hosts. which have au average ionization xuanmeter of log(q) = 7.36 4 0.15 cms. The ages of the voune stellar populations are representative of core-collapse progenitor ages (see Bloom et al.," The ionization parameter of the GRB 060505 site, $q$ ) = 7.49 cm/s, is unusually low as compared with average ionization parameter of the long GRB host galaxies $q$ ) = 7.95 $\pm$ 0.28 cm/s. The ionization parameter of the GRB 060505 site is consistent with the short GRB hosts, which have an average ionization parameter of $q$ ) = 7.36 $\pm$ 0.15 cm/s. The ages of the young stellar populations are representative of core-collapse progenitor ages (see Bloom et al."
+ 2002. Berecr et al.," 2002, Berger et al."
+ 2007)., 2007).
+ While the CRB sites is the vouugest at 5.3 + 0.3. this age is comparable o the 6.1 + 0.6 age of the young stellar population of he upper bulge. to within the errors.," While the GRB site's is the youngest at 5.3 $\pm$ 0.3, this age is comparable to the 6.1 $\pm$ 0.6 age of the young stellar population of the upper bulge, to within the errors."
+ Finally. the GRB site does not have a hieh SER with respect to the other regions of the host galaxy.," Finally, the GRB site does not have a high SFR with respect to the other regions of the host galaxy."
+ The SER of the CRB site is also found to be notably lower than the SFRs of the long CRB host ealaxics., The SFR of the GRB site is also found to be notably lower than the SFRs of the long GRB host galaxies.
+ We compare the optical diagnostic enission-line ratios of GRB 060505 with the hosts of wnambienous loue- aud short- duration WWoe show that the enüssion-ine ratios. metallicities. jionizatiou parameters aud star ormnatioun rates of the GRB 060505 environment are nore consistent with the two short-«duration GRB hosts hat have measured optical enissiou-Iine ratios rhan with one-diration CRBs.," We compare the optical diagnostic emission-line ratios of GRB 060505 with the hosts of unambiguous long- and short- duration We show that the emission-line ratios, metallicities, ionization parameters and star formation rates of the GRB 060505 environment are more consistent with the two short-duration GRB hosts that have measured optical emission-line ratios rhan with long-duration GRBs."
+ We compare the metallicity aud SFR of the CRB 160505 star-forming region with four other star-forming regions in the GRB 060505 host galaxy., We compare the metallicity and SFR of the GRB 060505 star-forming region with four other star-forming regions in the GRB 060505 host galaxy.
+ We fud no sienificaut difference in either metallicity or SER between he CRB 060505 region or the other star-forming regious. including the host galaxy bulec.," We find no significant difference in either metallicity or SFR between the GRB 060505 region or the other star-forming regions, including the host galaxy bulge."
+ We do not find conmipolliug evidence to sugeest that GRD 060505 originated iu a loug-duratiou core-collapse progenitor., We do not find compelling evidence to suggest that GRB 060505 originated in a long-duration core-collapse progenitor.
+ Our emission line diagnostic analysis sugecsts that the environment of CRB 060505 is more consistent with the host environments of compact-olject-merecr CRB progenitors., Our emission line diagnostic analysis suggests that the environment of GRB 060505 is more consistent with the host environments of compact-object-merger GRB progenitors.
+ A larger comparison sample of short and lone GRBs with cussion line spectra may shed further Leht on the nature of CRB 060505 aud other iutermediate-duration gaunuivcray bursts., A larger comparison sample of short and long GRBs with emission line spectra may shed further light on the nature of GRB 060505 and other intermediate-duration gamma-ray bursts.
+Oxveen is usually the major coolant iu the oxvgen abundant PNs.,Oxygen is usually the major coolant in the oxygen abundant PNs.
+ But what other coolants play a role?, But what other coolants play a role?
+ Iu Figure 13 we show the intensity of the line relative to for the major coolants in the SAIC. LAIC. and Galactie PNs. versus the oxvecu abundauce.," In Figure 13 we show the intensity of the line relative to for the major coolants in the SMC, LMC, and Galactic PNs, versus the oxygen abundance."
+ The intensitics are from the hottest models for each egalaxian mix., The intensities are from the hottest models for each galaxian mix.
+ Cloudy predicts that although the 5007 line is the major coolaut for the Galactic and LAIC PNs. in the SMC PNs the ALS LS line is the major coolaut.," Cloudy predicts that although the 5007 line is the major coolant for the Galactic and LMC PNs, in the SMC PNs the $\lambda$ 1548 line is the major coolant."
+ Other carbon lines are also important coolaut in the SAIC PNs. while the A L959 aud A3727 lines are significaut mostly for the Galactic models.," Other carbon lines are also important coolant in the SMC PNs, while the $\lambda$ 4959 and $\lambda$ 3727 lines are significant mostly for the Galactic models."
+ Clearly. wore detailed models aud observations of A1909 and AL550 are essential to confirm these predictions. but the interpretation is plausible and consistent with the available data.," Clearly, more detailed models and observations of $\lambda$ 1909 and $\lambda$ 1550 are essential to confirm these predictions, but the interpretation is plausible and consistent with the available data."
+ Iu Staughellinietal.(2002) we used the (H007|1959) / ratio to trace the nebular excitation. and by inference the stellar teiiperature. of the LMC PNs.," In \citet{sta02} we used the (5007+4959) / ratio to trace the nebular excitation, and by inference the stellar temperature, of the LMC PNs."
+ We cau not perform the same analysis here. since it is clear from Figure 11 that the ratio versus temperature relation is very uon-onotonuic.," We can not perform the same analysis here, since it is clear from Figure 11 that the ratio versus temperature relation is very non-monotonic."
+ Table 6 lists nebular densities determined from strong AA6716.6731 lines for cight SMC PNs (this paper) aud twelve LAC PNs (Stauehellinietal.2002).," Table 6 lists nebular densities determined from strong $\lambda\lambda$ 6716,6731 lines for eight SMC PNs (this paper) and twelve LMC PNs \citep{sta02}."
+. We assumed T.—104 K for these estimates., We assumed $_{\rm e}$ $^4$ K for these estimates.
+ However. the derived densities are extremely insensitive to T. (Osterbrock1989).," However, the derived densities are extremely insensitive to $_{\rm e}$ \citep{ost89}."
+. We find that <Neaypo-23. 15. whereas <No Dpe—328. a factor 1.5 different.," We find that $<{\rm N}_{\rm e}>_{\rm SMC}$ =3.45, whereas $<{\rm N}_{\rm e}>_{\rm LMC}=$ 3.28, a factor 1.5 different."
+ Caven the ια sample size. a factor of 1.5 is not likely to be very significant.," Given the small sample size, a factor of 1.5 is not likely to be very significant."
+ Furthermore. there is no discernible trend of N. with morphological type in SAIC or LMC PNs.," Furthermore, there is no discernible trend of $_{\rm e}$ with morphological type in SMC or LMC PNs."
+ Moreover. nebulae with bright lines are strongly biased to those of gcucrally high fiuxes or surface brightness. or low ionization.," Moreover, nebulae with bright lines are strongly biased to those of generally high fluxes or surface brightness, or low ionization."
+ These selection biases render trends in NC of limited significance., These selection biases render trends in $_{\rm e}$ of limited significance.
+ Iu addition we estimate masses of PNs in the LMC and SAIC using the method of (1991)..., In addition we estimate masses of PNs in the LMC and SMC using the method of \citet{bs94}.
+ The average ionized mass of the eight SAIC PNs. 0.3AL... is slightly larecr than that of LAIC PNs. 0.2AL.," The average ionized mass of the eight SMC PNs, 0.3, is slightly larger than that of LMC PNs, 0.2."
+.. Given the uncertainty in the data. and the simall data sample. we do uot believe that the mass discrepancy between SAIC aud LMC PNs is significant.," Given the uncertainty in the data, and the small data sample, we do not believe that the mass discrepancy between SMC and LMC PNs is significant."
+ Also. we fail to see auy obvious treud between uebular mass aud morphological type.," Also, we fail to see any obvious trend between nebular mass and morphological type."
+ A sample of 27 SMC PNs have been observed in imaging a slitless mode with Z$T/STIS to exiiuiue heim morphology. shape. aud fluxes. aud to study their evolution.," A sample of 27 SMC PNs have been observed in imaging a slitless mode with /STIS to examine their morphology, shape, and fluxes, and to study their evolution."
+ This morphological sample is the first sizable set of SAIC) PNs. aud represents almost half of the kuown SAIC PNs.," This morphological sample is the first sizable set of SMC PNs, and represents almost half of the known SMC PNs."
+ The images and spectra lave he same high quality aud resolution as our LMC sample., The images and spectra have the same high quality and resolution as our LMC sample.
+ We preseut the broad images aud monochromatic Huaees in the major cussion lines. aud determine that morphology is easily recognized m dost eniüssion ines.," We present the broad images and monochromatic images in the major emission lines, and determine that morphology is easily recognized in most emission lines."
+ We find that the ratio of svuuuetric to asviunetrice PNs is remarkably different iu the SAIC aud the LMC., We find that the ratio of symmetric to asymmetric PNs is remarkably different in the SMC and the LMC.
+" Specifically, bipolar PNs are much rarer in the SAIC than the LAIC (or the Galaxy)."," Specifically, bipolar PNs are much rarer in the SMC than the LMC (or the Galaxy)."
+ This new result las significant iuplicatious for the relation between stellar population aud PN morphology., This new result has significant implications for the relation between stellar population and PN morphology.
+ It is well known roni Galactic aud. LAIC PN studies that PN morphology correlates with the mass of the progenitor stars., It is well known from Galactic and LMC PN studies that PN morphology correlates with the mass of the progenitor stars.
+ Iu particular. bipolar PNs evolve frou relatively massive (21.5 M.) progenitors.," In particular, bipolar PNs evolve from relatively massive $\ge 1.5$ ) progenitors."
+ Thus the low incidence of bipolar PNs iu the SAIC probably reflects the low formation rate of these stars in the past 5 Car., Thus the low incidence of bipolar PNs in the SMC probably reflects the low formation rate of these stars in the past $\ge$ 5 Gyr.
+ This, This
+between the HLEODO RAIS anc energy (see Boivin ct al.,between the HFQPO RMS and energy (see Boirin et al.
+ 1998 for more details)., 1998 for more details).
+ Similar correlations have been observed. in other sources (Van cer Ixlis 1998)., Similar correlations have been observed in other sources (Van der Klis 1998).
+ at l., Looking at Fig.
+ one can see that HPQPOs were not. detected in observationsLooking with Fig.the largest count rates.," 1, one can see that HFQPOs were not detected in observations with the largest count rates."
+ We have derived a 30 upper, We have derived a $3\sigma$ upper
+ lle and Ma (2011)) have shown that it is possible to consider black holes as cuantized objects.,	 He and Ma ) have shown that it is possible to consider black holes as quantized objects.
+ Our approach here is to apply this idea to the universe., Our approach here is to apply this idea to the universe.
+ This is in line with the previous work that considers the universe to be a black hole 2010))., This is in line with the previous work that considers the universe to be a black hole ).
+" We find here (he equivalent quantum number for the universe aud present the bit as the euantum of the eravitational potential. as already noted in (2011)aa. This munber is close to 107,"," We find here the equivalent quantum number for the universe and present the bit as the quantum of the gravitational potential, as already noted in a. This number is close to $10^{122}$."
+ We note that this big munber is close to. if not the same. the worrving discrepancy that one has between the cosmological constant value. [rom cosmological observations. ancl the value obtained [rom the standard theory of particles 2009)).," We note that this big number is close to, if not the same, the worrying discrepancy that one has between the cosmological constant value, from cosmological observations, and the value obtained from the standard theory of particles )."
+ A likely explanation for (his discrepancy has also been advanced by Santos (2010aa. 2010bb and 2011)).," A likely explanation for this discrepancy has also been advanced by Santos a, b and )."
+ This strongly suggests that our universe can be considered to be an excited state of Planck's quantum black hole., This strongly suggests that our universe can be considered to be an excited state of Planck's quantum black hole.
+ Again. (his is reinforced by the fact that the natural units. Planck's units. when multiplied bv the big number 1095. eive the well known physical properties of the universe. mass. length and time.," Again, this is reinforced by the fact that the natural units, Planck's units, when multiplied by the big number $10^{61}$, give the well known physical properties of the universe, mass, length and time."
+ Also the entropy of the universe. as a black hole. is linked to this equantum number being of the order of the square of it. nearly the same as the estimation mace by Eagan ancl Lineweaver (2010)).," Also the entropy of the universe, as a black hole, is linked to this quantum number being of the order of the square of it, nearly the same as the estimation made by Eagan and Lineweaver )."
+ Finally we reler to, Finally we refer to
+where the positive (negative) root occurs if the particle is recediug from (approaching) tle observer.,where the positive (negative) root occurs if the particle is receding from (approaching) the observer.
+" As stated before. we cousider only the case in which the particle is moving along the branch of the hyperbola for which or),>0."," As stated before, we consider only the case in which the particle is moving along the branch of the hyperbola for which $x_{\mathrm{p}}>0$."
+ With this in mine. it follows from equation (17)) that which also implies that 0«/<x.," With this in mind, it follows from equation \ref{eq:|xp|}) ) that which also implies that $0<t<\infty$."
+" The apparent velocity of the particle is given by We. therefore.see that asf— x. ey1—c/2 and as /—0. c0,1——x."," The apparent velocity of the particle is given by We, therefore,see that as $t\rightarrow\infty$ , $v_{\mathrm{p}}\rightarrow c/2$ and as $t\rightarrow0^{+}$, $v_{\mathrm{p}}\rightarrow-\infty$."
+ This is consistent witl results following from equations (12) and (5)), This is consistent with results following from equations \ref{eq:x_p>0}) ) and \ref{eq:x_p<0}) ).
+ We now return to the situation in which the particle moves along the trajectory of equation (13)). but we now position our observer at (£.0.6) in liis rest frauue. rather than at (0.0.0).," We now return to the situation in which the particle moves along the trajectory of equation \ref{eq:hyperbola}) ), but we now position our observer at $\left(\xi,0,\zeta\right)$ in his rest frame, rather than at $(0,0,0)$."
+ If again. the particle emits a light signal at time 7 when it is at position Cep.0.0). we find that [rom equation (2)) where / is the time the signal arrives at (£.0. 6).," If again, the particle emits a light signal at time $\tau$ when it is at position $\left(x_{\mathrm{p}},0,0\right)$, we find that from equation \ref{eq:distance}) ) where $t$ is the time the signal arrives at $\left(\xi,0,\zeta\right)$ ."
+ From equation (13)) we fined that As a result.," From equation \ref{eq:hyperbola})) we find that so that equation \ref{eq:dist_gen_c}) ) becomes As a result,"
+'The period ratios between different QPPs may be meaningful.,The period ratios between different QPPs may be meaningful.
+" Around the first LPP shown in Fig.10, there occurred 3 classes of QPPs: VLP (paragraph A), LPP and SPP."," Around the first LPP shown in Fig.10, there occurred 3 classes of QPPs: VLP (paragraph A), LPP and SPP."
+" Their averaged periods are 4.13 min, 70 s, and 8 s, respectively."," Their averaged periods are 4.13 min, 70 s, and 8 s, respectively."
+ Their ratios are 31.0: 8.75: 1., Their ratios are 31.0: 8.75: 1.
+" The second LPP is just occurred in the gap between paragraph A and B of VLP, and during this gap we did not distinguish other kind of QPP."," The second LPP is just occurred in the gap between paragraph A and B of VLP, and during this gap we did not distinguish other kind of QPP."
+" During the third LPP, there occurred VLP (paragraph A, P=4.5 min), LPP (P=15 s) and VSP (P=90~110 ms)."," During the third LPP, there occurred VLP (paragraph A, $P=4.5$ min), LPP $P=15$ s) and VSP $P=90\sim110$ ms)."
+" The period ratio is about 2700: 150: 1, which is entirely different from the former ratios."," The period ratio is about 2700: 150: 1, which is entirely different from the former ratios."
+ Fig.11 shows an example of concurrence of two different classes of QPPs occurred at 02:38:50 — 02:40:00 UT., Fig.11 shows an example of concurrence of two different classes of QPPs occurred at 02:38:50 – 02:40:00 UT.
+" The higher hierarchy of QPP is a LPP with period of 16 s, and the lower hierarchy of QPP is a fast-VSP with period of 77 ms."," The higher hierarchy of QPP is a LPP with period of 16 s, and the lower hierarchy of QPP is a fast-VSP with period of 77 ms."
+" From this figure we find that the LPP is fine-structured with a train of fast-VSPs, and the period ratio is more than 200."," From this figure we find that the LPP is fine-structured with a train of fast-VSPs, and the period ratio is more than 200."
+" And in the train of fast-VSPs, we find that there is obviously global frequency drifting rate with value of about 2.5 MHz/s, and this drift rate is very slow comparing with the single pulse frequency drifting rates."," And in the train of fast-VSPs, we find that there is obviously global frequency drifting rate with value of about 2.5 MHz/s, and this drift rate is very slow comparing with the single pulse frequency drifting rates."
+ These evidences imply that there is no obvious regulations among the period ratios and the weakly links between the different classes of QPPs., These evidences imply that there is no obvious regulations among the period ratios and the weakly links between the different classes of QPPs.
+" Additionally, from careful investigation of the microwave observations there is no evidence of any LPP, SPP and VSP found around the time intervals of VLP paragraph C and D. It seems that all the LPP, SPP, and VSP take place around the time intervals of VLP paragraph A and B. Usually, the microwave burst can be regarded as a prompt signal of non-thermal energetic particles originating from the magnetic reconnection."," Additionally, from careful investigation of the microwave observations there is no evidence of any LPP, SPP and VSP found around the time intervals of VLP paragraph C and D. It seems that all the LPP, SPP, and VSP take place around the time intervals of VLP paragraph A and B. Usually, the microwave burst can be regarded as a prompt signal of non-thermal energetic particles originating from the magnetic reconnection."
+ So it is meaningful to quest for the relationships between the magnetic field and QPPs., So it is meaningful to quest for the relationships between the magnetic field and QPPs.
+" Here, we introduce the transport rate of magnetic helicity (dH/dt) to describe the magnetic field behavior."," Here, we introduce the transport rate of magnetic helicity $dH/dt$ ) to describe the magnetic field behavior."
+" It can be expressed as (Berger Field, 1984): al—§2(B-Ap)vzdS—¢ 2(v-Ap) "," It can be expressed as (Berger Field, 1984): $\frac{dH}{dt}=\oint2(\textbf{\textit{B}}\cdot\textbf{\textit{A}}_{\textrm{p}})\upsilon_{\textrm{z}}dS-\oint2(\mbox{\boldmath
+$ $}\cdot\textbf{\textit{A}}_{\textrm{p}})B_{\textrm{z}}dS$ ."
+B," Here, $\textbf{\textit{A}}_{\textrm{p}}$ is magnetic vector potential, ${\upsilon}$ is the fluid velocity, and $B_{\textrm{z}}$ is the normal component of the magnetic field."
+zdS.," In an open volume, the magnetic helicity may change by the passage of helical magnetic field lines through the surface (the first term) and by shuffling horizontal motion of field lines on the surface (the second term)."
+" Here, "," In an isolated active region $dH/dt$ is a nonpotential parameter which indicates the dynamic evolution of the magnetic field mainly in the flaring region (Berger Field, 1984)."
+Ap," As the flare/CME event occurred in an isolated active region of AR 10930, we may estimate the quantity for the event by calculating from the above expression."
+ is ma," The work of Zhang et al (2008) presents that the temporal profile of $dH/dt$ is consistent with that of microwave burst in the flare/CME event on Dec.13, 2006."
+g, So we superposed the curve of $dH/dt$ in the same duration in Fig.10 (dashed-dotted curve).
+netic vecto," Here, the cadence of dH/dt is 2 minutes, which is calculated from the observations of the Solar Optical Telescope on board Hinode (SOT/Hinode, Tsuneta et al, 2008; Kosugi et al, 2007)."
+r potenti," From the comparison of $dH/dt$ and QPPs, we find that when QPPs take place $dH/dt$ is positive."
+al, This fact indicates a continuous injection of magnetic helicity.
+, This may lead to accumulation of nonpotential energy and the magnetic reconnections in the flaring active region.
+, There are some variations in the profile of $dH/dt$ during the flaring event.
+", "," However, as the cadence of $dH/dt$ is only 2 minutes, we could not confirm any pulsating features in $dH/dt$."
+"v is the fluid velocity, and Bz is "," Actually, we find that VLP, LPP and SPP have much of similarities, such as broad emission bandwidth, almost at the same level of Q-factor $\leq10$ ), weakly circular polarization, and relatively slow frequency drift rates, except that they have different durations, periods, and different magnitude of pulsating emission flux."
+t, They belong possibly to a same class of QPP in a sense.
+he 1," In order to confirm such viewpoint, we plot all the QPPs in a logarithmic period-duration space in Fig.12."
+," We find that all of the VLPs, LPPs, SPPs, and part of slow-VSPs are distributed around a line."
+," However, almost all fast-VSPs and most part of the slow-VSPs are distributed far away from the above line dispersively."
+," This fact implies that fast-VSPs and part of slow-VSPs may be different originally from VLPs, LPPs, and SPPs."
+," Hereby, we may classify all QPPs into two groups: group I includes VLPs, LPPs, SPPs, and part of slow-VSPs of which the longer the duration corresponds to the longer period; group II includes fast-VSPs and most part of slow-VSPs of which the period is dispersive respect to the duration."
+," Table 2 presents a brief summary of VLP, LPP, SPP, slow-VSP, and fast-VSP."
+" From above investigations, we find that there are 5 classes of QPPs with different timescales associated with the flare/CME event: VLP, LPP, SPP, slow- and fast-VSP."," From above investigations, we find that there are 5 classes of QPPs with different timescales associated with the flare/CME event: VLP, LPP, SPP, slow- and fast-VSP."
+" They form a broad hierarchy of timescales of hecto-second, deca-second, a few seconds, deci-second, and centi-second."," They form a broad hierarchy of timescales of hecto-second, deca-second, a few seconds, deci-second, and centi-second."
+ And this is similar to the discovery of microwave burst timescales by Kruger et al (1994)., And this is similar to the discovery of microwave burst timescales by Kruger et al (1994).
+" They classified the microwave burst timescales as tens of minutes, a few minutes, a few seconds and sub-seconds which represent to main burst phase, main burst pulse, subpulse and spiky burst elements, respectively."," They classified the microwave burst timescales as tens of minutes, a few minutes, a few seconds and sub-seconds which represent to main burst phase, main burst pulse, subpulse and spiky burst elements, respectively."
+" But, these bursts have no obvious periodicities."," But, these bursts have no obvious periodicities."
+" However, our investigations indicate that the most remarkable feature of QPPs is the quasi-periodicity of the repetitive pulses."," However, our investigations indicate that the most remarkable feature of QPPs is the quasi-periodicity of the repetitive pulses."
+" Then, what is the generation mechanism of the above various classes of QPPs?"," Then, what is the generation mechanism of the above various classes of QPPs?"
+" It is well known that harmonic motion of a classic mechanic oscillator is due to a restoring force proportional to the displacement, and any elastic body can be excited to oscillate in eigen-modes."," It is well known that harmonic motion of a classic mechanic oscillator is due to a restoring force proportional to the displacement, and any elastic body can be excited to oscillate in eigen-modes."
+" When a magnetic field presents in plasma systems, the characteristic eigen-frequencies will be determined by the magnetic field, plasma density, temperature, and geometrical configurations of the magnetized plasma system."," When a magnetic field presents in plasma systems, the characteristic eigen-frequencies will be determined by the magnetic field, plasma density, temperature, and geometrical configurations of the magnetized plasma system."
+" Aschwanden (1987) presented an extensive review of pulsation models and classified them as three groups: (1) MHD flux tube oscillations, which modulate the radio emissivity with a standing or propagating MHD waves, e.g. slow magnetoacoustic mode, fast kink mode, fast sausage mode (Roberts, Edwin Benz, 1984; Nakariakov Melnikov, 2009). ("," Aschwanden (1987) presented an extensive review of pulsation models and classified them as three groups: (1) MHD flux tube oscillations, which modulate the radio emissivity with a standing or propagating MHD waves, e.g. slow magnetoacoustic mode, fast kink mode, fast sausage mode (Roberts, Edwin Benz, 1984; Nakariakov Melnikov, 2009). ("
+"2) Periodic self-organizing systems of plasma instabilities of wave-particle or wave-wave interactions interlocked by a Lotka-Volterra type of coupled equations (Aschwanden Benz, 1988). (","2) Periodic self-organizing systems of plasma instabilities of wave-particle or wave-wave interactions interlocked by a Lotka-Volterra type of coupled equations (Aschwanden Benz, 1988). ("
+"3) Modulation of periodic acceleration (repetitive injection of particles into the emission source region) which may possibly generatedfrom repetitive magnetic reconnections, for example, the pulsed acceleration in","3) Modulation of periodic acceleration (repetitive injection of particles into the emission source region) which may possibly generatedfrom repetitive magnetic reconnections, for example, the pulsed acceleration in"
+"(BZ) process in the force-free magnetosphere, realistic conversion mechanisms from the magnetic energy to the other form are still not clear, because in the force-free approximation the inertia of matters is ignored.","(BZ) process in the force-free magnetosphere, realistic conversion mechanisms from the magnetic energy to the other form are still not clear, because in the force-free approximation the inertia of matters is ignored."
+ We are now discussing MHD accretion onto a black hole., We are now discussing MHD accretion onto a black hole.
+" The MHD inflow takes the fluid energy into the hole (that is, the kinetic energy, the internal energy and the rest mass energy of the fluid), which is positive at the plasma source."," The MHD inflow takes the fluid energy into the hole (that is, the kinetic energy, the internal energy and the rest mass energy of the fluid), which is positive at the plasma source."
+" However, the situation of energy extraction from a rotating black hole can be also described in the MHD magnetosphere."," However, the situation of energy extraction from a rotating black hole can be also described in the MHD magnetosphere."
+" In this case, the MHD inflow (E« 0) is realized, where it is required that the point of the flow considered is located close to the inner light surface in addition to the condition for the BZ process; both the point and the inner light surface need to posit inside the ergosphere (Takahashietal.1990)."," In this case, the MHD inflow $E<0$ ) is realized, where it is required that the point of the flow considered is located close to the inner light surface in addition to the condition for the BZ process; both the point and the inner light surface need to posit inside the ergosphere \citep{TNTT90}."
+". By considering negative energy MHD inflows, the rotational energy of the black hole can be carried in the surrounding magnetosphere as the outward energy flux to the distant region directly (McKinney&Gammie2004;McKinney2006) or indirectly by way of an equatorial"," By considering negative energy MHD inflows, the rotational energy of the black hole can be carried in the surrounding magnetosphere as the outward energy flux to the distant region directly \citep{MG04,McKinney06} or indirectly by way of an equatorial"
+distributed cross-bar are greatly reduced in comparison.,distributed cross-bar are greatly reduced in comparison.
+ The WIDAR signal processing path is as follows., The WIDAR signal processing path is as follows.
+ 96 Gbps of sampled. pavload data arrive from each antenna via optical fiber., 96 Gbps of sampled payload data arrive from each antenna via optical fiber.
+ As described in subsection 3.3.. the data are sent to 4 SCDs per antenna via (he deformatter (the fiber optic receiver module (FORM) shown in Figure 6).," As described in subsection \ref{sub:dts}, the data are sent to 4 StBs per antenna via the deformatter (the fiber optic receiver module (FORM) shown in Figure \ref{fig:widar}) )."
+ An FPGA cross-bar allows (these signals to be routed (o two G4-bit. wide. 256 MlIIz data paths on the board.," An FPGA cross-bar allows these signals to be routed to two 64-bit wide, 256 MHz data paths on the board."
+ Coarse delay. tracking to within zone-half a sample is implemented with an FPGA and DDR. SDRAM on a delay module (DM) mezzanine card on the SUD. Baseline clelay al the level of -zEone-sixteenth a sub-saimple is implemented by continuous. real-time tracking of the residual delay error at (he center of each sub-band using the correlator chip's phase rotators [29].," Coarse delay tracking to within $\pm$ one-half a sample is implemented with an FPGA and DDR SDRAM on a delay module (DM) mezzanine card on the StB. Baseline delay at the level of $\pm$ one-sixteenth a sub-sample is implemented by continuous, real-time tracking of the residual delay error at the center of each sub-band using the correlator chip's phase rotators [29]."
+ The delav-corrected signals are then distributed to dual. IS-EPGA filter banks (FB).," The delay-corrected signals are then distributed to dual, 18-FPGA filter banks (FB)."
+ Each FB FPGA ean further delav-correct the data to allow lor a sub-band-specifie beam delay-center offset on the skv., Each FB FPGA can further delay-correct the data to allow for a sub-band-specific beam delay-center offset on the sky.
+ Each FPGA also contains four. 512-tap FIR stages with 16-bits carried between stages.," Each FPGA also contains four, 512-tap FIR stages with 16-bits carried between stages."
+ One or all of the stages can be used to produce sub-bands from 128 MlIIZ wide to as narrow as 31.25 kIIz., One or all of the stages can be used to produce sub-bands from 128 MHz wide to as narrow as 31.25 kHz.
+ Data exit the MB into the distributed cross-bar and is then routed to the BIBs., Data exit the StB into the distributed cross-bar and is then routed to the BlBs.
+ High-density. harchnetric connectors and mating cabling allow for 1.024 Gbps per pair of BIBs. fitting the numerology of the correlator nicely.," High-density, hardmetric connectors and mating cabling allow for 1.024 Gbps per pair of BlBs, fitting the numerology of the correlator nicely."
+ The pair of DIBs processes all 32 antennas (496 baselines) for a pair (2 x 128 MIIz). with 64 Ghps of astronomical data ancl up to 32 Gbps of control and Gmiug data into each BID. Signals are synchronized using dvnamic phase alignment. ancl distributed by two eross-bar and phasing FPGAs to the 3x8 array of NF correlator chips.," The pair of BlBs processes all 32 antennas (496 baselines) for a sub-band pair (2 x 128 MHz), with 64 Gbps of astronomical data and up to 32 Gbps of control and timing data into each BlB. Signals are synchronized using dynamic phase alignment, re-timed, and distributed by two cross-bar and phasing FPGAs to the 8x8 array of XF correlator chips."
+ The F672 BGA XF correlator chip is a 4 Mgate standard-cell device fabricated in 130 nm CMOS., The F672 BGA XF correlator chip is a 4 Mgate standard-cell device fabricated in 130 nm CMOS.
+ Each device contains 2048 lags”. each of which consists of a 4-bit multiplier. a 3-level phase rotator. and dual 22- accumulators.," Each device contains 2048 lags"", each of which consists of a 4-bit multiplier, a 3-level phase rotator, and dual 22-bit accumulators."
+ The phase rotators perform phase tracking (fringe rotation). sub-sample delay correction. and frequency. shift removal.," The phase rotators perform phase tracking (fringe rotation), sub-sample delay correction, and frequency shift removal."
+ All signals lor integration control originate in the SUD. ancl integrations can be synchronized (o svstem timing or any timing model. such as a pulsar ephemeris.," All signals for integration control originate in the StB, and integrations can be synchronized to system timing or any timing model, such as a pulsar ephemeris."
+ Data from each correlator chip are integrated further in a LTA FPGA and DDR SDRAM., Data from each correlator chip are integrated further in a low-cost LTA FPGA and DDR SDRAM.
+ The final results are contained in UDP/IP packets. each containing 128 complex-lag accumulation results.," The final results are contained in UDP/IP packets, each containing 128 complex-lag accumulation results."
+ The packets are sent to the correlator computing cluster on 1 Gigabit Ethernet via a commercial switch., The packets are sent to the correlator computing cluster on 1 Gigabit Ethernet via a commercial switch.
+ The cluster carries oul the data normalization. Fourier transform. ancl interference excision.," The cluster carries out the data normalization, Fourier transform, and interference excision."
+ The resultant data are sent to an archive for final image processing., The resultant data are sent to an archive for final image processing.
+ The crosshar FPGAs also perform the phasing of the antenna array., The crossbar FPGAs also perform the phasing of the antenna array.
+ The phased signal is output to several destinations. primarily Ethernet packets to a VLBI data recorder.," The phased signal is output to several destinations, primarily Ethernet packets to a VLBI data recorder."
+ The correlator also provides four. wideband. auto-correlation products for every baseband pair.," The correlator also provides four, wideband, auto-correlation products for every baseband pair."
+ Each data product has 1024 spectral channels. but with a factor of four sensitivity loss," Each data product has 1024 spectral channels, but with a factor of four sensitivity loss"
+To estimate the interstellar extinction to ILD 6833. we exploited (he spectroscopic study by Fulbright (2000). who derived Ti; = 4450 Ix. log g = 1.1. and |Fe/II] = —1.04.,"To estimate the interstellar extinction to HD 6833, we exploited the spectroscopic study by Fulbright (2000), who derived $T_{\rm eff}$ = 4450 K, log $g$ = 1.4, and [Fe/H] = $-1.04$."
+ The star has Y= 6.75 and B—V = 1.17. and its metallicity is a neaur-perfect match for that of NGC 6723 ([Fe/II] = —1.03 να Ivans 2003).," The star has $V$ = 6.75 and $B-V$ = 1.17, and its metallicity is a near-perfect match for that of NGC 6723 ([Fe/H] = $-1.03$ Kraft Ivans 2003)."
+ To obtain the same elfective temperature from the Alonso et ((1999) calibrations requires E(B—V) = 0.11 [E(5—y) = 0.076). so (D—V) = 1.06.," To obtain the same effective temperature from the Alonso et (1999) calibrations requires $B-V$ ) = 0.11 $b-y$ ) = 0.076], so $B-V$ $_{0}$ = 1.06."
+ The available color-magnitude diagram of the cluster reveals that Mqο—0.4 al that color. which we adopt.," The available color-magnitude diagram of the cluster reveals that $_{\rm V} \approx\ -0.4$ at that color, which we adopt."
+ Note that as a result. our photometric estimate of log q is 1.6. quite similar to that obtained by Fulbright (2000). who derived a value of 1.4.," Note that as a result, our photometric estimate of log $g$ is 1.6, quite similar to that obtained by Fulbright (2000), who derived a value of 1.4."
+ In the case of IDE 232078. we relied on the analvsis by Burris et ((2000).," In the case of HDE 232078, we relied on the analysis by Burris et (2000)."
+ Thev determined the temperature and metallicity using spectroscopic methods. but. emploved photometry to estimate the gravity. finding Tar=4000. log y = +0.3. and |Fe/1l] = —1.54.," They determined the temperature and metallicity using spectroscopic methods, but employed photometry to estimate the gravity, finding $T_{\rm eff} = 4000$, log $g$ = +0.3, and [Fe/H] = $-1.54$."
+ These values enable us to determine the interstellar extinction. under the assumption that ihe Alonso et ((1999) color-temperature calibration should vield the same temperature if we have estimated the reddening correctly.," These values enable us to determine the interstellar extinction, under the assumption that the Alonso et (1999) color-temperature calibration should yield the same temperature if we have estimated the reddening correctly."
+ The globular cluster M3 has nearly identical metallicity ([Fe/H] = —1.47: Ixraft et 11992). and a very low reddening [E(J—V) = 0.01 mag]. so in principle we only need to use the Alonso et ((1999) relations between V—A and Toy to estimate temperatures of its red giants.," The globular cluster M3 has nearly identical metallicity ([Fe/H] = $-1.47$; Kraft et 1992), and a very low reddening $B-V$ ) = 0.01 mag], so in principle we only need to use the Alonso et (1999) relations between $V-K$ and $T_{\rm eff}$ to estimate temperatures of its red giants."
+ Then we simply identily (V— A) for stus with the same effective temperature as IIDIEE 232078., Then we simply identify $V-K$ $_{0}$ for stars with the same effective temperature as HDE 232078.
+ Using available unpublished M3 DV photometry and 2MAÀSS A magnitudes. we obtain E(B— V) = 0.56 mag for WDE 232073. as well as My = —2.15.," Using available unpublished M3 $BV$ photometry and 2MASS $K$ magnitudes, we obtain $B-V$ ) = 0.56 mag for HDE 232078, as well as $M_{\rm V}$ = $-2.15$."
+ Our photometric estimate for the exavityv is log g = 40.6. somewhat higher than that estimated by Burris et ((2000).," Our photometric estimate for the gravity is log $g$ = +0.6, somewhat higher than that estimated by Burris et (2000)."
+ We relv on the photometric estimation of My. values and bolometric corrections {ο estimate (he eravilies of our program stars. so we must ask how accurate such estimates are. especially eiven the relatively large distances to our program stars.," We rely on the photometric estimation of $M_{\rm V}$ values and bolometric corrections to estimate the gravities of our program stars, so we must ask how accurate such estimates are, especially given the relatively large distances to our program stars."
+ We have (wo methods available: trigonometric parallaxes from HIPPARCOS. and spectroscopic gravities derived during the course of metallicity determinations.," We have two methods available: trigonometric parallaxes from HIPPARCOS, and spectroscopic gravities derived during the course of metallicity determinations."
+ Figure 1. plots My. values estimated by ATT and. in the special cases here and in C2003. by us. against (those derived directly. from trigonometric parallaxes.," Figure \ref{fig1} plots $M_{\rm V}$ values estimated by ATT and, in the special cases here and in C2003, by us, against those derived directly from trigonometric parallaxes."
+ The lo error bars are almost all very. large. since (vpical parallaxes are only one to two milli-arcseconds. and the uncertainties are often. comparable to or even lareer than the parallaxes. so that (he formal error encompasses physically unrealistic negative parallaxes.," The $\sigma$ error bars are almost all very large, since typical parallaxes are only one to two milli-arcseconds, and the uncertainties are often comparable to or even larger than the parallaxes, so that the formal error encompasses physically unrealistic negative parallaxes."
+ We have not plotted, We have not plotted
+=O010.02 7.,$\approx 0.01-0.02$ $^{-3}$.
+ Lo illustrate the ellect that such a constant density core would have in terms of the rotation curve of the MW. we artificially smooth the inner density profile of the MW dark halo model along its evolution in such away that at the present epoch it has such a core.," To illustrate the effect that such a constant density core would have in terms of the rotation curve of the MW, we artificially smooth the inner density profile of the MW dark halo model along its evolution in such a way that at the present epoch it has such a core."
+ The final rotation curve decomposition obtained after the formation of a MW disc in this halo is shown in Fig., The final rotation curve decomposition obtained after the formation of a MW disc in this halo is shown in Fig.
+ 2a by the thin curves. which are analogous to the thick curves of the NEW case.," 2a by the thin curves, which are analogous to the thick curves of the NFW case."
+ The rotation curve is now almost completely [lat and he disc component dominates over the halo component up o the solar radius. i.c. the dashed ancl clotted curves now intersect at approximately the solar radius. as direct studies of our galaxy show (c.g. Ixuijken Cilmore 1989).," The rotation curve is now almost completely flat and the disc component dominates over the halo component up to the solar radius, i.e., the dashed and dotted curves now intersect at approximately the solar radius, as direct studies of our galaxy show (e.g. Kuijken Gilmore 1989)."
+ We can renee conclude that a comparison of dynamical studies in he MW and galactic evolutionary mocels olfers vet another evidence of the excessive central. concentration of CDM idoes predicted by the standard. cosmology., We can hence conclude that a comparison of dynamical studies in the MW and galactic evolutionary models offers yet another evidence of the excessive central concentration of CDM haloes predicted by the standard cosmology.
+ Fortunately. he inclusion of a shallow core in the halo has no major consequences upon the SELL. as will be shown below.," Fortunately, the inclusion of a shallow core in the halo has no major consequences upon the SFH, as will be shown below."
+ As it ds seen in able 1. the properties of the MW fiducial model predict reasonably well the data inferred rom. observations. with the exception of the rotation curve.," As it is seen in Table 1, the properties of the MW fiducial model predict reasonably well the data inferred from observations, with the exception of the rotation curve."
+ In Fig., In Fig.
+ 3a we show the evolution of the SER. per unit area and of the gas infall rate. both at the radius &y=8.5 kpe. MoCo.1) (solid line) and. NM{τοςϱ) (point-dashed line). respectively.," 3a we show the evolution of the SFR per unit area and of the gas infall rate, both at the radius $R_0=8.5$ kpc, $\sfr (R_0,t)$ (solid line) and $\infrate (R_0,t)$ (point-dashed line), respectively."
+ ‘The disc at the radius £y begins to form at à look-back time 11 Gye (22), The disc at the radius $R_0$ begins to form at a look-back time $\sim 11$ Gyr $z\approx 2$ ).
+ An interesting test derives from taking into account an integrated colour index for the solar neigibourhood., An interesting test derives from taking into account an integrated colour index for the solar neighbourhood.
+ Using recent Padova stellar evolutionary mocels (e.g. Girardi et al., Using recent Padova stellar evolutionary models (e.g. Girardi et al.
+ 1996) we computed svnthetic colour. magnitude ciagrams for the average SELL shown by the solid curve in Fig.," 1996) we computed synthetic colour magnitude diagrams for the average SFH, shown by the solid curve in Fig."
+ Sa. taking a solar metallicity for the last 2 Gyr. and one third solar before this point. as a first approximation to the enrichment history.," 3a, taking a solar metallicity for the last 2 Gyr, and one third solar before this point, as a first approximation to the enrichment history."
+ C'aleulating the DAx. colour index for realizations contzuning upwards of 200.000 stars we obtain a value of D.IK=3.15. in excellent agreement with the solar netghbourheος observational estimate for this quantity of D.WK=3.13 (Binney Morrifield. 1999).," Calculating the $B-K$ colour index for realizations containing upwards of 200,000 stars we obtain a value of $B-K=3.15$, in excellent agreement with the solar neighbourhood observational estimate for this quantity of $B-K=3.13$ (Binney Merrifield 1999)."
+ Repeating he experiment or SEL resulting from ALALIs cleviating rom the average one. we obtain practically identical values or the dy band uminosity. reflecting equal integrals uncer he SELL curves. out rather different values for the B. band unminositv. rellec‘tine dillerent recent Skis.," Repeating the experiment for SFHs resulting from MAHs deviating from the average one, we obtain practically identical values for the $K-$ band luminosity, reflecting equal integrals under the SFH curves, but rather different values for the $B-$ band luminosity, reflecting different recent SFHs."
+ Taking the ALALIS which clifer most from the average (of the ones eiven in Fig., Taking the MAHs which differ most from the average (of the ones given in Fig.
+ Ib). we obtain dillerences of 0.3 mag in the oedicted 2dy of the solar neighbourhood. rather larger han the observaional uncertainties of the measured. value.," 1b), we obtain differences of 0.3 mag in the predicted $B-K$ of the solar neighbourhood, rather larger than the observational uncertainties of the measured value."
+ This important result leads us to expect the average ALALL as the optimal choice for the Galaxy., This important result leads us to expect the average MAH as the optimal choice for the Galaxy.
+ The recent. observational inferences of Binney. Dehen Bertelli (2000 henceforth DBDOO) use kinematical cata from the Llipparcos catalogue to break the age metallicity degeneracy at the oldest. turn olf. identified: also from the Lipparcos satellite by combining photometric data with theoretical isochrones.," The recent observational inferences of Binney, Dehen Bertelli (2000 henceforth DBD00) use kinematical data from the Hipparcos catalogue to break the age metallicity degeneracy at the oldest turn off, identified also from the Hipparcos satellite by combining photometric data with theoretical isochrones."
+ Their results for the age of the solar neighbourhood are shown by the shaded bar in Fig., Their results for the age of the solar neighbourhood are shown by the shaded bar in Fig.
+ 3a. and are Clearly in good agreement with our predictions for the average NLALI.," 3a, and are clearly in good agreement with our predictions for the average MAH."
+ The maximum rate of gas accretion at Ly is attained at &10.5 Civrs. and strongly decreases. after this epoch.," The maximum rate of gas accretion at $R_0$ is attained at $\approx
+10.5$ Gyrs, and strongly decreases after this epoch."
+ Phe SELL attains a broad maximum at S-6 Civis (ς=0.9 0.6) ane then smoothly declines by a [actor of z2 towards the present day., The SFH attains a broad maximum at 8-6 Gyrs $z=0.9 - 0.6$ ) and then smoothly declines by a factor of $\approx 2$ towards the present day.
+ Ehe level of 3.1 pe7 | we obtain for this quantity lies well within the range determined: observationallv. shown by the vertical shaded box at t—0 Cyr (Talbot. 1980. Rana 1987).," The level of 3.1 $^{-2}$ $^{-1}$ we obtain for this quantity lies well within the range determined observationally, shown by the vertical shaded box at t=0 Gyr (Talbot 1980, Rana 1987)."
+ A similar SEI is obtained for the MW. model with a shallow core (thick dashed line). which illustrates the robustness of our results regarding SELls. to the existing discrepancies in dark halo density profiles between theory and observations.," A similar SFH is obtained for the MW model with a shallow core (thick dashed line), which illustrates the robustness of our results regarding SFHs, to the existing discrepancies in dark halo density profiles between theory and observations."
+ Our SER at the solar radius is almost incdistinguishable from an average of an annulus with inner and outer racii of 7.0 ancl LO kpe. respectively.," Our SFR at the solar radius is almost indistinguishable from an average of an annulus with inner and outer radii of 7.0 and 10 kpc, respectively."
+ The local SELL inferred. [rom observations might refer to an annulus of this size because due to stellar cilfusion. stars which are today in the solar neighbourhood could have been formed. at Calactocentric radii dillering by up to 1.5 kpe (e.g. Dehnen Binney 1998h).," The local SFH inferred from observations might refer to an annulus of this size because due to stellar diffusion, stars which are today in the solar neighbourhood could have been formed at Galactocentric radii differing by up to $\sim 1.5$ kpc (e.g. Dehnen Binney 1998b)."
+The cross section for this reaction is known with ereat accuracy (Wishart 1979).,The cross section for this reaction is known with great accuracy (Wishart 1979).
+ We compare in Fie., We compare in Fig.
+ d our fif (solid line) to the the rate tabulated by de Jones (1972) (triangles)., 1 our fit (solid line) to the the rate tabulated by de Jong (1972) (triangles).
+ The expression eiven by Teemark et al. (, The expression given by Tegmark et al. (
+1997) is shown by a dashed noe.,1997) is shown by a dashed line.
+ There is substantial aerecment between the value of this rate coefficient (to within  10%)) at temperatures below ~10° EK aud the experimental measurement at T=300 WS performed by Schineltekopf et al. (, There is substantial agreement between the value of this rate coefficient (to within $\sim 10$ ) at temperatures below $\sim$ $^3$ K and the experimental measurement at $T=300$ K performed by Schmeltekopf et al. (
+1967).,1967).
+ At hieher temperatures earlier results by Browne Dalearuo (1969) indicated. a slow rise in the value of the rate. whereas the more recent calcilatious of Launav et al. (," At higher temperatures earlier results by Browne Dalgarno (1969) indicated a slow rise in the value of the rate, whereas the more recent calculations of Launay et al. ("
+1991) show a significant decline,1991) show a significant decline.
+" Án uncertaintv of about 50ον, however. is still presen in the data due to the uncertainty on the interaction potential."," An uncertainty of about 50, however, is still present in the data due to the uncertainty on the interaction potential."
+ Our fit (solid line in Fie., Our fit (solid line in Fig.
+ 1) is based on the results by Launax ct al. (, 1) is based on the results by Launay et al. (
+1991). and is in reasonable agrecien with the fit of Shapiro laus (1987) to the data o| Bieniclk (1980).,"1991), and is in reasonable agreement with the fit of Shapiro Kang (1987) to the data of Bieniek (1980)."
+ Ou the other haud. Abel et al. (," On the other hand, Abel et al. ("
+1997) favor a rate increasing with temperature for Z7>1000 Ik. assuming a rate based ou the cross section of Browne Dalearne (1969).,"1997) favor a rate increasing with temperature for $T>1000$ K, assuming a rate based on the cross section of Browne Dalgarno (1969)."
+the data obtained bv Alcaino et al.,"the data obtained by Alcaino et al.,"
+ no systematic deviations between them are secu., no systematic deviations between them are seen.
+ The GB piruneters are determined from Ves.R) diagrauus in a wav similar to that used for Ves.(2.V) diagrns., The GB parameters are determined from $V vs. (B-R)$ diagrams in a way similar to that used for $V vs. (B-V)$ diagrams.
+ The relatious between these parameters and metallicity have been derived using Table 1 aud the Epà values for some clusters. specified in the previous paragraph.," The relations between these parameters and metallicity have been derived using Table 1 and the $E_{B-V}$ values for some clusters, specified in the previous paragraph."
+ Color excesses Ey;47 are transformed to Lpη through the Crebol aud Roberts (1995) ratio: Epp/Epv = 1.62., Color excesses $E_{B-V}$ are transformed to $E_{B-R}$ through the Grebel and Roberts (1995) ratio: $E_{B-R}/E_{B-V} $ = 1.62.
+ The equatious that connect the GB parameters aud metallicity of the clusters (fig 3 and fig 1) can be used todetermine the metallicity [FeΠΠ aud color excesses Ep g, The equations that connect the GB parameters and metallicity of the clusters (fig 3 and fig 4) can be used todetermine the metallicity $[Fe/H]$ and color excesses $E_{B-R}$ :
+ The equatious that connect the GB parameters aud metallicity of the clusters (fig 3 and fig 1) can be used todetermine the metallicity [FeΠΠ aud color excesses Ep ge, The equations that connect the GB parameters and metallicity of the clusters (fig 3 and fig 4) can be used todetermine the metallicity $[Fe/H]$ and color excesses $E_{B-R}$ :
+We modeled the complete data set (as obtained from the literature) using a two-WKeplerian LET model.,We modeled the complete data set (as obtained from the literature) using a two-Keplerian LTT model.
+ In their analysis and by construction. Qianetal.(2011). forced the inner companion to be in a circular orbit.," In their analysis and by construction, \citet{Qian2011} forced the inner companion to be in a circular orbit."
+ While this requirement reduces the number of free parameters. it does not account. for the possibility of eccentricity excitation of the inner planet by an outside perturber (Malmberg&Davies2009).," While this requirement reduces the number of free parameters, it does not account for the possibility of eccentricity excitation of the inner planet by an outside perturber \citep{MalmbergDavies2009}."
+. In our study. we allowed the inner planet to also attain some eccentricity.," In our study, we allowed the inner planet to also attain some eccentricity."
+ Lt is however. important to point out that in the two-Ixepler LPL model. no mutual gravitational perturbations are included.," It is however, important to point out that in the two-Kepler LTT model, no mutual gravitational perturbations are included."
+ In our first approach to find a best-fitting mocel. we generated. 111.544 initial guesses.," In our first approach to find a best-fitting model, we generated 111,844 initial guesses."
+ Each LAL guess was allowed a maximum of 500 iterations prior to termination., Each LM guess was allowed a maximum of 500 iterations prior to termination.
+ We carried out the Monte Carlo simulations on a supercomputer lasting for about 14 davs of non-stop computing., We carried out the Monte Carlo simulations on a supercomputer lasting for about 14 days of non-stop computing.
+ In detail. we used the two-sided numerical derivative option available in algorithm.," In detail, we used the two-sided numerical derivative option available in algorithm."
+ We also chose to assign individual maxiniun step-sizes to each mocdel parameter., We also chose to assign individual maximum step-sizes to each model parameter.
+ For example. we allowed the eccentricity to only change by at most per iteration.," For example, we allowed the eccentricity to only change by at most per iteration."
+ Maximum change in an angular quantity was 57., Maximum change in an angular quantity was $5^{\circ}$.
+ Similar reasonable values were assigned to other parameters., Similar reasonable values were assigned to other parameters.
+ We believe that our initial choices of the value of the parameters renders the search-space of 47 to have a fine enough grid for each iteration of the fitting algorithm., We believe that our initial choices of the value of the parameters renders the search-space of $\chi_{r}^2$ to have a fine enough grid for each iteration of the fitting algorithm.
+ We also believe that with such a Large number of initial guesses the full X7 space will be sulliciently explored., We also believe that with such a large number of initial guesses the full $\chi_{r}^2$ space will be sufficiently explored.
+ During the fitting process. we held all twelve parameters of the svstem free. and decreased the numerical accuracy parameters in from their default values.," During the fitting process, we held all twelve parameters of the system free, and decreased the numerical accuracy parameters in from their default values."
+ We also decreased the parameter step-size with which the numerical derivatives were evaluated., We also decreased the parameter step-size with which the numerical derivatives were evaluated.
+ We then carried out a Monte Carlo experiment based on a total of 111844 randomly distributed initial guesses., We then carried out a Monte Carlo experiment based on a total of 111844 randomly distributed initial guesses.
+ For each converged initial guess with «7«10.0. we recorded the resulting 47. the number of iterations. the initial guess parameters. and the final converged model parameters as returned by the LAL minimisation algorithm.," For each converged initial guess with $\chi_{r}^2 \le 10.0$, we recorded the resulting $\chi_{r}^2$, the number of iterations, the initial guess parameters, and the final converged model parameters as returned by the LM minimisation algorithm."
+ In order to avoid any bias elfect. we removed converged iterations that reached a lower or upper boundary value.," In order to avoid any bias effect, we removed converged iterations that reached a lower or upper boundary value."
+ That means. once à parameter reached any of the boundaries within its defined range (see above). its value was held constant and the parameter was no longer changed by the minimisation algorithm for carrying out the remaining of the iterations.," That means, once a parameter reached any of the boundaries within its defined range (see above), its value was held constant and the parameter was no longer changed by the minimisation algorithm for carrying out the remaining of the iterations."
+ Although we have no strict explanation for this “stickiness” effect. we judge that it has the potential to skew the final distribution of model parameters.," Although we have no strict explanation for this “stickiness” effect, we judge that it has the potential to skew the final distribution of model parameters."
+ The idea is that [rom a mathematical point of view. the minimisation algorithm attempts to cross the lower or upper boundary to obtain a better fit based on the steepest descent in. parameter space.," The idea is that from a mathematical point of view, the minimisation algorithm attempts to cross the lower or upper boundary to obtain a better fit based on the steepest descent in parameter space."
+ Hence it remains constant at these locations but with no physical meaning., Hence it remains constant at these locations but with no physical meaning.
+ As an example. one can consider the orbital cecentricity.," As an example, one can consider the orbital eccentricity."
+ Since in the systems of our interest (cireumbinary. planetary systems) eccentricities smaller than zero do not correspond to acceptable orbits. iterations that were stuck at the lower boundary. of ο=0 were removed.," Since in the systems of our interest (circumbinary planetary systems) eccentricities smaller than zero do not correspond to acceptable orbits, iterations that were stuck at the lower boundary of $e=0$ were removed."
+ This allows us to avoid introducing a bias towards a preferred direction in parameter space that has no reasonable physical significance., This allows us to avoid introducing a bias towards a preferred direction in parameter space that has no reasonable physical significance.
+ In addition.MPFIT algorithm does not return any formal uncertainties for fitted parameters that reach a lower or upper boundary.," In addition, algorithm does not return any formal uncertainties for fitted parameters that reach a lower or upper boundary."
+ After this cleaning process. we ended up with a data file containing a total of 48485 initial guesses.," After this cleaning process, we ended up with a data file containing a total of 48485 initial guesses."
+ Our best Gt resulted. in a yy=143 which had. an occurrence. frequency of fourteen times in total., Our best fit resulted in a $\chi_{r}^2 = 1.43$ which had an occurrence frequency of fourteen times in total.
+ We show the corresponding ο6 diagram of this case in Fig. 2.., We show the corresponding $O-C$ diagram of this case in Fig. \ref{BestFitOmC}.
+ Fitted. model and. derived: parameters are shown in Table 1. with formal uncertainties as obtained from the solutions covariance matrix., Fitted model and derived parameters are shown in Table \ref{fitparamtable} with formal uncertainties as obtained from the solutions covariance matrix.
+ We point out the close agreement between the two fitted. periods (6.5 ancl 15.5 vears. respectively) and the periods obtained from the Lomb-Scargle frequency analysis shown in Fig. 1..," We point out the close agreement between the two fitted periods (6.5 and 15.5 years, respectively) and the periods obtained from the Lomb-Scargle frequency analysis shown in Fig. \ref{periodogram}."
+ We visually inspected the remaining thirteen mocdels and did not find any differences., We visually inspected the remaining thirteen models and did not find any differences.
+ All final fitted: parameters agreed with one another., All final fitted parameters agreed with one another.
+ However. we noticed a qualitative dillerence when comparing our best fit with Fig.," However, we noticed a qualitative difference when comparing our best fit with Fig."
+ 2 in (2011)., 2 in \citet{Qian2011}.
+. These authors show a better description of the newly added timing data but they do not provide a quantitative X7 value for their best fit., These authors show a better description of the newly added timing data but they do not provide a quantitative $\chi_{r}^2$ value for their best fit.
+ When comparing our results. we noticed that the inner companion in our model can attain eccentric orbits. which might explain the dilference between the fits presented in this work and in Qianetal.(2011).," When comparing our results, we noticed that the inner companion in our model can attain eccentric orbits, which might explain the difference between the fits presented in this work and in \cite{Qian2011}."
+.. In. addition. in Fig. 2..," In addition, in Fig. \ref{BestFitOmC},"
+ it appears that a single ΤΠ (sinusoidal) curve will provide a poor description of the timing data., it appears that a single LTT (sinusoidal) curve will provide a poor description of the timing data.
+ Neither the individual short (inner planet) nor the long-period (outer planet) LET model will be able to provide a satisfactory fit by itself., Neither the individual short (inner planet) nor the long-period (outer planet) LTT model will be able to provide a satisfactory fit by itself.
+ Only by combining the two L'T'T models one can obtain a satisfactory fit., Only by combining the two LTT models one can obtain a satisfactory fit.
+ Our final orbital parameters for the two planetary companions are relatively close to ones reported by Qianetal.(2011) except for a slight increase (decrease) in the inner (outer) companions orbital eccentricity., Our final orbital parameters for the two planetary companions are relatively close to ones reported by \cite{Qian2011} except for a slight increase (decrease) in the inner (outer) companion's orbital eccentricity.
+ The outer companion's mass and semi-major axis are also larger in our work., The outer companion's mass and semi-major axis are also larger in our work.
+ We examined the orbital stability of our best-fitting model (Table 1)) using he orbital integration package (Chambers&Migliorini1997:Chambers|1999).," We examined the orbital stability of our best-fitting model (Table \ref{fitparamtable}) ) using the orbital integration package \citep{MERCURY-1, MERCURY-2}."
+.. We integrated. the three-body equations of motion using the variable-step Bulirsh-Stoer (BS2) algorithm., We integrated the three-body equations of motion using the variable-step Bulirsh-Stoer (BS2) algorithm.
+ Initial step-size was chosen to be of a cay and he accuracy parameter was set to 10.17ο, Initial step-size was chosen to be of a day and the accuracy parameter was set to $10^{-16}$.
+ To be consistent with the fitted. LEP orbits. we considered the binary pair as one single object with à mass equal to the sum of the masses of the two stars.," To be consistent with the fitted LTT orbits, we considered the binary pair as one single object with a mass equal to the sum of the masses of the two stars."
+ We carried out integrations for two scenarios: i) the nominal parameters for the best-fitting model with (S=1.43). and ii) an optimistic scenario where we considered the lower init in masses and eccentricitv and upper limit in the semi-major axis for both companions.," We carried out integrations for two scenarios: i) the nominal parameters for the best-fitting model with $(\chi_{r}^2 = 1.43)$, and ii) an optimistic scenario where we considered the lower limit in masses and eccentricity and upper limit in the semi-major axis for both companions."
+ We integrated both systems for 107 vers., We integrated both systems for $10^4$ years.
+ The results are shown in Fig. 3.., The results are shown in Fig. \ref{TwoOrbitFigs}.
+ This short integration time is sullicient to demonstrate an unstable configuration or the first scenario., This short integration time is sufficient to demonstrate an unstable configuration for the first scenario.
+ The upper panel of the figure shows the orbits based on the nominal best-fitting parameters and clearly shows orbital instability of at least one of the planet companion that resulted in its eventual escape., The upper panel of the figure shows the orbits based on the nominal best-fitting parameters and clearly shows orbital instability of at least one of the planet companion that resulted in its eventual escape.
+ The optimistic scenario shown in the lower panel of Pig., The optimistic scenario shown in the lower panel of Fig.
+ 3 resulted in a relatively stable orbital architecture when compared with the nominal scenario., 3 resulted in a relatively stable orbital architecture when compared with the nominal scenario.
+ The corresponding Ixepler elements of the orbit in this case and their time evolutions are shown in, The corresponding Kepler elements of the orbit in this case and their time evolutions are shown in
+(2001) and for groups. of galaxies in the Zwicky catalogue by Pacilla (2001).,(2001) and for groups of galaxies in the Zwicky catalogue by Padilla (2001).
+ After giving a brief flavour of the clustering of 2dECGIS galaxies in redshift space in section 4.. we focus our attention on clustering in real space in the remainder of the paper.," After giving a brief flavour of the clustering of 2dFGRS galaxies in redshift space in Section \ref{sec:redspace}, we focus our attention on clustering in real space in the remainder of the paper."
+ Phe clustering signal in real space is inferred by integrating £(o.7) in the x direction (i.e. along the line of sight): For the samples that we consider. the integral converges bv pair separations of wm=50h!Mpc.," The clustering signal in real space is inferred by integrating $\xi(\sigma,\pi)$ in the $\pi$ direction (i.e. along the line of sight): For the samples that we consider, the integral converges by pair separations of $\pi\,\ge\,50\,$."
+ The. projected correlation function can then be written as an integral over the spherically averaged real space correlation function. £C). (Davis Peebles 1983).," The projected correlation function can then be written as an integral over the spherically averaged real space correlation function, $\xi(r)$, (Davis Peebles 1983)."
+ Lowe assume that the real space correlation function is a power law (which is a fair approximation for APAL galaxies out to separations around ro10h Alpe. see eg. Baugh 1996). then Eq.," If we assume that the real space correlation function is a power law (which is a fair approximation for APM galaxies out to separations around $r\sim10\,h^{-1}\,$ Mpc, see e.g. Baugh 1996), then Eq."
+ 7. can be wrltten as ∖∖⋎↓↥∢⊾↓⋅⋖⋅↓∖↿∖⇁≀⋅∃⊲↓⊳∖↿↓↕⋖⋅⊔⊳∖⊔⋜↧⊔∶⋜⋯↓⊔⋯⇂⋅⇂⇂↓↥≼⇍⇂↕∢≱↓↥⊳⋜⋯∠⇂∖∖⊽⋖⋅↓⋯∖⇁∢⊾ ⊔⊳∖⋖⋅∠⇂≾↿∖∣⋮⊐∶↿∖∣⋮⊔∣⋮⊐⋅⊳∖∖⋎↓↕⋖⋅↓⋅∢⊾∣⋮⊔⊀↓⊳∖↿↓↥∢⊾↓⋅⋖⋅⋜↧↓⊳∖↓≻⋯⇍⋖⋅≼⇍∪↓⋅↓⋅⋖⋅⋜∐⊲↓∪⊔ length and 5 is equal to the slope of the projected correlation function Ἔσ)έσ.," \ref{eq:projxi} can be written as where $\Gamma(x)$ is the usual Gamma function, and we have used $\xi(r) = (r_{0}/r)^{\gamma}$ , where $r_{0}$ is the real space correlation length and $\gamma$ is equal to the slope of the projected correlation function $\Xi(\sigma)/\sigma$."
+ As we demonstrate in Section. 4.2.. the projected. correlation function is well described. by a power law.," As we demonstrate in Section \ref{sec:xi.s}, the projected correlation function is well described by a power law."
+ We study a range of samples containing cillerent numbers of galaxies and covering dillercnt volumes of the Universe., We study a range of samples containing different numbers of galaxies and covering different volumes of the Universe.
+ Lt is imperative to include sampling Iuctuations when estimating the errors on the measured. correlation function. to allow a meaningful comparison of the results obtained. from. different saniples.," It is imperative to include sampling fluctuations when estimating the errors on the measured correlation function, to allow a meaningful comparison of the results obtained from different samples."
+ “Vhis contribution to the errors has often been neglected in previous work., This contribution to the errors has often been neglected in previous work.
+ Following (2001)... we employ a sample of 22 mock 2dGIU catalogues drawn from the κΕΔΕ Hubble Volume simulation (Evrard 2002) to estimate the error bars on the measured correlation functions.," Following , we employ a sample of 22 mock 2dFGRS catalogues drawn from the $\Lambda$ CDM Hubble Volume simulation (Evrard 2002) to estimate the error bars on the measured correlation functions."
+ The construction of these mock catalogues is explained in Baugh (2002. in preparation: see also Cole 1905 and Norberg 2002).," The construction of these mock catalogues is explained in Baugh (2002, in preparation; see also Cole 1998 and Norberg 2002)."
+ These catalogues have the same selection criteria and the same clustering amplitude as measured for galaxies in the Dux-limited. 2dECGIUS., These catalogues have the same selection criteria and the same clustering amplitude as measured for galaxies in the flux-limited 2dFGRS.
+ We have experimentecl with ensenibles of mock catalogues. constructed. with dilferent clustering strengths to ascertain how best to assign error bars when the measuredclustering hasa different amplitude, We have experimented with ensembles of mock catalogues constructed with different clustering strengths to ascertain how best to assign error bars when the measuredclustering hasa different amplitude
+preseuce of a local heat source that is likely to be star formation (Woltectal.2008).. and in CRB-DLA systems there is obviously evidence for ongoing star formation.,"presence of a local heat source that is likely to be star formation \citep{wolfe08a}, and in GRB-DLA systems there is obviously evidence for ongoing star formation."
+ Qur result simply siieeestsOO that the star formation must be takine place iu other parts of these galaxies. which have significantly higher column deusities aud molecular contents than the sightlines we most commonly observe as DLAs.," Our result simply suggests that the star formation must be taking place in other parts of these galaxies, which have significantly higher column densities and molecular contents than the sightlines we most commonly observe as DLAs."
+" We thank €. MAIcKvce aud the referee. SSchawo. for helpful conunents,"," We thank C. McKee and the referee, Schaye, for helpful comments."
+ Support for this work was provided bv the Alfred SSloan Foundation (ΑΠ). by NÀSA/JPL through the Spitzer Theoretical Research Program (MBIS). bv the >National Science. Foundation through erauts AST-08077309 (NIRE) and AST-0709235 (IND). aud by au NSERC Discovery GCraut (SLE).," Support for this work was provided by the Alfred Sloan Foundation (MRK), by NASA/JPL through the Spitzer Theoretical Research Program (MRK), by the National Science Foundation through grants AST-0807739 (MRK) and AST-0709235 (JXP), and by an NSERC Discovery Grant (SLE)."
+"To see why 5,5 dillers [rom μμ. let us write them in more intuitive forms: (z2)-— —: due =0|[z.z). where we have used D4(à=1ο)τςλατ) (Nang2000b).. and 2’)[A(z)—A(2)] is the angular diameter distance of a source at 2 from a fictitious observer al z/ lor A=0 (Schneideretal.1992).","To see why $\kappa_{min}$ differs from $\tilde{\kappa}_{min}$, let us write them in more intuitive forms: = - _0^z =0|z',z), where we have used $D_A(\tilde{\alpha}=1|z)=r(z)/(1+z)$ \citep{SNfluxavg}, and $D_A(\tilde{\alpha}=0|z',z)= (1+z')\, [\lambda(z)-\lambda(z')]$ is the angular diameter distance of a source at $z$ from a fictitious observer at $z'$ for $\tilde{\alpha}=0$ \citep{Sch92}."
+. Similarly. we can write where r(\(2)—ας)/(14+2) is the angular diameter distance between z and z'/ in a completely smooth universe (i.e.. filled beam. à= 1).," Similarly, we can write = - _0^z where $r\left( \chi(z)-\chi(z')\right)/(1+z)$ is the angular diameter distance between $z$ and $z'$ in a completely smooth universe (i.e., filled beam, $\tilde{\alpha}=1$ )."
+" Since the (rue minim convergence should correspond to the minimum magnification. which occurs when (he beam connecting the source and observer isempfg (ie. à— 0). &jyju(z) is the true minimum convergence. while &,,;,Cz) is Ils approximation when the fillecl beam (instead of the empty beam) angular ciameter distance is used along the line of sight from the observer to the source."," Since the true minimum convergence should correspond to the minimum magnification, which occurs when the beam connecting the source and observer is (i.e., $\tilde{\alpha}=0$ ), $\tilde{\kappa}_{min}(z)$ is the true minimum convergence, while $\kappa_{min}(z)$ is its approximation when the filled beam (instead of the empty beam) angular diameter distance is used along the line of sight from the observer to the source."
+" Kimin(2) has been mistakenly associated with j5,5, in the literature.", $\kappa_{min}(z)$ has been mistakenly associated with $\mu_{min}$ in the literature.
+" Note that |8,,5,| unin]. hence L/(1—yin)? underestimates the minimum magnification 4655."," Note that $|\kappa_{min}| < |\tilde{\kappa}_{min}|$ , hence $1/(1-\kappa_{min})^2$ underestimates the minimum magnification $\mu_{min}$."
+" Some of the consequences of the subtle difference between £,,5, and &,,;, wil be presented in Sec", Some of the consequences of the subtle difference between $\tilde{\kappa}_{min}$ and $\kappa_{min}$ will be presented in Sec.
+LV.,IV.
+" Fie.2 shows that the correction to &,,;5,0:) required for the accurate calculation of the minimun convergence £,,5,0:) is quite modest (less (han ab 2< 1) in cosmological models (hat fit current observational data (the ACDM model).", Fig.2 shows that the correction to $\kappa_{min}(z)$ required for the accurate calculation of the minimum convergence $\tilde{\kappa}_{min}(z)$ is quite modest (less than at $z\leq 1$ ) in cosmological models that fit current observational data (the $\Lambda$ CDM model).
+ Let us define a modified reduced convergence given by lt which should be compared with the reduced convergence defined by Valageas(2000) and MunshianclJain (2000)..," Let us define a modified reduced convergence given by 1 + which should be compared with the reduced convergence defined by \cite{Valageas00} and \cite{MJ00}, ,"
+davs.,days.
+ This value is confine by our reanalysis of archival spectroscopic data., This value is confirmed by our reanalysis of archival spectroscopic data.
+" Our reanalysis of the spectroscopy also vields a velocity scuuamplitude for the companion star of ου=d4ua304gms ο, which when combined with he orbital period. gives an optical mass functiou of FAL)=2.73£0.56AL..."," Our reanalysis of the spectroscopy also yields a velocity semiamplitude for the companion star of $K_2=435\pm 30$ km $^{-1}$, which when combined with the orbital period, gives an optical mass function of $f(M)=2.73\pm 0.56\,M_{\odot}$."
+ The spectral type of he companion star is not well constrained because we only have six template spectra available to us., The spectral type of the companion star is not well constrained because we only have six template spectra available to us.
+" A template spectrmu with a spectral type of IKIV xovides the best match: next best matches are oxovided by template spectra with spectral types of COV and W200. respectively,"," A template spectrum with a spectral type of K4V provides the best match; next best matches are provided by template spectra with spectral types of G5V and K2III, respectively."
+ The πο spectrum suggests that the accretion disk may contribute a substantial amount of light iu £A (~80% or more)., The summed spectrum suggests that the accretion disk may contribute a substantial amount of light in $R$ $\approx 80\%$ or more).
+" Using the R-haud light curve. we find a lower limit to the inclination of /=5043"". which gives an upper limit to the mass of the black hole of Af,<7.3AL... assunine a niass ratio of GQ=10."," Using the $R$ -band light curve, we find a lower limit to the inclination of $i=50\pm 3^{\circ}$, which gives an upper limit to the mass of the black hole of $M_1\le 7.3\,M_{\odot}$, assuming a mass ratio of $Q=10$."
+ If we assune that the accretion disk contributes of the light iu P. then the iuclinatiou would be about 707 or more. aud the black hole mass would be about LAL. (again assunine a lass ratio of Q= 10).," If we assume that the accretion disk contributes of the light in $R$, then the inclination would be about $70^{\circ}$ or more, and the black hole mass would be about $4\,M_{\odot}$ (again assuming a mass ratio of $Q=10$ )."
+ We thank Matt. IIolinau. and Paul Schechter for providing supplementary photometric data aud Mauricio Navarette for help with the observations., We thank Matt Holman and Paul Schechter for providing supplementary photometric data and Mauricio Navarette for help with the observations.
+ This work was supported in part by NASA eraut NAG5-993, This work was supported in part by NASA grant NAG5-9930.
+llowever. all the work done so far relates only to the classical regime.,"However, all the work done so far relates only to the classical regime."
+ Quantum ellects have remained out of reach., Quantum effects have remained out of reach.
+ The most famous quantum effect is. of course. Hawking (1974) radiation.," The most famous quantum effect is, of course, Hawking (1974) radiation."
+ Unfortunately. or massive DlIls such as the ones discovered so fur. the ILawking temperature is so small e10TOM./M) IX} that there is no hope of ever detecting this radiation.," Unfortunately, for massive BHs such as the ones discovered so far, the Hawking temperature is so small $\sim 10^{-7}(M_\odot/M)$ K] that there is no hope of ever detecting this radiation."
+ The story is different in the case of DIIs with AZ<10’ e. Such DIIs will evaporate by the Hawking nechanism in a time less (han the age of the universe and will explode with final Iuminosities ~10?eves! (Shapiro Teukolskv 1983).," The story is different in the case of BHs with $M<10^{15}$ g. Such BHs will evaporate by the Hawking mechanism in a time less than the age of the universe and will explode with final luminosities $\sim 10^{20} ~{\rm
+erg\,s^{-1}}$ (Shapiro Teukolsky 1983)."
+ If the universe made such mini-DIIs in sufficient iunmbers. we might expect to see a few of them cie in fleeting bursts of hieh energy emission.," If the universe made such mini-BHs in sufficient numbers, we might expect to see a few of them die in fleeting bursts of high energy emission."
+ Given our present ποΣοπέΠας of the formation of compact objects by. stellar collapse. and given the constraints on density [Inetuations in the early universe as determined from observations of the cosmic microwave background. there does not seem to be much chance of finding 10Pe DIIs in the universe today.," Given our present understanding of the formation of compact objects by stellar collapse, and given the constraints on density fluctuations in the early universe as determined from observations of the cosmic microwave background, there does not seem to be much chance of finding $10^{15}$ g BHs in the universe today."
+ However. nature has a long history of surprising us nowhere more (rue (han in astrophvsics so perhaps we should not eive up hope altogether.," However, nature has a long history of surprising us — nowhere more true than in astrophysics — so perhaps we should not give up hope altogether."
+ This work was supported in part by NASA erant. NAG5-LO780 and NSF grant. AST 0307433., This work was supported in part by NASA grant NAG5-10780 and NSF grant AST 0307433.
+standard error at each Ireeuencey. point.,standard error at each frequency point.
+ That is. to find the best-lit power-law slope lor a eiven power spectrum. we minimize (he quantity with respect toC and o.," That is, to find the best-fit power-law slope for a given power spectrum, we minimize the quantity with respect to$C$ and $\alpha$."
+ This form tends to weight most heavily the hiehest decade of frequencies because the frequency points are separated by a constant Av=(8000).3., This form tends to weight most heavily the highest decade of frequencies because the frequency points are separated by a constant $\Delta \nu = (8000M)^{-1}$.
+ A Κον physical question we wish to explore is: how well correlated is the variability at different radii?, A key physical question we wish to explore is: how well correlated is the variability at different radii?
+ Answering this question will help us understand the light eurve's power spectrum and provide a context for comparison to other models (e.g..2))., Answering this question will help us understand the light curve's power spectrum and provide a context for comparison to other models \cite{1997MNRAS.292..679L}) ).
+ Correlations are performed by first decomposing the light curves PDS into partitions., Correlations are performed by first decomposing the light curve's PDS into partitions.
+" Let there be Vv partitions. each with a different light curve £5(/) so that our total light curve is their simple sum: F(/)=SOF,\)."," Let there be $N$ partitions, each with a different light curve $F_n\left(t\right)$ so that our total light curve is their simple sum: $F\left(t\right) = \sum_{n=1}^N F_n\left(t\right)$."
+" The total power spectrum P(77) can therefore be expressed as a sum of the partitions PDSs (5,) plus a sum that depends on how well the different modes are correlated. (55): Tere At),=vy—Uy is the dilferencein phase at frequency v between two partitions.", The total power spectrum $P(\nu)$ can therefore be expressed as a sum of the partitions' PDSs $S_a$ ) plus a sum that depends on how well the different modes are correlated $S_b$ ): Here $\Delta \psi_{n m} = \psi_n - \psi_m$ is the differencein phase at frequency $\nu$ between two partitions.
+ Note that even though5 P(») is independent of our partition scheme. the relative sizes of ↙ ∣⇆⊽∣≀↧↴∐≺⇂∣," Note that even though $P(\nu)$ is independent of our partition scheme, the relative sizes of $S_a$ and $S_b$ are not."
+"⇆⊽∣↗≀↕↴↕⋅≼↲∐∪↥⋅∐≯⊔∐↲⇀∣↥∣≀⋯↲≀↧↴∐∪↓≯⋟∖⊽↕∐↓∐≀∐⋅∐↓≀↧↴∩∐∐⋯⇂≼↲⋅⊔∐↲∐∣⇆⊽∣∕∕∣⇆⊽��↾⇢∩≀↧⊔∖⊽↼∖⊽⇢↕⋊↽↴aft (5, N(N —1)while$,  IN)and ρου 4 QasN — 1(thereis only one partition and one signal is perfectly coherent with itself)."," If the $A_n$ are all of similar magnitude, then $S_a / S_b \rightarrow 0$ as $N \rightarrow \infty$ $S_b \sim N\left(N-1\right)$ while $S_a \sim N$ ) and $S_b / S_a \rightarrow 0$ as $N \rightarrow 1$ (there is only one partition and one signal is perfectly coherent with itself)."
+ If the partitions are perfectly incoherent. then Dc ο.," If the partitions are perfectly incoherent, then $P \simeq S_a$ ."
+" Conversely. if thev are perfectly. coherent. (hen 2~5$, for No> 2."," Conversely, if they are perfectly coherent, then $P \simeq S_b$ for $N>2$ ."
+ sample light curves and their corresponding power spectra can be seen in Figures 3 , Sample light curves and their corresponding power spectra can be seen in Figures \ref{fig:light-curves-incl} 
+1998).,.
+. The secouc order coutribution we are going to discuss in thisLeller is very similar to this effect. only that the source of the fluctuation in the free electron deusity is clilfereut.," The second order contribution we are going to discuss in this is very similar to this effect, only that the source of the fluctuation in the free electron density is different."
+ The source of reionization is thought to be the UV radiation of early objects like qttasars aic proto-galaxies as hosts of au early generation of stars., The source of reionization is thought to be the UV radiation of early objects like quasars and proto-galaxies as hosts of an early generation of stars.
+ The uuclear aud gravitational euergy of these objects is transformed into radiatiou which subsequently ionizes the hydrogen in spheres which surround them (Teemark.Silk&Blanchard1991:Rees1996:Aghanimetal.HaimanLoeb1997.1993:Loeb1997:Silk&Rees1998:HaimanIxnox 1999).," The nuclear and gravitational energy of these objects is transformed into radiation which subsequently ionizes the hydrogen in spheres which surround them \citep{tegmark:94,rees:96,aghanim:96,haiman:97,haiman:98,loeb:97,silk:98,haiman:99}."
+. Que way to study the consequeuces oL inhomogeneous reionization is to use models which describe the distribution of ionized regious by a stuall number of [ree parameters (Ciruzinov&Hu1998:Ixuox.ScocciniarroDodel-son1998:Haitian&Ixnox 1999).," One way to study the consequences of inhomogeneous reionization is to use models which describe the distribution of ionized regions by a small number of free parameters \citep{hu:98,knox:98,haiman:99}."
+. We adopt the model by Ciruzinov&Hu(1998).. which describes inhomogeneous reionizatiou as a set of uncorrelated patches of a certain fixed size 2.," We adopt the model by \citet{hu:98}, which describes inhomogeneous reionization as a set of uncorrelated patches of a certain fixed size $R$."
+ The number density of these patches grows with time and finally the whole uuiverse becomes relouized in a homogeneous way., The number density of these patches grows with time and finally the whole universe becomes reionized in a homogeneous way.
+ More realistic models include varyiug patch sizes (Aghanimetal.1996) and contaiu correlatious of the iouized regions., More realistic models include varying patch sizes \citep{aghanim:96} and contain correlations of the ionized regions.
+ It turus out that these correlatious lead to a somewhat different sigual ou smallest scales for the temperature anisotropies (Ixnox.Scoccimarro 1903)., It turns out that these correlations lead to a somewhat different signal on smallest scales for the temperature anisotropies \citep{knox:98}.
+. However. the naive uncorrelated model by Ciruzinov&Hu(1908). gives a good estimate of the effect on the CMB (Haiman&Inox1999).," However, the naive uncorrelated model by \citet{hu:98} gives a good estimate of the effect on the CMB \citep{haiman:99}."
+. We defiue the ionization fraction ore to be the ratio of the number density of [ree electrous ne and the overall (free aud bound) number density of electrous n. i.e. cre=no/n.," We define the ionization fraction $x_{\rm e}$ to be the ratio of the number density of free electrons $\e$ and the overall (free and bound) number density of electrons $n$, i.e. $x_{\rm e} = n_{\rm e}/n$."
+ Since we want to study the effects of inhomogeneous reionization ou the CMB to second order. we not only. need the mean background ionization fraction ore. we also lave to know the variance of the distribution.," Since we want to study the effects of inhomogeneous reionization on the CMB to second order, we not only need the mean background ionization fraction $\xeb$, we also have to know the variance of the distribution."
+ Cruzinov&Hu(1998). give the secoud moment to be 2 = le where =[fdl/a is the conformal time. with « the scale factor. tmin=minty.72) aud Haas=Waxy. gs).," \citet{hu:98} give the second moment to be > = ], where $\eta \equiv \int dt/a$ is the conformal time, with $a$ the scale factor, $\eta_{\rm min} \equiv {\rm min}(\eta_1,\eta_2)$ and $\eta_{\rm max} \equiv {\rm max}(\eta_1,\eta_2)$ ."
+ The correlation (2)) drops off exponentially if the distance between two points is larger than the size of a patch A., The correlation \ref{ecor}) ) drops off exponentially if the distance between two points is larger than the size of a patch $R$.
+ IC or yo is in one of the homogeneous regimes (nearly ueutral or complete ionization) the correlation is just the product. of the mean values., If $\eta_1$ or $\eta_2$ is in one of the homogeneous regimes (nearly neutral or complete ionization) the correlation is just the product of the mean values.
+ In the followingwe will describe the inbomogeueity as the fluctuation of the free electron number density,	 In the followingwe will describe the inhomogeneity as the fluctuation of the free electron number density
+"do not contain solar neighbourghood stars and thus include stars at heliocentric distances that are larger on average than for planes that do go through the Sun with φροιε=90°,270° (this is illustrated for Φροιε~0? and d$poe~90? in Fig. 3)).","do not contain solar neighbourghood stars and thus include stars at heliocentric distances that are larger on average than for planes that do go through the Sun with $\phi_{pole}=90^\circ,270^\circ$ (this is illustrated for $\phi_{pole}\sim0^\circ$ and $\phi_{pole}\sim90^\circ$ in Fig. \ref{fig:gc3_planes}) )."
+" Therefore the number of observable stars that contribute to poles φροιε~90?,270? will be larger than in the perpendicular direction."," Therefore the number of observable stars that contribute to poles $\phi_{pole}\sim90^\circ,270^\circ$ will be larger than in the perpendicular direction."
+ These two effects reinforce one another and give rise to the morphology of the pole count map of Fig. 2.., These two effects reinforce one another and give rise to the morphology of the pole count map of Fig. \ref{fig:gc3_gal_true}.
+ The GC3 method’s capability for detecting tidal streams clearly depends on how large is the contribution of the stream with respect to the Galactic background., The GC3 method's capability for detecting tidal streams clearly depends on how large is the contribution of the stream with respect to the Galactic background.
+" In Fig. 2,,"," In Fig. \ref{fig:gc3_gal_true}, ,"
+" the mode of pole counts is 4.5x10° stars/pole, which is a measure of the typical contribution of Galactic stars in this error-free map."," the mode of pole counts is $\sim4.5\times10^6$ stars/pole, which is a measure of the typical contribution of Galactic stars in this error-free map."
+" On the other hand, a dwarf galaxy with a stellar population similar to the Galactic Halo (i.e. an age of 13.6 Gyr and iron abundance [Fe/H]= —1.7) and total luminosity of 10"" Lo has ~1.2x10° stars brighter than My=5, according to the stellar population synthesis models of Bruzual&Charlot(2003)."," On the other hand, a dwarf galaxy with a stellar population similar to the Galactic Halo (i.e. an age of 13.6 Gyr and iron abundance $[\mathrm{Fe}/\mathrm{H}]=-1.7$ ) and total luminosity of $10^7$ $_\odot$ has $\sim1.2\times10^6$ stars brighter than $M_V=5$, according to the stellar population synthesis models of \citet{bru03}."
+". The number of stars, however, will depend on the distance distribution of the stars in the stream, but will necesarily be smaller, specially since most of these are rather faint (My2 4) main-sequence stars."," The number of stars, however, will depend on the distance distribution of the stars in the stream, but will necesarily be smaller, specially since most of these are rather faint $M_V \gtrsim 4$ ) main-sequence stars."
+" Therefore, the signature of such a dwarf galaxy in a pole count map would be several times below the typical number of Galactic background stars and thus very hard to detect."," Therefore, the signature of such a dwarf galaxy in a pole count map would be several times below the typical number of Galactic background stars and thus very hard to detect."
+" As exhibited in this example, the number of Galactic contaminants in a cell is an important limitation of the GC3 method."," As exhibited in this example, the number of Galactic contaminants in a cell is an important limitation of the GC3 method."
+" In the following section we propose a modification of the method, to help discriminate stream stars from the smooth halo using the main attribute of the original GC3 method, namely the use of angular momentum conservation through a geometrical approach."," In the following section we propose a modification of the method, to help discriminate stream stars from the smooth halo using the main attribute of the original GC3 method, namely the use of angular momentum conservation through a geometrical approach."
+" The GC3 method of Johnston et al.,"," The GC3 method of Johnston et al.,"
+" as explained above, uses solely positional information."," as explained above, uses solely positional information."
+" In order to improve the ‘signal’ of a stellar stream over the ‘noise’ of the smooth background halo population, we propose the inclusion of velocity information as well."," In order to improve the `signal' of a stellar stream over the `noise' of the smooth background halo population, we propose the inclusion of velocity information as well."
+" Since the stellar stream is roughly confined to a plane that contains the center of symmetry of the potential, as seen from this center both its position and velocity vectors will be contained in the orbital plane."," Since the stellar stream is roughly confined to a plane that contains the center of symmetry of the potential, as seen from this center both its position and velocity vectors will be contained in the orbital plane."
+" A grid of all the possible great circle cells is devised, with each cell uniquely determined by itspole,, that is the unit vector L which is perpendicular to the plane it defines."," A grid of all the possible great circle cells is devised, with each cell uniquely determined by its, that is the unit vector $\mathbf{\hat{L}}$ which is perpendicular to the plane it defines."
+" The modified GC3 (mGC3) method proposed here for counting the number of stars associated with each pole, can be expressed in terms of the following Galactocentric position and velocity criteria where f and ¥ are unit vectors parallel to the star position and velocity vectors, ó. and 6, are the tolerances which allow for the width of each great circle associated with a cell."," The modified GC3 (mGC3) method proposed here for counting the number of stars associated with each pole, can be expressed in terms of the following Galactocentric position and velocity criteria where $\mathbf{\hat{r}}$ and $\mathbf{\hat{v}}$ are unit vectors parallel to the star position and velocity vectors, $\delta_r$ and $\delta_v$ are the tolerances which allow for the width of each great circle associated with a cell."
+ The pole vector L that corresponds to the cell which best coincides with a stream is parallel to the streams's total agular momentum., The pole vector $\mathbf{\hat{L}}$ that corresponds to the cell which best coincides with a stream is parallel to the streams's total agular momentum.
+" As with the original GC3, the mGC3 criteria in (2)) will hold for tidal streams evolved in a spherical or slightly flattened potential."," As with the original GC3, the mGC3 criteria in \ref{criteria}) ) will hold for tidal streams evolved in a spherical or slightly flattened potential."
+" For the Milky Way (MW), according to several authors, the orbit of the Sgr tidal stream constrains the inner halo (< 60kpc) to be slightly flattened with a Z-axis flattening q;=0.85—0.95 for oblate models (Johnston,Law&Majewski2005;; Martinez-Delgadoetal.2004;Ibataet 2001)), and ᾳ-=1.25 for prolate models (Helmi2004;Law,Johnston&Majewski2005);; albeit more recently Law,Majewski&Johnston(2009) find that a triaxial halo with 4.=1.25 and qy=1.5 provides a better fit of the radial velocity and distance distribution of observed Sgr stream stars."," For the Milky Way (MW), according to several authors, the orbit of the Sgr tidal stream constrains the inner halo $\lesssim 60$ kpc) to be slightly flattened with a Z-axis flattening $q_z=0.85-0.95$ for oblate models \citealt*{joh05}; \citealt{md04,iba01}) ), and $q_z=1.25$ for prolate models \citep*{hel04,law05}; albeit more recently \citet*{law09}
+ find that a triaxial halo with $q_z=1.25$ and $q_y=1.5$ provides a better fit of the radial velocity and distance distribution of observed Sgr stream stars."
+" On the other hand, depending on dynamical age, orbital inclination and how close the stellar system comes to the Galactic plane, the disc’s potential can perturb the symmetry of the overall potential felt by the stream."," On the other hand, depending on dynamical age, orbital inclination and how close the stellar system comes to the Galactic plane, the disc's potential can perturb the symmetry of the overall potential felt by the stream."
+" Johnstonetal.(2008) classify stream morphologies resulting from their N-body simulations in three categories ‘great circle’, ‘cloudy’ and ‘mixed’, illustrated in their Fig."," \citet{joh08} classify stream morphologies resulting from their N-body simulations in three categories `great circle', `cloudy' and `mixed', illustrated in their Fig."
+" 1, as wellas transition types in between these categories, as shown in their Fig."," 1, as wellas transition types in between these categories, as shown in their Fig."
+ 2., 2.
+" The mGC3 criteria proposed here will hold for the ‘great circle’ streams, as well as for the"," The mGC3 criteria proposed here will hold for the `great circle' streams, as well as for the"
+dust. temperatures of 35. 40. 45 and SOWKK are about 1.5. 1.5. 2.0 ancl 2.3 respectively.,"dust temperatures of 35, 40, 45 and K are about 1.5, 1.5, 2.0 and 2.3 respectively."
+ LE the specific form. of the redshift evolution of the activity parameter (£7)+ Shown in equation (18) is assumed. then a dust. temperature of 35 or KIN is most consistent with the data. the same temperature that was required. for. consistency. by. both Blain et (0109900) and. Trentham et ((1999). and. is in agreement with the dust temperatures derived for redshift QSOs by Benford et ((1999).," If the specific form of the redshift evolution of the activity parameter $(F\sigma)^{-1}$ shown in equation (18) is assumed, then a dust temperature of 35 or K is most consistent with the data, the same temperature that was required for consistency by both Blain et (1999c) and Trentham et (1999), and is in agreement with the dust temperatures derived for high-redshift QSOs by Benford et (1999)."
+ The increase in the value of the activity. parameter (Po)+ asa function of redshift can be interpreted in terms of two extreme scenarios. or as à combination of both.," The increase in the value of the activity parameter $(F\sigma)^{-1}$ as a function of redshift can be interpreted in terms of two extreme scenarios, or as a combination of both."
+ In the first scenario. the fraction. of dark halo mergers that lead. to a luminous phase in a dusty galaxy f° ds. fixed. but that the duration. of the luminous phase m ids less ab high. redshifts.," In the first scenario, the fraction of dark halo mergers that lead to a luminous phase in a dusty galaxy $F$ is fixed, but that the duration of the luminous phase $\sigma$ is less at high redshifts."
+ This is plausible. based on the results of simulations of galaxy mergers (c.g. Alihos 1999: Dokki et 11999): on average. the tvpical mass of a merging pair of galaxies is expected to be less at high redshifts in an hierarchical scenario of galaxy evolution. and the gas content of the galaxies is expected. to be greater.," This is plausible, based on the results of simulations of galaxy mergers (e.g. Mihos 1999; Bekki et 1999); on average, the typical mass of a merging pair of galaxies is expected to be less at high redshifts in an hierarchical scenario of galaxy evolution, and the gas content of the galaxies is expected to be greater."
+ As a result. the dynamical time of a merger would be expected to decrease with increasing redshift. and the viscosity of the ISAL would be expected to increase.," As a result, the dynamical time of a merger would be expected to decrease with increasing redshift, and the viscosity of the ISM would be expected to increase."
+ Both of these factors might be expected to increase the star formation cllicicney ofa merger with increasing redshift., Both of these factors might be expected to increase the star formation efficiency of a merger with increasing redshift.
+ In the second. scenario. he duration of the luminous phase associated with a merger 7 is independent of redshift. but the fraction of mergers that induce such a phase £ is reduced as redshift increases.," In the second scenario, the duration of the luminous phase associated with a merger $\sigma$ is independent of redshift, but the fraction of mergers that induce such a phase $F$ is reduced as redshift increases."
+ Lt is »rhaps more plausible that the second of these scenarios could produce the large change in the activity parameter (Fo) toby a actor of about 100 from z= Oto 2=3 that is required to fit the data.," It is perhaps more plausible that the second of these scenarios could produce the large change in the activity parameter $(F\sigma)^{-1}$ , by a factor of about 100 from $z=0$ to $z=3$ that is required to fit the data."
+ This is because the duration of the uminous phase ofa merger-induced starburst e must exceed he lifetime of a reasonably massive star. Le. σος10τ vvr.," This is because the duration of the luminous phase of a merger-induced starburst $\sigma$ must exceed the lifetime of a reasonably massive star, i.e. $\sigma \gs 10^7$ yr."
+ Lf star formation activity powers a significant fraction of the SCUBA galaxies. as sccms reasonable. then this limit to the value of the merger duration e is constrained to be greater than about 10 ver. only a few times less than the values ofσ ab 2=O listed in L1.," If star formation activity powers a significant fraction of the SCUBA galaxies, as seems reasonable, then this limit to the value of the merger duration $\sigma$ is constrained to be greater than about $^7$ yr, only a few times less than the values of $\sigma$ at $z=0$ listed in 1."
+ Phus it seems likely that a large raction of the increase in the value of the activity parameter (Fo)5. which is required at high redshifts to account. [or he observed counts. should be attributed to a reduction in he fraction £ of dark halo mergers that generate a Luminous galaxy.," Thus it seems likely that a large fraction of the increase in the value of the activity parameter $(F\sigma)^{-1}$, which is required at high redshifts to account for the observed counts, should be attributed to a reduction in the fraction $F$ of dark halo mergers that generate a luminous galaxy."
+ We speculate that this may be connected. with the ower tvpical metallicity expected at higher redshifts., We speculate that this may be connected with the lower typical metallicity expected at higher redshifts.
+ In a ower metallicity system the cooling of dense gas would be expected to be less efficient. ancl so a large amount of high- star formation may be unable to take place during the short merger process.," In a lower metallicity system the cooling of dense gas would be expected to be less efficient, and so a large amount of high-mass star formation may be unable to take place during the short merger process."
+ Counts predicted by the four models. listed. in 11. emploving the values of ps listed above. are compared. with observations at wavelengths of 15. 60. 175. 450. 850. 1300 ancl jam in 99.," Counts predicted by the four models listed in 1, employing the values of $p_\sigma$ listed above, are compared with observations at wavelengths of 15, 60, 175, 450, 850, 1300 and $\mu$ m in 9."
+ While these models do not. present unique solutions. fewer parameters are involved in the mocel than the number of separate pieces of constraining data.," While these models do not present unique solutions, fewer parameters are involved in the model than the number of separate pieces of constraining data."
+ In future. by comparing the predictions of the mocdels with observations. especially with the redshift distributions of the SCUBA ealaxics (Barger et 11990: Lilly οἱ 11999: Smail et 11999. in preparation). the models can be developed to account more accurately for the increasing amount of available data.," In future, by comparing the predictions of the models with observations, especially with the redshift distributions of the SCUBA galaxies (Barger et 1999b; Lilly et 1999; Smail et 1999, in preparation), the models can be developed to account more accurately for the increasing amount of available data."
+ There is a tight correlation between the flux densities. of low-redshilt galaxies in the racio ancl far-infrarecl wavebancds (sce the review by Condon 1992)., There is a tight correlation between the flux densities of low-redshift galaxies in the radio and far-infrared wavebands (see the review by Condon 1992).
+ Thus the counts of faint ealaxies observed in the radio waveband should not be overproduced. when the SEDs of the galaxies in the 35-. 40-. 45- and 50-Ix. models presented. here are extended. into the radio waveband using this correlation.," Thus the counts of faint galaxies observed in the radio waveband should not be overproduced when the SEDs of the galaxies in the 35-, 40-, 45- and 50-K models presented here are extended into the radio waveband using this correlation."
+ It is permissible to underpredict the counts. as there will be a contribution from AGN to the faint counts. which need not be associated with powerful.restframe far-infrared) emission. from dust.," It is permissible to underpredict the counts, as there will be a contribution from AGN to the faint counts, which need not be associated with powerfulrestframe far-infrared emission from dust."
+ Partridge et ((1997) report a s.4-Gllz galaxy count of, Partridge et (1997) report a 8.4-GHz galaxy count of
+important features of the satellite evolution such as mass loss history and density. profile evolution.,important features of the satellite evolution such as mass loss history and density profile evolution.
+ In this paper we use the higher resolution simulation to investigate satellite disruption particularly focusing on the detailed. physical processes allecting satellite evolution such as. resonant dvnaniies., In this paper we use the higher resolution simulation to investigate satellite disruption particularly focusing on the detailed physical processes affecting satellite evolution such as resonant dynamics.
+ When a time-dependent force acts on a bound system such as a galaxy or a dark matter halo. resonant interactions play an important role in the svstem evolution.," When a time-dependent force acts on a bound system such as a galaxy or a dark matter halo, resonant interactions play an important role in the system evolution."
+ Recent studies of resonant dynamics in galaxy evolution claim that high resolution simulations are required to accurately reproduce these resonant ellects (Weinberg&Ixatz2007a.h) and Weinberg&Ixatz(2007a) provide a procedure to determine minimum particle number guidelines.," Recent studies of resonant dynamics in galaxy evolution claim that high resolution simulations are required to accurately reproduce these resonant effects \citep{WK07a,WK07b} and \citet{WK07a} provide a procedure to determine minimum particle number guidelines."
+ Since our idealised. high resolution simulations are designed to satisfy hese particle number guidelines. they allow us to investigate he role of resonant dynamics in satellite cisruptions.," Since our idealised, high resolution simulations are designed to satisfy these particle number guidelines, they allow us to investigate the role of resonant dynamics in satellite disruptions."
+ Resonant interactions couple a time-dependent »rturbing force. with orbits in. the/— system., Resonant interactions couple a time-dependent perturbing force with orbits in the system.
+ The requency spectrum of the time-dependent perturbing orce characterises the interaction., The frequency spectrum of the time-dependent perturbing force characterises the interaction.
+ For a satellite orbiting in its host halo. the satellites orbital frequencies. anc possibly he rate of orbital decay. determine the time-dependence of he external force.," For a satellite orbiting in its host halo, the satellite's orbital frequencies, and possibly the rate of orbital decay, determine the time-dependence of the external force."
+ A general. eccentric satellite orbit in a spherical halo has both a racial ancl azimuthal frequency. making the resonant coupling for an eccentric orbit complex.," A general, eccentric satellite orbit in a spherical halo has both a radial and azimuthal frequency, making the resonant coupling for an eccentric orbit complex."
+ Empirically. we may characterise the overall effects of the interaction as a and aPorque. A resonant shock represents coupling with the radial orbita frequency and. a resonant torque represents coupling with the azimuthal orbital frequency.," Empirically, we may characterise the overall effects of the interaction as a and a. A resonant shock represents coupling with the radial orbital frequency and a resonant torque represents coupling with the azimuthal orbital frequency."
+ A resonant shock is a generalisation of the stancare impulsive. gravitational shock.," A resonant shock is a generalisation of the standard impulsive, gravitational shock."
+ During a resonant shock. some orbits within the satellite gain energy through resonan coupling even though they are not in the impulsive limit. i.c. the time scale. of the perturbation near pericentre is much longer than the internal orbital time scale.," During a resonant shock, some orbits within the satellite gain energy through resonant coupling even though they are not in the impulsive limit, i.e. the time scale of the perturbation near pericentre is much longer than the internal orbital time scale."
+ A resonan torque couples the rotation in the external potential to orbits in the satellite., A resonant torque couples the rotation in the external potential to orbits in the satellite.
+ For a resonant torque. the magnitude of the external potential does not have to change: a change in the position angle of the satellite frame relative to the centre of the halo is sullicient to produce a torque.," For a resonant torque, the magnitude of the external potential does not have to change; a change in the position angle of the satellite frame relative to the centre of the halo is sufficient to produce a torque."
+ Resonan shocks have been previously considered. ancl included. in the impulsive. approximation as an adiabatic correction (e.g.g.mSpitzerLOST:Weinberg.1994a.b:Cinedin&Ostriker0r 1999).," Resonant shocks have been previously considered and included in the impulsive approximation as an adiabatic correction \citep[e.g.][]{Spitzer87,Weinberg94a,Weinberg94b,GO99}."
+. Llowever. resonant torques have not been. similarly considered in satellite evolution studies although they have been extensively investigated in the dynamics of barred ealaxies.," However, resonant torques have not been similarly considered in satellite evolution studies although they have been extensively investigated in the dynamics of barred galaxies."
+ In this study. we will carefully investigate these resonant elfects on satellite evolution using this distinction.," In this study, we will carefully investigate these resonant effects on satellite evolution using this distinction."
+ CGlobular. cluster. evolution in a host galaxy is well established (e.g.SpitzerLOST:Chernoll&Weinberg1990).," Globular cluster evolution in a host galaxy is well established \citep[e.g.][]{Spitzer87,Chernoff.Weinberg90}."
+. A elobular cluster experiences both tidal truncation and heating by both compressive gravitational shocks ancl ticlal shocks., A globular cluster experiences both tidal truncation and heating by both compressive gravitational shocks and tidal shocks.
+ Because satellite halo evolution in a host halo is similar to globular cluster evolution in a host galaxy. many galaxy Formation studies employ simple analytic formulae taken from these elobular cluster. evolution studies. to estimate satellite ealaxy evolution.," Because satellite halo evolution in a host halo is similar to globular cluster evolution in a host galaxy, many galaxy formation studies employ simple analytic formulae taken from these globular cluster evolution studies to estimate satellite galaxy evolution."
+ However. unlike globular clusters. the satellitehost. mass ratio is vanishinglv small.," However, unlike globular clusters, the satellite–host mass ratio is vanishingly small."
+ “Phis breaks the spatial symmetry. in mass loss. as shown in Choietal.(2007).. and changes the relative importance of resonant. coupling.," This breaks the spatial symmetry in mass loss, as shown in \citet{Choi.etal:07}, and changes the relative importance of resonant coupling."
+ In this paper. we present numerical simulation results of satellite galaxy. disruption.," In this paper, we present numerical simulation results of satellite galaxy disruption."
+ In refsec:mnethod.. we present an overview of our numerica techniques: the N-body. simulation code. the generation of initial conditions. and the relevant perturbation theory.," In \\ref{sec:method}, we present an overview of our numerical techniques: the N-body simulation code, the generation of initial conditions, and the relevant perturbation theory."
+ In relsce:cire.. we present the results of a. circular. orbi simulation.," In \\ref{sec:circ}, we present the results of a circular orbit simulation."
+ We show that resonant torque clleets result in significant satellite mass loss., We show that resonant torque effects result in significant satellite mass loss.
+ In re[sec:ecc.. we present the results of eccentric orbi simulations.," In \\ref{sec:ecc}, we present the results of eccentric orbit simulations."
+ We show that satellite heating bv eravitationa shocks at pericentre. which also includes internal structure evolution. is the dominant process responsible for disrupting 1e satellite.," We show that satellite heating by gravitational shocks at pericentre, which also includes internal structure evolution, is the dominant process responsible for disrupting the satellite."
+ In refsecistripping.. we show that the process of satellite μαrippingDpping is an outsicde-in process in satellite space.," In \\ref{sec:stripping}, we show that the process of satellite stripping is an outside-in process in satellite space."
+"5| ‘sine this finding. we suggest an explanation for the ""missing LMC stellar stream."," Using this finding, we suggest an explanation for the `missing' LMC stellar stream."
+ We also cliscuss the evolution of the satellite density profile., We also discuss the evolution of the satellite density profile.
+ In refseciniassloss.. we provide an improved analytic estimate for satellite mass loss. and we summarise in refsecisummary..," In \\ref{sec:massloss}, we provide an improved analytic estimate for satellite mass loss, and we summarise in \\ref{sec:summary}."
+ Our N-body simulations use a three-climensional selt-consistent field code (SCE.alsoknownasanexpansioncode:berg 1900)...," Our N-body simulations use a three-dimensional self-consistent field code \citep[SCF, also known as an expansion
+code;][]{cbrock72,cbrock73,ho92,Weinberg99}."
+ The expansion code calculates. bi-orthogonal basis sets of censitv-potential pais and. computes the gravitational potential of the svstem using these basis sets., The expansion code calculates bi-orthogonal basis sets of density-potential pairs and computes the gravitational potential of the system using these basis sets.
+ There are two reasons that an expansion code is an appropriate potential solver for our study., There are two reasons that an expansion code is an appropriate potential solver for our study.
+ First. the expansion basis can be chosen to follow the structure over an interesting range of scales ancl simultaneously. suppress small-scale noise.," First, the expansion basis can be chosen to follow the structure over an interesting range of scales and simultaneously suppress small-scale noise."
+ ln contrast. the noise [from two-body scaltering can arise at all scales in direct-sumamation. tree. ancl mesh based: codes.," In contrast, the noise from two-body scattering can arise at all scales in direct-summation, tree, and mesh based codes."
+ Small-scale noise can also give rise to dilfusion in conserved quantities. which can then lead to unphysical outcomes. particularly for. studies. of long-term galaxy evolution. (Weinberg&Ixatz.20072.b).," Small-scale noise can also give rise to diffusion in conserved quantities, which can then lead to unphysical outcomes particularly for studies of long-term galaxy evolution \citep{WK07a,WK07b}."
+. Second. the expansion. code is computationally ellicient: 16 computational time increases only linearly. with particle number and with modest overhead.," Second, the expansion code is computationally efficient; the computational time increases only linearly with particle number and with modest overhead."
+ Hence. the expansion code permits the use of à much larger number of particles iui most other codes for the same computational cost.," Hence, the expansion code permits the use of a much larger number of particles than most other codes for the same computational cost."
+ The expansion code is not adaptive., The expansion code is not adaptive.
+ The largest clliciency obtains when the basis resembles the galaxy., The largest efficiency obtains when the basis resembles the galaxy.
+ Otherwise. the expansion code requires a large number of basis function pairs. which introduces. small-scale noise and results in a ercater computational overhead.," Otherwise, the expansion code requires a large number of basis function pairs, which introduces small-scale noise and results in a greater computational overhead."
+ These constraints have been minimised by some recent improvements (Weinberg1999:Choietal.2007).," These constraints have been minimised by some recent improvements \citep{Weinberg99,Choi.etal:07}."
+.. First. emploving a numerical solution of the Sturm-Liouville equation. the initial galaxy model can be used as the zeroth-order basis function for the expansion code.," First, employing a numerical solution of the Sturm-Liouville equation, the initial galaxy model can be used as the zeroth-order basis function for the expansion code."
+ Then. the expansion code requires only a modest. number terms to accurately compute the potential.," Then, the expansion code requires only a modest number terms to accurately compute the potential."
+ Second. we accurately trace the densitypotential centre of the expansion during the course ofα simulation (seeChoietal.2007.for doetails)..," Second, we accurately trace the density–potential centre of the expansion during the course of a simulation \citep[see][for details]{Choi.etal:07}. ."
+II colt deusitv for CNAME clouds is 5<101? ? (note that the seusitivity of this survev is AN(IID~LOM 2).,"HI column density for CNM clouds is $5\times10^{19}$ $^{-2}$ (note that the sensitivity of this survey is $\Delta N({\rm HI})\sim
+10^{18}$ $^{-2}$ )."
+ IITOS provided statistical distributions of the fundamental parameters: I cobuun density. temperature. turbulence. and magnetic field.," HT05 provided statistical distributions of the fundamental parameters: HI column density, temperature, turbulence, and magnetic field."
+" They found that the observed. probability density function. (PDF) of ΑΠ). for sources prinarelv at high Calactic latitudes. is approximated by Nt between the Πες given by: where the subscript 20 πασάς units of 1079 7,"," They found that the observed probability density function (PDF) of $N({\rm HI})$, for sources primarely at high Galactic latitudes, is well-approximated by $N^{-1}$ between the limits given by: where the subscript 20 means units of $10^{20}$ $^{-2}$."
+" UT05 regarded the lower laut CV,554,5;,) to be real. not au observational result mnüposed by the lack of sensitivity."," 	HT05 regarded the lower limit ${\Nperp}_{20,min}$ ) to be real, not an observational result imposed by the lack of sensitivity."
+ However. the presence of three weak absorption components towards 30286/one of ITTO5 sightlucsshows that their seusitivitv estimate wass. optimistic.," However, the presence of three weak absorption components towards 3C286—one of HT05's sightlines---shows that their sensitivity estimate was optimistic."
+ ος PDF of equation 1. certainly applies to their sources with stroug absorption lines which have received a large amount of iutegration time for nieasuriug Zeeman splitting., HT05's PDF of equation \ref{Nobsobsnumbers} certainly applies to their sources with strong absorption lines which have received a large amount of integration time for measuring Zeeman splitting.
+ But it appears that there is a new population of CNAL components that is unurepresenuted in their survey., But it appears that there is a new population of CNM components that is unrepresented in their survey.
+ McBee aud Ostriker (1977) predicted the CNM to form ~95%( of the total ΠΠ., McKee and Ostriker (1977) predicted the CNM to form $\sim95\%$ of the total HI.
+ Observers have traditionally fouud sanaller colunn-deusitv fractious. for example Ieiles&Trolaud(20035). fouud that the majority of sightlincs have the CNAL cobunun deusitv fraction < 30%. while the WNAL fractionis ~61%.," Observers have traditionally found smaller column-density fractions, for example \cite{Heiles03b} found that the majority of sightlines have the CNM column density fraction $<30\%$ , while the WNM fraction is $\sim61\%$."
+ This WNM column ceusity fraction translates to a volue fraction of ~0.5., This WNM column density fraction translates to a volume fraction of $\sim0.5$.
+ Woltireetal.(2003) have updated the theory of ISAL neutral phases and slightly refined values for the WNM pressure and density: these refinements bring the observed WNM volume fraction to 0.5., \cite{Wolfire03} have updated the theory of ISM neutral phases and slightly refined values for the WNM pressure and density; these refinements bring the observed WNM volume fraction to 0.8.
+" However. sightlines towards 30286 iid 3€287 have much smaller fractious of the CNM column density. ~L% and X. respectively,"," However, sightlines towards 3C286 and 3C287 have much smaller fractions of the CNM column density, $\sim 4\%$ and $\lesssim 2\%$, respectively."
+ These echo the sunall fractions foundby ITTO5. who found 19 (out of of9) lines ofsight with no detectable CNME 3C286 was onc these. so perhaps we should take their non-doetectious to indicate a possible represcutative upper limit of a few percent for this CNAL/total III ratio.," These echo the small fractions found by HT05, who found 19 (out of 79) lines of sight with no detectable CNM; 3C286 was one of these, so perhaps we should take their non-detections to indicate a possible representative upper limit of a few percent for this CNM/total HI ratio."
+ These sinall CNAL fractions might be a problem for the theory., These small CNM fractions might be a problem for the theory.
+ Alternatively. the particular sightlines having small CNAL fractions Ιστ lie in special regious affected by energetic processes that have temporarily destroved the CNAL.," Alternatively, the particular sightlines having small CNM fractions might lie in special regions affected by energetic processes that have temporarily destroyed the CNM."
+ In fact. sightlines towards 30286 and 36287 pass through the North Calactic Pole region where Eulkarui&Fich(1985) already noticed that about of III is iufalliug towards the plane. while wari ΕΠ occupies about1056.," In fact, sightlines towards 3C286 and 3C287 pass through the North Galactic Pole region where \cite{Kulkarni85a} already noticed that about of HI is infalling towards the plane, while warm HI occupies about."
+.. This question needs further investigation., This question needs further investigation.
+ A second population of low-colui CNM clouds are those often referred to as the tiuv-seale atomic structure (TSAS)., 	A second population of low-column CNM clouds are those often referred to as the tiny-scale atomic structure (TSAS).
+ Tho TSAS size scale is inferred from time variability of TT absorption profiles against pulsus or VLBI unuagiug of extra-ealactic continui sources and ranges frou a few AU to a few lundireds of AU., The TSAS size scale is inferred from time variability of HI absorption profiles against pulsars or VLBI imaging of extra-galactic continuum sources and ranges from a few AU to a few hundreds of AU.
+ TSAS vpically las NIT) somewhat higher than our Table J compoucuts. namely from ~3.1075 to ~ὃς10M ? (Ieiles1997).," TSAS typically has $N({\rm HI})$ somewhat higher than our Table 1 components, namely from $\sim3\times10^{18}$ to $\sim2\times10^{19}$ $^{-2}$ \citep{Heiles97}."
+". Receutlv. Staniniroviéetal.(2005). found oersistent variations in the ΤΗ optical depth profiles of PSR D1929|10 which indicate structure in the cold II with NUD~2«1055 2,"," Recently, \cite{Stanimirovic05} found persistent variations in the HI optical depth profiles of PSR B1929+10 which indicate structure in the cold HI with $N({\rm HI})\sim2\times10^{18}$ $^{-2}$."
+ With these column deusities aud small sizes. the TSAS must be overpessured with respect o most of the ISM.," With these column densities and small sizes, the TSAS must be overpessured with respect to most of the ISM."
+" Thus TSAS features are thought to be very deuse and over-pressured. with »(IIT) aud P reaching values up to —104 ? and ~10° IS 5mὉ, respectively."," Thus TSAS features are thought to be very dense and over-pressured, with $n({\rm HI})$ and $P$ reaching values up to $\sim10^{4}$ $^{-3}$ and $\sim10^{6}$ K $^{-3}$, respectively."
+ Tf we assuue that our low colin deusity CNM clouds towards 3C286 and 230287 are similarly overpressured. then the inplied size scale is even smaller than the typical TSAS scale because our column deusities are somewhat sinaller.," 	If we assume that our low column density CNM clouds towards 3C286 and 3C287 are similarly overpressured, then the implied size scale is even smaller than the typical TSAS scale because our column densities are somewhat smaller."
+ Alternatively. if our CNM. components are at the standard ISM pressure. with oZ~3000 cm? K (Jenkins&Tripp2001).. then their HII volume clensities are pc30100 cin? aud the nuplied size scales are MIDsO1000 AU.," Alternatively, if our CNM components are at the standard ISM pressure, with $nT \sim 3000$ $^{-3}$ K \citep{Jenkins01}, then their HI volume densities are $n \sim 20-100$ $^{-3}$ and the implied size scales are ${N({\rm HI}) \over n({\rm HI})} \sim
+800-4000$ AU."
+ Even these larger sizes are not too nmch hieher than the inferred sizes for much of the TSAS., Even these larger sizes are not too much higher than the inferred sizes for much of the TSAS.
+ Another indication that the low cobluun density clouds could be related to TSAS comes from direct interferometric imaging., 	Another indication that the low column density clouds could be related to TSAS comes from direct interferometric imaging.
+ Over a ranee of velocities BINOS found compact enuission chumps with a FWIIM of 1 aand au intrinsic size of 30 arcsec: at an assed distance of 100 pe this corresponds to a plane-ou-the-sky size of about 3000 AU., Over a range of velocities BK05 found compact emission clumps with a FWHM of 1–2 and an intrinsic size of 30 arcsec; at an assumed distance of 100 pc this corresponds to a plane-on-the-sky size of about 3000 AU.
+ This morphology is distinctly differeut from the Uam sheet-like model of the CNM suggested bv Helles&Trolaud(2003b).. aud may poit again to a distinctly different origin of these clouds.," This morphology is distinctly different from the clumpy sheet-like model of the CNM suggested by \cite{Heiles03b}, and may point again to a distinctly different origin of these clouds."
+ The compact chuups at apparcutly random locations may bo sugeestive of fluctuations in the distribution of HIT optical depth over a wide rauge of spatial scales. rather than the preseuce of distinct. physical cutitics.," The compact clumps at apparently random locations may be suggestive of fluctuations in the distribution of HI optical depth over a wide range of spatial scales, rather than the presence of distinct physical entities."
+ If the low colhuun density HI clouds and TSAS features have a conunuon origin. then TSAS is significantly more abundant in the ISM. and with significantly lower optical depths. than what is expected theoretically if the TSAS is simply the low-size-scale exteusiou of the interstellar turbulence spectrum as cüseussed by. Deshpande(2000).," If the low column density HI clouds and TSAS features have a common origin, then TSAS is significantly more abundant in the ISM, and with significantly lower optical depths, than what is expected theoretically if the TSAS is simply the low-size-scale extension of the interstellar turbulence spectrum as discussed by \cite{Deshpande00a}."
+. Perhaps the TSAS is characterized by a rauge of column densities and pressures. with the classical absorption observations presented here sampling the low-deusity eud and the pulsar/VLBI obscrvatious. which are less seusitive. the (possibly ιο rarer) high-deusity eud.," 	Perhaps the TSAS is characterized by a range of column densities and pressures, with the classical absorption observations presented here sampling the low-density end and the pulsar/VLBI observations, which are less sensitive, the (possibly much rarer) high-density end."
+ We discuss some current theoretical approaches which can lead to the production of cold clouds in the ISM with column densities similar to what is fouud observationallv., We discuss some current theoretical approaches which can lead to the production of cold clouds in the ISM with column densities similar to what is found observationally.
+ These comprise two ecueral classes: clouds that are long-lived and can be characterized as semi-permanent structures. and transicut clouds.," These comprise two general classes: clouds that are long-lived and can be characterized as semi-permanent structures, and transient clouds."
+ Iu the selEaccepted theory of Mckee&Ostriker(1977) the CNAL resideswithin warmer phases., In the well-accepted theory of \cite{McKee77a} the CNM resideswithin warmer phases.
+ However. it," However, it"
+"background emission-line galaxies and QSOs, which retain prominent emission features in low S/N spectra.","background emission-line galaxies and QSOs, which retain prominent emission features in low S/N spectra."
+" Our aim is to compare the spatial distribution, kinematics and colours of spectroscopically confirmed Virgo and Fornax"," Our aim is to compare the spatial distribution, kinematics and colours of spectroscopically confirmed Virgo and Fornax"
+"In recent years. considerable attention has been paid to what appears to be a population of massive (M.107—10!AZ; ο. rich (Mau)MeuM)~ O.3—-0.7). rapidly star forming (M.|0-—200A; yr!) ""disk"" galaxies (flattened and apparently rotationally supported. albeit with the ratio of rotational to random velocity Vio u few) at redshifts z~2—3222222222.","In recent years, considerable attention has been paid to what appears to be a population of massive $M_{\ast}\sim10^{10}-10^{11}\,\msun$ ), gas-rich $M_{\rm gas}/(M_{\rm gas}+M_{\ast})\sim0.3-0.7$ ), rapidly star forming $\dot{M}_{\ast}\sim10-200\,\msun\,{\rm yr^{-1}}$ ) “disk” galaxies (flattened and apparently rotationally supported, albeit with the ratio of rotational to random velocity $V/\sigma\sim\,$ a few) at redshifts $z\sim2-3$."
+ The most extreme star-forming populations at these redshifts (with M.100—300042..vr appear to be major merger-induced starbursts. but these “disky” systems dominae the intermediate-Iuminosity »opulation which includes a large fr:action of the total SFR density?).," The most extreme star-forming populations at these redshifts (with $\dot{M}_{\ast}\sim100-3000\,\msun\,{\rm yr^{-1}}$ ) appear to be major merger-induced starbursts, but these “disky” systems dominate the intermediate-luminosity population which includes a large fraction of the total SFR density."
+" The morphology of the young stars (observed in the rest-frame UV and optical light) and. star-orming or molecular gas HHe vis characteristically irregular Celumpy"" or “clump-chain” morphologies}. with a signiticant Traction (tens of percent) of the ight in ~a few massive Mz). ~kpe-scale clumps or blo"," The morphology of the young stars (observed in the rest-frame UV and optical light) and star-forming or molecular gas $\alpha$ ) is characteristically irregular (“clumpy” or “clump-chain” morphologies), with a significant fraction (tens of percent) of the light in $\sim$ a few massive $\gg10^{8}\,\msun$ ), $\sim$ kpc-scale clumps or blobs."
+ Theoretically. the origin of these morphologies is controversial.," Theoretically, the origin of these morphologies is controversial."
+ Some early models argued they were irregular because of ongoing mergers(2): observations of local analogues (and studies in which these systems are mock-observed as if they were at high redshifts) suggest that this is true for least some fraction of the observed systems(222)., Some early models argued they were irregular because of ongoing mergers; observations of local analogues (and studies in which these systems are mock-observed as if they were at high redshifts) suggest that this is true for least some fraction of the observed systems.
+. However. the constraints on the structure and kinematics of the massive ς—2 systems suggest that a sizeable fraction are more (222). although this is not definitive evidence against mergers (?2)).," However, the constraints on the structure and kinematics of the massive $z \sim 2$ systems suggest that a sizeable fraction are more ), although this is not definitive evidence against mergers )."
+ And the merger rate may not be sufficiently high to explain the abundance of these intermediate-SFR systems222?)., And the merger rate may not be sufficiently high to explain the abundance of these intermediate-SFR systems.
+. An alternative explanation for the observed morphologies is that clumpiness stems from disk fragmentation., An alternative explanation for the observed morphologies is that clumpiness stems from disk fragmentation.
+ Any gas disk in which cooling is efficient (cooling time less than the dynamical time so that turbulent support dissipates in a crossing time. which is easily satistied here) will fragment at the Toomre scale. Ry—oπο) or M;~atποM(yz for disks which equilibrate at marginal stability (Toomre Q— D. this is just Ry~fiiRuin CM;~PNMaa In the Milky Way. these clumps correspond to massive GMCs (8 1θ0ρο. M~10AY. ο.," Any gas disk in which cooling is efficient (cooling time less than the dynamical time so that turbulent support dissipates in a crossing time, which is easily satisfied here) will fragment at the Toomre scale, $R_{T}\sim \sigma^{2}/\pi\,G\,\Sigma$ or $M_{T}\sim \sigma^{4}/\pi\,G^{2}\,\Sigma$; for disks which equilibrate at marginal stability (Toomre $Q\sim1$ ), this is just $R_{T}\sim f_{\rm gas}\,R_{\rm disk}$ $M_{T}\sim f_{\rm gas}^{3}\,M_{\rm disk}$ In the Milky Way, these clumps correspond to massive GMCs $R\sim100\,$ pc, $M\sim10^{6}\,\msun$ )."
+" But in a massive. gas-rich disk. the scale-length and mass of the resulting clumps can be as large as ~kpe and ~105—10""Αι. respectively. similar to the structures observed."," But in a massive, gas-rich disk, the scale-length and mass of the resulting clumps can be as large as $\sim$ kpc and $\sim10^{8}-10^{9}\,\msun$, respectively, similar to the structures observed."
+ This instability has been well-known for ~40 years., This instability has been well-known for $\sim40$ years.
+ The implications of such clumping in the ISM are less clear. however.," The implications of such clumping in the ISM are less clear, however."
+ If there were no feedback from stars to expel mass from GMCs/clumps. they would (locally) collapse in a few free-fall times and turn most of their mass into stars.," If there were no feedback from stars to expel mass from GMCs/clumps, they would (locally) collapse in a few free-fall times and turn most of their mass into stars."
+ They would then, They would then
+ , 
+We observed the lowest transition of para-water. the 1j;—Ooo transition at 1113.3430 GHz line. in the upper sideband of the Band 4b HIFI receiver.,"We observed the lowest transition of para-water, the $1_{11} - 0_{00}$ transition at 1113.3430 GHz line, in the upper sideband of the Band 4b HIFI receiver."
+ This observation. of duration 1167 s including overheads. was carried out on 2010 March. 5. using the dual beam switeh. (DBS) mode and the Wide Band Spectrometer (WBS).," This observation, of duration 1167 s including overheads, was carried out on 2010 March 5, using the dual beam switch (DBS) mode and the Wide Band Spectrometer (WBS)."
+" The WBS has a spectral resolution of 1.1 MHz. corresponding to a velocity resolution of 0.30kms""! at the frequency of the 1),—Quo transition."," The WBS has a spectral resolution of 1.1 MHz, corresponding to a velocity resolution of $0.30\, \rm km\,s^{-1}$ at the frequency of the $1_{11} - 0_{00}$ transition."
+" The telescope beam. of half-power-beam-width (HPBW) ~20"". was centered on V Cygni at coordinates a=20h4Im18.3s.048""08'29""(12000)."," The telescope beam, of half-power-beam-width (HPBW) $\sim 20^{\prime\prime}$, was centered on V Cygni at coordinates $\rm \alpha=20h\,41m\,18.3s, \delta = 48^0\, 08^\prime \,29^{\prime\prime} (J2000)$."
+ The reference positions for this observation were located at offsets of 3’ on either side of the source., The reference positions for this observation were located at offsets of $3^{\prime}$ on either side of the source.
+ The data were processed using the standard HIFI pipeline to Level 2. providing fully calibrated spectra of the source.," The data were processed using the standard HIFI pipeline to Level 2, providing fully calibrated spectra of the source."
+ The Level 2 data were analysed further using the Herschel Interactive Processing Environment (HIPE: Ott 2010). version 2.4. along with ancillary IDL routines that we have developed.," The Level 2 data were analysed further using the Herschel Interactive Processing Environment (HIPE; Ott 2010), version 2.4, along with ancillary IDL routines that we have developed."
+ Having found the signals measured in the two orthogonal polarizations to be in excellent agreement. we combined them to obtain an average spectrum.," Having found the signals measured in the two orthogonal polarizations to be in excellent agreement, we combined them to obtain an average spectrum."
+" Figure | shows the WBS spectrum of para-H>O 1),—Oo obtained toward V Cygni. with the frequency scale expressed as Doppler velocities relative to the Local Standard of Rest (LSR) and the intensity scale expressed as antenna temperature."," Figure 1 shows the WBS spectrum of $_2$ O $1_{11}-0_{00}$ obtained toward V Cygni, with the frequency scale expressed as Doppler velocities relative to the Local Standard of Rest (LSR) and the intensity scale expressed as antenna temperature."
+subtracted. The vertical dashed line indicates the LSR velocity of the source. as determined by Bieging Wilson (2001. hereafter BWOI) from observations of the CO J=2—| line.," The vertical dashed line indicates the LSR velocity of the source, as determined by Bieging Wilson (2001, hereafter BW01) from observations of the CO $J=2-1$ line."
+ The integrated antenna temperature is found to be 1.69x0.17 K km s!.., The integrated antenna temperature is found to be $1.69 \pm 0.17$ K km $^{-1}$.
+ In modeling the water emission from V Cygnr. we have used the methods described by Gonzállez-Alfonso. Neufeld Melnick (2007. hereafter GNM).," In modeling the water emission from V Cygni, we have used the methods described by Gonzállez-Alfonso, Neufeld Melnick (2007, hereafter GNM)."
+ To determine. the HO rotational populations and the resultant HO emission spectrum. we included the effects of radiative pumping -- through the v»—] and và=| vibrational states — and of collisional excitation by H2. together with a treatment of radiative transfer based upon that discussed by Gonzállez-Alfonso Cernicharo (1997).," To determine the $_2$ O rotational populations and the resultant $_2$ O emission spectrum, we included the effects of radiative pumping -- through the $\nu_2=1$ and $\nu_3=1$ vibrational states – and of collisional excitation by $_2$, together with a treatment of radiative transfer based upon that discussed by Gonz\'allez-Alfonso Cernicharo (1997)."
+ The assumed parameters for the source are summarized in Table | and discussed below., The assumed parameters for the source are summarized in Table 1 and discussed below.
+ A variety of distance estimates for V Cygni have appeared in the literature., A variety of distance estimates for V Cygni have appeared in the literature.
+ BWOI obtained values of 406 pe and 456 pe respectively by using the P—K (period — absolute K magnitude) and P—L (Period — bolometric luminosity) relationships presented by Groenewegen Whitelock (1996; hereafter GW96)., BW01 obtained values of 406 pc and 456 pc respectively by using the $P-K$ (period – absolute K magnitude) and $P-L$ (Period – bolometric luminosity) relationships presented by Groenewegen Whitelock (1996; hereafter GW96).
+ However. our own fit to the spectral energy distribution (SED). shown in Figure 2. implies that a distance of 342 pe is required to fit the GW96 P-L relationship.," However, our own fit to the spectral energy distribution (SED), shown in Figure 2, implies that a distance of 342 pc is required to fit the GW96 $P-L$ relationship."
+" Here. we used a dust radiative transfer model to fit a combination of flux measurements fromνόσο,Hipparcos.MSX. 2MASS. ISO.IRAS. and JCMT."," Here, we used a dust radiative transfer model to fit a combination of flux measurements from, 2MASS, and JCMT."
+ This model. full details of which will be presented in a future publication (Schmidt et 22010). assumes a power-law distribution of grain radii with index —3.5 between 0.005 and 0.25 um. a maximum dust temperature of 1050 K. an inner radius for the dust shell of 2x10'* em. and a total optical depth Ay=9.45 mag.," This model, full details of which will be presented in a future publication (Schmidt et 2010), assumes a power-law distribution of grain radii with index $-3.5$ between 0.005 and 0.25 $\mu$ m, a maximum dust temperature of 1050 K, an inner radius for the dust shell of $2 \times 10^{14}$ cm, and a total optical depth $A_V = 9.45$ mag."
+ A smaller distance estimate (370 pc) than those of BWOI was also obtained by Schétier Olofsson (2000)., A smaller distance estimate (370 pc) than those of BW01 was also obtained by Schöiier Olofsson (2000).
+ All of the values mentioned above are consistent with the Hipparcos-measured parallax of 3.69+1.77 mas. which implies a lo distance range of 180 — 520 pe.," All of the values mentioned above are consistent with the }-measured parallax of $3.69 \pm 1.77$ mas, which implies a $1\sigma$ distance range of 180 – 520 pc."
+ Accordingly. we adopt a value of 400 pe for the estimated distance of V Cyent. with a likely uncertainty — 20%.," Accordingly, we adopt a value of 400 pc for the estimated distance of V Cygni, with a likely uncertainty $~\sim 20\%$ ."
+ This estimate of the distance then requires a bolometric luminosity of 8500L. to match the observed SED shown in Figure 2. and a stellar radius of 3.42x10!em. given," This estimate of the distance then requires a bolometric luminosity of $8500\,L_\odot$ to match the observed SED shown in Figure 2, and a stellar radius of $3.42 \times 10^{13} \, \rm cm,$ given"
+were compared to the actual observations.,were compared to the actual observations.
+ These. model observations are based on a number of geometrical and kinematical parameters that describe the IHE disk. such as the rotation curve and the inclination ancl position angles as a function of radius.," These model observations are based on a number of geometrical and kinematical parameters that describe the HI disk, such as the rotation curve and the inclination and position angles as a function of radius."
+ The other input. parameters. are: the central position and the systemic velocity. which can be derived with reasonable accuracy from the observations: the velocity dispersion as a function of radius. which we kept at a lixed valueof6 kms|J. typical of dwarf irregulars (Weldrake et al.," The other input parameters are: the central position and the systemic velocity, which can be derived with reasonable accuracy from the observations; the velocity dispersion as a function of radius, which we kept at a fixed value of 6 km $^{-1}$, typical of dwarf irregulars (Weldrake et al."
+ 2003): the surface density as a function of radius. which is derived. from the total HE map: the scale. height of the LIE disk. which we assumed to be 300 pc. a typical value for clhwarl galaxies: the racial motions as a function of radius.," 2003); the surface density as a function of radius, which is derived from the total HI map; the scale height of the HI disk, which we assumed to be 300 pc, a typical value for dwarf galaxies; the radial motions as a function of radius."
+ With the choice of parameters from Section 3. the position angle was incorrect. ancl the inclination was underestimated. as can be seen in Figs.," With the choice of parameters from Section \ref{velfi_first}, the position angle was incorrect, and the inclination was underestimated, as can be seen in Figs."
+ 4 and 5.., \ref{velfi} and \ref{mom0_velfi}.
+ Phese two parameters were then iteratively changed in order to give a better representation of the total LEE map: the results are show in Figs., These two parameters were then iteratively changed in order to give a better representation of the total HI map; the results are show in Figs.
+" 6 and 7 [or the high- and low-resolution data cubes respectively, in which the models clearly match the data much better."," \ref{mom0} and \ref{mom0_lores}
+ for the high- and low-resolution data cubes respectively, in which the models clearly match the data much better."
+ Compared to the position angles shown in Fig. 3..," Compared to the position angles shown in Fig. \ref{rotcur},"
+ the new ones are dillerent by values ranging from in the inner parts to about iin the outer ones., the new ones are different by values ranging from in the inner parts to about in the outer ones.
+ The inclination was changed to a fixed. value. of70... in. order to obtain good agreement with the total LIE maps. both at high and. low resolution.," The inclination was changed to a fixed value of, in order to obtain good agreement with the total HI maps, both at high and low resolution."
+ The racial dependence of the position angle is shown in the middle panel of Fie. 12.., The radial dependence of the position angle is shown in the middle panel of Fig. \ref{parametri}.
+ In Fig., In Fig.
+ Swe show low- total intensity maps with different choices of the orientation. parameters., \ref{pa_incl} we show low-resolution total intensity maps with different choices of the orientation parameters.
+ Baseck upon these models. we estimate the uncertainty of the inclination. ancl position angles to be about4.," Based upon these models, we estimate the uncertainty of the inclination and position angles to be about."
+. In the innermost [ew data points we conservatively doubled this value. given the smaller number of beams that sample the total LIE map and the higher complexity of the LL kinematics ancl distribution.," In the innermost few data points we conservatively doubled this value, given the smaller number of beams that sample the total HI map and the higher complexity of the HI kinematics and distribution."
+ Note that the assumption of a constant velocity dispersion of 6 km + is consistent with the observations: even though it is apparently increasing towards the centre. a velocity dispersion map of a model data cube with a constant velocity dispersion of6 kms + exhibits a very similar behaviour (Fie. 9)).," Note that the assumption of a constant velocity dispersion of 6 km $^{-1}$ is consistent with the observations: even though it is apparently increasing towards the centre, a velocity dispersion map of a model data cube with a constant velocity dispersion of 6 km $^{-1}$ exhibits a very similar behaviour (Fig. \ref{vdisp}) ),"
+ because of the broadening of the velocity. profiles due to the beam where the velocity gradients are the largest., because of the broadening of the velocity profiles due to the beam where the velocity gradients are the largest.
+ In order to account for the kinematies. ancl therefore o explain the cilferenee between morphological anc kinematical position angles. one must resort to non-circular motions.," In order to account for the kinematics, and therefore to explain the difference between morphological and kinematical position angles, one must resort to non-circular motions."
+ A data cube built with the inclination ancl position angles that match the total HE map is inconsistent with he observations if non-circular motions are kept to zero (Vig. 10))., A data cube built with the inclination and position angles that match the total HI map is inconsistent with the observations if non-circular motions are kept to zero (Fig. \ref{cubes_hires_vr0}) ).
+ An estimate of their amplitude can be found by »forming a harmonic decomposition of the velocity field (Schoenmakers ct al., An estimate of their amplitude can be found by performing a harmonic decomposition of the velocity field (Schoenmakers et al.
+ LOOT. Wong. Blitz Bosma 2004). keeping the position ancl inclination angles fixed at the values derived in the previous Section. since leaving them. as [ree parameters leads (ο unconstrained fits.," 1997, Wong, Blitz Bosma 2004), keeping the position and inclination angles fixed at the values derived in the previous Section, since leaving them as free parameters leads to unconstrained fits."
+ With this method the observed line-of-sight velocity Vy. is mocelled N where vc ds the azimuthal angle., With this method the observed line-of-sight velocity $V_{\rm los}$ is modelled by where $\psi$ is the azimuthal angle.
+ 1n. this analysis we considered. terms up to jf=3., In this analysis we considered terms up to $j=3$.
+ The results are illustrated in Figs., The results are illustrated in Figs.
+" H1. for the high-resolution cube (the low-resolution cube gives similar results). where the 5, component (indicativo of radial motions) is non-zero and varies approximately between 5 and 10 km with a region between aand wwhere they are close to 5I."," \ref{harmonic} for the high-resolution cube (the low-resolution cube gives similar results), where the $s_1$ component (indicative of radial motions) is non-zero and varies approximately between 5 and 10 km $^{-1}$, with a region between and where they are close to 5."
+ Phe very small values of c» and so indicate that the kinematics of 33741 is very symmetric., The very small values of $c_2$ and $s_2$ indicate that the kinematics of 3741 is very symmetric.
+ Wong ct al. (, Wong et al. (
+2004) show that an error in the fitted. inclination. angle causes the c; term to be proportional to e;sin/. which seems to be the case here.,"2004) show that an error in the fitted inclination angle causes the $c_3$ term to be proportional to $c_1 ~ 
+\sin i$, which seems to be the case here."
+ llowever. as seen in Fig. 6..," However, as seen in Fig. \ref{mom0},"
+ the inclination is quite well constrained from the total HIE. map., the inclination is quite well constrained from the total HI map.
+" We interpret. this SPULLOUS C4X6€,sin? term as the dillerence between an analysis that assumes an infinitely thin cüsk (the harmonic decomposition of the velocity field) and one that takes this ellect into account (the data cube modelling).", We interpret this spurious $c_3 \propto c_1~\sin i$ term as the difference between an analysis that assumes an infinitely thin disk (the harmonic decomposition of the velocity field) and one that takes this effect into account (the data cube modelling).
+ Indeed. within the framework of the harmonic analysis this termi. would indicate an incorrect value of the inclination.," Indeed, within the framework of the harmonic analysis this term would indicate an incorrect value of the inclination."
+ However. 1f harmonic decomposition assumes an infinitely thin isk: if the disk thickness is actually non-negligible. this assumption will result. in an underestimated inclination angle.," However, the harmonic decomposition assumes an infinitely thin disk: if the disk thickness is actually non-negligible, this assumption will result in an underestimated inclination angle."
+" Therefore. if one fixes the inclination at its correct value (accounting for the thickness of the disk). the output harmonic decomposition procedure will display à spurious ""OaXC, sin elfect."," Therefore, if one fixes the inclination at its correct value (accounting for the thickness of the disk), the output harmonic decomposition procedure will display a spurious $c_3 \propto c_1~\sin i$ ” effect."
+ Note that the plotted errors in Fig., Note that the plotted errors in Fig.
+ 11. are formal errors coming from the fitting procedure and they are likely to be smaller than the real uncertainties in the parameters (see ee. Bureau Carignan 1999)., \ref{harmonic} are formal errors coming from the fitting procedure and they are likely to be smaller than the real uncertainties in the parameters (see e.g. Bureau Carignan 1999).
+ Starting from the results of the harmonic decomposition of the velocity Ποια we built different model cata cubes. iteratively improving the input parameters of the mocdels until acceptable results were found.," Starting from the results of the harmonic decomposition of the velocity field we built different model data cubes, iteratively improving the input parameters of the models until acceptable results were found."
+ Some parameters were adjusted. at. an earlier stage (the position ancl inclination angles). while others (he rotation anc radial velocities) needed: some additional small changes.," Some parameters were adjusted at an earlier stage (the position and inclination angles), while others (the rotation and radial velocities) needed some additional small changes."
+ An automatic and, An automatic and
+Fiewe Ὁ plots the color-color diagram in the Johuson/IXrou-Cousius photometric svsteu for a white dwart of mass O0O.7A/.. with ages indicated aloug the sequence.,Figure 3 plots the color-color diagram in the Johnson/Kron-Cousins photometric system for a white dwarf of mass $0.7M_{\odot}$ with ages indicated along the sequence.
+ The wild deviatious from black body colors are evident in this diagram as the oldest aud coolest white dwarts get dramatically blucr in the (WF) color aud somewhat bluer in (WR)., The wild deviations from black body colors are evident in this diagram as the oldest and coolest white dwarfs get dramatically bluer in the $(V-I)$ color and somewhat bluer in $(V-R)$.
+ When white dwarts are observed in au open or globular cluster it is not strictly correct to compare their location iu the cluster coloranagnuitude diagram with a theoretical 'ooling sequence of sonie Inass as lias generally been doue oei the past (ee. Richer 1995. 1997. 1998: Cool 1996: Reuziu 1996).," When white dwarfs are observed in an open or globular cluster it is not strictly correct to compare their location in the cluster color-magnitude diagram with a theoretical cooling sequence of some mass as has generally been done in the past (e.g. Richer 1995, 1997, 1998; Cool 1996; Renzini 1996)."
+ The reason for this is that the oldest white chwarts iu these clusters have evolved from the most massive stars originallv iu the cluster (up to the aN nass that produces white dwarfs). aud because OLE Massive progenitors produce more massive renmauts. the older white dwarfs should be more massive.," The reason for this is that the oldest white dwarfs in these clusters have evolved from the most massive stars originally in the cluster (up to the maximum mass that produces white dwarfs), and because more massive progenitors produce more massive remnants, the older white dwarfs should be more massive."
+ This has ecnerally not. been a problem with the clusters observed thus far as the range in white chwarf masses has Όσοι relatively siuall. aud. in any case. the initial mass function of clusters is expected to vield mauy fewer massive stars and hence few massive white chwarfs.," This has generally not been a problem with the clusters observed thus far as the range in white dwarf masses has been relatively small, and, in any case, the initial mass function of clusters is expected to yield many fewer massive stars and hence few massive white dwarfs."
+ However. when larec eround-based telescopes or UST eventually penetrate to the termination poiut of the white dxvarf cooling sequence in a elobular cluster. aud thus cover a wide range in white dwarf masses. it will be extremely inportaut to have white caf isochroucs ready to interpret the data as opposed to just cooling sequences Ημ ," However, when large ground-based telescopes or HST eventually penetrate to the termination point of the white dwarf cooling sequence in a globular cluster, and thus cover a wide range in white dwarf masses, it will be extremely important to have white dwarf isochrones ready to interpret the data as opposed to just cooling sequences for some mass."
+For these reasons we have calculated isochrones for white dwarfs in star chisters with a wide range in age., For these reasons we have calculated isochrones for white dwarfs in star clusters with a wide range in age.
+ All the isochrones shown are derived from solar metallicity models., All the isochrones shown are derived from solar metallicity models.
+ The isochrones were constructed by (a) beeiuniug with a white dwarf mass of 0.937... the iiaxinimn mass model that we had available. aud (b) usiug au initial-Bnal mass relation constructed from Uerwie’s 1995 data at the ligh nass end mated to the results from Cübson 1999 for M67 aud MI at the low mass eud to determine the mass of the maim sequence progenitor.," The isochrones were constructed by (a) beginning with a white dwarf mass of $0.9M_{\odot}$, the maximum mass model that we had available, and (b) using an initial-final mass relation constructed from Herwig's 1995 data at the high mass end mated to the results from Gibson 1999 for M67 and M4 at the low mass end to determine the mass of the main sequence progenitor."
+ Iu usiug these data we are musing metal rich and metal poor relations. however. at the moment this is all that can be done if we wish to stick with cupirical results. (," In using these data we are mixing metal rich and metal poor relations, however, at the moment this is all that can be done if we wish to stick with empirical results. ("
+0) The stellar evolutionary models of Doimiuguez 1999 were then cluploved to deterimine the lifetime of the main sequence star (κ) up to the eud of the ACB. (,c) The stellar evolutionary models of Dominguez 1999 were then employed to determine the lifetime of the main sequence star $A_{ms}$ ) up to the end of the AGB. (
+"d) The age of the white dwart is then simply T;,;,«ων where των i5 the age of the isochrone that we are calculating. (",d) The age of the white dwarf is then simply $T_{iso} - A_{ms}$ where $T_{iso}$ is the age of the isochrone that we are calculating. (
+"0) The absolute uaenitude and color of the white dwarf of interest was hen obtained from the cooling model for a white dwarf of nass OAL. and age Tye, - Ais. C",e) The absolute magnitude and color of the white dwarf of interest was then obtained from the cooling model for a white dwarf of mass $0.9 M_{\odot}$ and age $T_{iso}$ - $A_{ms}$. (
+F) We then decremented he mass. interpolating within the models. aud repeated he process until a white dwarf mass of 0.5AL.. was reached (the mininumuau white dwarf mass model available). at which voit the caleulations were halted.,"f) We then decremented the mass, interpolating within the models, and repeated the process until a white dwarf mass of $0.5M_{\odot}$ was reached (the minimum white dwarf mass model available), at which point the calculations were halted."
+ Figure [| illustrates these isochrones for a range of ages ikelv to be of interest in any application., Figure 4 illustrates these isochrones for a range of ages likely to be of interest in any application.
+ The hook to the due in the isochrones for ages less than about 7 Cors is not due to the effects of Πω opacity (these stars are too hot for Πο to forma) but is caused by the fact that the white dwarfs at the bottoms of these curves come froni massive main sequence stars., The hook to the blue in the isochrones for ages less than about 7 Gyrs is not due to the effects of $H_2$ opacity (these stars are too hot for $H_2$ to form) but is caused by the fact that the white dwarfs at the bottoms of these curves come from massive main sequence stars.
+ These produce heavier white dwarfs which follow cooling sequences which lie below those of lighter degenerates (more massive white dwiirfs have sialler radii and are thus less huninous at a given temperature)., These produce heavier white dwarfs which follow cooling sequences which lie below those of lighter degenerates (more massive white dwarfs have smaller radii and are thus less luminous at a given temperature).
+ It is oulv for ages older than 8 Cars where the effect of the FZ Opacity Is sen., It is only for ages older than 8 Gyrs where the effect of the $H_2$ opacity is seen.
+ Figure 5 illustrates some detail for the 11 Gyr isochrone. showing the mass of the white dwarf itself aud that of its main sequence precursor.," Figure 5 illustrates some detail for the 11 Gyr isochrone, showing the mass of the white dwarf itself and that of its main sequence precursor."
+ Tables 3 and { list selected isochrones iu both the Joliusou-I&rou/C'ousius svsteni aud in the UST filters., Tables 3 and 4 list selected isochrones in both the Johnson-Kron/Cousins system and in the HST filters.
+" The columns in these tables are the white dwarf mass CAL). the mass of the progenitor (AL..) and Z5.,,;. A. VW Ry aud (V. D) of the white dwarf."," The columns in these tables are the white dwarf mass $M_{\odot}$ ), the mass of the progenitor $M_{\odot}$ ) and $T_{eff}$, $M_V$, $V-R$ ) and $V-I$ ) of the white dwarf."
+ Details of the initialfinal mass relation that we used cau be obtaimed from these Tables., Details of the initial-final mass relation that we used can be obtained from these Tables.
+"events wilh simple ""svnchrotron. bubble’ models.",events with simple `synchrotron bubble' models.
+ In. such cases the decay of the radio emission arises due to adiabatic expansion losses. and so weaker radio emission is associated with increasingly large physical structures which display less and less variability on short timescales.," In such cases the decay of the radio emission arises due to adiabatic expansion losses, and so weaker radio emission is associated with increasingly large physical structures which display less and less variability on short timescales."
+ Such events are also physically decoupled and observationally uncorrelated with the accretion disc emission following the dramatic events of the outburst., Such events are also physically decoupled and observationally uncorrelated with the accretion disc emission following the dramatic events of the outburst.
+ For example. 10 davs after ejection. material travelling at ~0.96 will be 107 em (=1500 A.U. =1027 Sehwarzschild radii) from the central black hole.," For example, 10 days after ejection, material travelling at $\sim
+0.9c$ will be $> 10^{16}$ cm $\equiv 1500$ A.U. $\equiv 10^{10}$ Schwarzschild radii) from the central black hole."
+ Phe picture is not quite so simple however. as several transients whose radio emission is dominated by optically thin ejections sometimes briefly. display inverted. spectra (Ixuulkers et al.," The picture is not quite so simple however, as several transients whose radio emission is dominated by optically thin ejections sometimes briefly display inverted spectra (Kuulkers et al."
+ 1999)., 1999).
+" ""his is interpreted as the ejection of new components which are initially optically thick at radio wavelengths. but which rapidly. become optically thin."," This is interpreted as the ejection of new components which are initially optically thick at radio wavelengths, but which rapidly become optically thin."
+ Such multiple ejections of discrete. blobs has been directly. observed. in c.g. GIUS 1915|105 (Mirabel guez 1994: Fender ct al., Such multiple ejections of discrete blobs has been directly observed in e.g. GRS 1915+105 (Mirabel guez 1994; Fender et al.
+ 19992)., 1999a).
+ lt has long been accepted that radio emission from. X-ray binaries is svnchrotron emission from material ejected from. the system (LHjellming Ian 1995 and references therein)., It has long been accepted that radio emission from X-ray binaries is synchrotron emission from material ejected from the system (Hjellming Han 1995 and references therein).
+ Such ejections. at. relativistic bulk velocities. have been directly. observed in several cases (Hljellming Lan 1995: Mirabel guez 1999: Fender 2000 and references therein).," Such ejections, at relativistic bulk velocities, have been directly observed in several cases (Hjellming Han 1995; Mirabel guez 1999; Fender 2000 and references therein)."
+ llowever. a simple homosgenous svnchrotronemitting source should exhibit à two component spectrum. with a spectral index of (2.5 below some frequency at which. self-absorption becomes significant. and a above this frequency (with a turnover within one decade in frequency).," However, a simple homogenous synchrotron–emitting source should exhibit a two component spectrum with a spectral index of +2.5 below some frequency at which self-absorption becomes significant, and $\alpha$ above this frequency (with a turnover within one decade in frequency)."
+" For a power-law distribution of electrons of the form Αα)xE""dE. this optically thin spectral index isa=(1p)/2."," For a power-law distribution of electrons of the form $N(E) dE \propto
+E^{-p} dE$, this optically thin spectral index is $\alpha = (1-p)/2$."
+ Observed optically thin spectral indices are twpically in the range Lsas0.5. corresponding to electron energy indices of 28pS3.," Observed optically thin spectral indices are typically in the range $-1 \la \alpha \la -0.5$, corresponding to electron energy indices of $2 \la p \la 3$."
+ This range is approximately consistent with both simple theories of particle acceleration. and the observed cosmic-ray energy index (Blandford Eichler 1987)., This range is approximately consistent with both simple theories of particle acceleration and the observed cosmic-ray energy index (Blandford Eichler 1987).
+ The radio spectra. observed. from these black ho systems in the Low/Llarcl state is however quite diferent. showing a llat spectrum which probably extends to very high frequencies.," The radio spectra observed from these black hole systems in the Low/Hard state is however quite different, showing a flat spectrum which probably extends to very high frequencies."
+" Blancllord Ixónnigl. (1979) showed. in a model developed: for. AGN. that a simple ‘isothermal’ conicaljet can produce a Hat spectrum even with an electron distribution N(dExE.ας, "," Blandford Könnigl (1979) showed, in a model developed for AGN, that a simple `isothermal' conical jet can produce a flat spectrum even with an electron distribution $N(E) dE \propto E^{-2} dE$."
+Reynolds (1982) explored in more detail the observed spectra from winds and jets with a variety of ecometrics. magnetic fields and energeties (see also Cawthorne 1991 for a review).," Reynolds (1982) explored in more detail the observed spectra from winds and jets with a variety of geometries, magnetic fields and energetics (see also Cawthorne 1991 for a review)."
+ Hjellming Johnston (1988) and Faleke Biermann (1996. 1999) have cliseussect the application of such models to N-rav. binaries.," Hjellming Johnston (1988) and Falcke Biermann (1996, 1999) have discussed the application of such models to X-ray binaries."
+ Given that we recover the negative spectral index curing major optically thin outbursts. some form of the partially conical jet model does seem the most likely origin of the flat spectral component in X-ray. binaries also.," Given that we recover the negative spectral index during major optically thin outbursts, some form of the partially conical jet model does seem the most likely origin of the flat spectral component in X-ray binaries also."
+ In the simplest case. a flat spectrum implies a characteristic size scale at any frequeney which is proportional tov +.," In the simplest case, a flat spectrum implies a characteristic size scale at any frequency which is proportional to $\nu^{-1}$."
+ Thus if 1¢ Hat spectra extend from radio to near-infrared or optica requencies. this implies that the jet must be self-similar over 10 same range in physical size. ic.," Thus if the flat spectra extend from radio to near-infrared or optical frequencies, this implies that the jet must be self-similar over the same range in physical size, ie."
+ z5 orders of magnitude., $\ga 5$ orders of magnitude.
+ In Fender et al. (, In Fender et al. (
+2000) it was noted that the flat. spectra 'omponent observed. from Cvg N-1 was much flatter than wt observed from ‘Hat spectrum! GN.,2000) it was noted that the flat spectral component observed from Cyg X-1 was much flatter than that observed from `flat spectrum' AGN.
+ This is primarily ue to the high-frequency. turnover. which occurs aroun 1e millimetre band for AGN (Bloom et al.," This is primarily due to the high-frequency turnover, which occurs around the millimetre band for AGN (Bloom et al."
+ 1994). not being observed for (νο X-1 (or any other X-ray. binary to date).," 1994), not being observed for Cyg X-1 (or any other X-ray binary to date)."
+ As 1e highest-frequencey. emission will arise from the smalles ohvsical scales for a conical jet or similar model. this lack of observed turnover in X-ray binaries is probably. due to uieher densities and/or magnetic fields in the accretion Low around a LOAL. black hole as compared to the z:LOAD. ack holes in AGN.," As the highest-frequency emission will arise from the smallest physical scales for a conical jet or similar model, this lack of observed turnover in X-ray binaries is probably due to higher densities and/or magnetic fields in the accretion flow around a $\sim 10$ $_{\odot}$ black hole as compared to the $\geq 10^{6}$ $_{\odot}$ black holes in AGN."
+ Lf the turnover for the X-ray. binaries ranspires to be around the optical or near-infrared. bands. his would implv an approximate empirical scaling of the ugh frequency cut-oll: μιςκMig.; ," If the turnover for the X-ray binaries transpires to be around the optical or near-infrared bands, this would imply an approximate empirical scaling of the high frequency cut-off, $\nu_{\rm HIGH} \propto M_{\rm BH}^{1/2}$."
+Earlier interpretations of the μαι radio spectra were very closely related: to these conical jet models., Earlier interpretations of the flat radio spectra were very closely related to these conical jet models.
+" The flat spectral components observed from Gis 2023|338 and GRO J0422]32 have been referred to as ""second stage racio emission (Iljellming Han 1995). and it was suggested that they originate in à wind from the accretion disc through which the observer sees to dillerent. depths as a function of frequency (Lijellming Elan 1995)."," The flat spectral components observed from GS 2023+338 and GRO J0422+32 have been referred to as `second stage' radio emission (Hjellming Han 1995), and it was suggested that they originate in a wind from the accretion disc through which the observer sees to different depths as a function of frequency (Hjellming Han 1995)."
+ Furthermore it was suggested that the slowlydecreasing flux density was as a result of a decreasing physical size scale., Furthermore it was suggested that the slowly–decreasing flux density was as a result of a decreasing physical size scale.
+ Both a spherical wind and conical jet will have an electron density which falls as r7 (for no pair processes). and. will produce analogous spectra (Revnolcls 1982).," Both a spherical wind and conical jet will have an electron density which falls as $r^{-2}$ (for no pair processes), and will produce analogous spectra (Reynolds 1982)."
+ However. the direct. observations of apparently collimated jets from Cve δα. LE 1740.7-2042 and GIU 1758-258 provide strong evidence for a collimated geometry.," However, the direct observations of apparently collimated jets from Cyg X-1, 1E 1740.7-2942 and GRS 1758-258 provide strong evidence for a collimated geometry."
+ Furthermore. the observed. linear polarisation would not arise in a spherically symmetric source.," Furthermore, the observed linear polarisation would not arise in a spherically symmetric source."
+ The radiative luminosity is the only quantity we can directly measure from the flat spectral component., The radiative luminosity is the only quantity we can directly measure from the flat spectral component.
+ As non-racliative (i.e. acliahbatic expansion) losses are likely to dominate. this will be a very conservative lower limit on the power into the jet.," As non-radiative (i.e. adiabatic expansion) losses are likely to dominate, this will be a very conservative lower limit on the power into the jet."
+ The observed. radiative luminosity of Cvg N-1 is ~2107 erg | when observed up to 15 Gllz in the radio band (for a distance of ~2 kpe)., The observed radiative luminosity of Cyg X-1 is $\sim 2 \times 10^{30}$ erg $^{-1}$ when observed up to 15 GHz in the radio band (for a distance of $\sim 2$ kpc).
+ However. as arguec above. there is evidence that the [at spectral. componen extends. across the millimetre and. infrared. bands to. the nearinfrared or even optical bands.," However, as argued above, there is evidence that the flat spectral component extends across the millimetre and infrared bands to the nearinfrared or even optical bands."
+ The same flat spectra component extending to the optical ραπ will have a racliative luminosity of ~71075 erg , The same flat spectral component extending to the optical V-band will have a radiative luminosity of $\sim 7 \times 10^{34}$ erg $^{-1}$.
+The ratio of observed. (radiative) power. g. to tota," The ratio of observed (radiative) power, $\eta$ , to total"
+As noted before. we do not discuss the local monopole term.,"As noted before, we do not discuss the local monopole term."
+ By subdividing supernovae into low and high redshift bins (Zehavietal.1998;Jha2007).. a significant variation in the local Hubble parameter has been detected.," By subdividing supernovae into low and high redshift bins \citep{Zehavi:1998gz,jrk07}, a significant variation in the local Hubble parameter has been detected."
+ This does not have any impact on our results (see section ?2)). and since it was discussed thoroughly in Jhaetal.(2007) we refer to that paper for further details.," This does not have any impact on our results (see section \ref{sec:window}) ), and since it was discussed thoroughly in \citet{jrk07} we refer to that paper for further details."
+ Results from galaxy velocity surveys on scales of order 4000-6000 kms! generally agree that the magnitude of the dipole is of order 300 kms! in the direction /~300°. b~ (see for instance Zaroubi(2002) and references therein).," Results from galaxy velocity surveys on scales of order 4000-6000 $\kms$ generally agree that the magnitude of the dipole is of order 300 $\kms$ in the direction $l \sim 300^\circ$, $b \sim
+20^\circ$ (see for instance \citet{zaroubi} and references therein)."
+ This result is compatible with the SN Ia dipole direction and magnitude within 26., This result is compatible with the SN Ia dipole direction and magnitude within $2\sigma$.
+ A reconstruction. of the very local velocity field (<3000kms ) was done by Tonryetal.(2000) measuring surface brightness fluctuations in 300 early type galaxies. predominantly in groups and clusters.," A reconstruction of the very local velocity field $<3000 \kms$ ) was done by \citet{tonry:2000} measuring surface brightness fluctuations in 300 early type galaxies, predominantly in groups and clusters."
+ They used an explicit flow model witha Virgo Attractor and a Great Attractor which contain the main local mass concentrations., They used an explicit flow model with a Virgo Attractor and a Great Attractor which contain the main local mass concentrations.
+ Furthermore they added dipole and quadrupole terms to account for the gravitational pull and shear from large scale structure further away., Furthermore they added dipole and quadrupole terms to account for the gravitational pull and shear from large scale structure further away.
+ They find a very low value of the dipole (~150km s) and the quadrupole polar vector (~50km sl). but it may be related to having the dipole in the same direction as the attractors. and the attractors accounting for the major part of the shear (quadrupole term) in the model.," They find a very low value of the dipole $\sim 150 \kms$ ) and the quadrupole polar vector $\sim 50
+\kms$ ), but it may be related to having the dipole in the same direction as the attractors, and the attractors accounting for the major part of the shear (quadrupole term) in the model."
+ The dipole has also been measured using velocity field reconstruction of the 2mass catalogue., The dipole has also been measured using velocity field reconstruction of the 2mass catalogue.
+ At a distance of 4000- kms the dipole direction is found to be roughly /— 2507. b~35—A07. again compatible with our result within 2c (Erdogduetal.2006a.b:Pike&Hudson 2005)..," At a distance of 4000-6000 $\kms$ the dipole direction is found to be roughly $l \sim
+250^\circ$ , $b \sim 35-40^\circ$, again compatible with our result within $2\sigma$ \citep{Erdogdu:2006nd,Erdogdu:2006qs,pike}."
+ Cluster samples like SMAC probe larger distances. and find directions which are generally compatible with the SN Ia result.," Cluster samples like SMAC probe larger distances, and find directions which are generally compatible with the SN Ia result."
+ For example Hudsonetal.(2004) find /=260°+13°. b=0°+11°.," For example \citet{hudson} find $l = 260^\circ \pm 13^\circ$, $b = 0^\circ \pm 11^\circ$."
+ However. they find an amplitude of 687+203km s'. significantly higher than our result (although again compatible at 27).," However, they find an amplitude of $687 \pm 203 \kms$ , significantly higher than our result (although again compatible at $2\sigma$ )."
+ We have analysed mock supernova surveys in order to study the number of supernovae needed to probe the large scale velocity field of the local universe. quantified in terms of the angular power spectra as a function of redshift.," We have analysed mock supernova surveys in order to study the number of supernovae needed to probe the large scale velocity field of the local universe, quantified in terms of the angular power spectra as a function of redshift."
+ We then proceeded to use the best available database of low-redshift supernovae. the JRK sample. to probe the local dipole and quadrupole of the velocity field at three different distances.," We then proceeded to use the best available database of low-redshift supernovae, the JRK sample, to probe the local dipole and quadrupole of the velocity field at three different distances."
+ The present method has several advantages over galaxy surveys., The present method has several advantages over galaxy surveys.
+ The uncertainty on each individual supernova lummosity is much smaller than the systematic uncertainties in determining galaxy lummosities so that a much smaller sample is sufficient., The uncertainty on each individual supernova luminosity is much smaller than the systematic uncertainties in determining galaxy luminosities so that a much smaller sample is sufficient.
+ We find that Finally. we note that new surveys like Pan-STARRs. SkyMapper and LSST will measure about 10.000 type la supernovae at z«0.1 per year. and if proper light curves and redshifts can be measured for even a small fraction of these events they will provide an extremely powerful tool for studying the dynamics of the local universe.," We find that Finally, we note that new surveys like Pan-STARRs, SkyMapper and LSST will measure about 10,000 type Ia supernovae at $z < 0.1$ per year, and if proper light curves and redshifts can be measured for even a small fraction of these events they will provide an extremely powerful tool for studying the dynamics of the local universe."
+ We thank the Danish Centre of Scientific Computing (DCSC) for granting the computer resources used., We thank the Danish Centre of Scientific Computing (DCSC) for granting the computer resources used.
+ TH thanks the DARK Cosmology Centre for hospitality during. the course of this work., TH thanks the DARK Cosmology Centre for hospitality during the course of this work.
+SJ is grateful for support at KIPAC and SLAC via the Panofsky Fellowship.,SJ is grateful for support at KIPAC and SLAC via the Panofsky Fellowship.
+ The Dark Cosmology Centre is funded by the DNRF., The Dark Cosmology Centre is funded by the DNRF.
+on the X-ray luminosities were determined by generating a Monte Carlo Markov Chain using the built-in functionality.,on the X-ray luminosities were determined by generating a Monte Carlo Markov Chain using the built-in functionality.
+ After the chain had converged. we calculated the luminosity (without absorption) from each set of values in the chain.," After the chain had converged, we calculated the luminosity (without absorption) from each set of values in the chain."
+ The quoted luminosity is the median luminosity from the chain. and the uncertainties were calculated using the 15.85 and 84.15 percentiles.," The quoted luminosity is the median luminosity from the chain, and the uncertainties were calculated using the 15.85 and 84.15 percentiles."
+ The rest-frame 0.5.2 keV luminosities of thermal minihaloes of Perseus core galaxies from Sun et al (2007) are also given in Table 4.., The rest-frame $0.5-2$ keV luminosities of thermal minihaloes of Perseus core galaxies from Sun et al (2007) are also given in Table \ref{tab:Luminosities}.
+ These are the total (power-law+MEKAL})) luminosities., These are the total ) luminosities.
+ Sun et al (2007) give no data for sources 6. 7. 9. L1. 12. 13 and found only two minihalo objects.," Sun et al (2007) give no data for sources 6, 7, 9, 11, 12, 13 and found only two minihalo objects."
+ We found two more., We found two more.
+ The absolute total B band magnitudes of the galaxies are also shown in Table 4.. calculated using the luminosity distance and the total B band flux density from NED.," The absolute total B band magnitudes of the galaxies are also shown in Table \ref{tab:Luminosities}, calculated using the luminosity distance and the total B band flux density from NED."
+ There are no reported B band fluxes for sources 5.6. 11. 12.13.," There are no reported B band fluxes for sources 5, 6, 11, 12, 13."
+ As the first 200 ks of the observations occurred around 2 yr before the remainder. we investigated whether the detected sources were time variable.," As the first 200 ks of the observations occurred around 2 yr before the remainder, we investigated whether the detected sources were time variable."
+ We determined the X-ray flux of the sources for the two epochs of observation and found the fluxes were the same within the uncertainties., We determined the X-ray flux of the sources for the two epochs of observation and found the fluxes were the same within the uncertainties.
+ Therefore no significant variability was found., Therefore no significant variability was found.
+ We now make a systematic investigation of the properties of the detected 13 early-type galaxies., We now make a systematic investigation of the properties of the detected 13 early-type galaxies.
+ In Fig., In Fig.
+ 5 we plot for each source the unabsorbed power-law X-ray luminosity versus the K band luminosity., \ref{fig:Kbandlumin} we plot for each source the unabsorbed power-law X-ray luminosity versus the K band luminosity.
+ La was derived from Kay measured within the 20 mag : isophote (taken from Two Micron All Sky Survey. 2MASS). using My. = 3.33 mag.," $_K$ was derived from $_{20}$ measured within the 20 mag $^{-2}$ isophote (taken from Two Micron All Sky Survey, 2MASS), using $_{K\odot}$ = 3.33 mag."
+ We also plot a straight line showing the linear relation between the X-ray luminosity of LMXB ina galaxy and the galaxy K band luminosity (from Kim Fabbiano 2004)., We also plot a straight line showing the linear relation between the X-ray luminosity of LMXB in a galaxy and the galaxy K band luminosity (from Kim Fabbiano 2004).
+ Kim Fabbiano (2004) find a range in K band luminosity of 7 to 40x10!Le. which is a similar range to our galaxies.," Kim Fabbiano (2004) find a range in K band luminosity of 7 to $40\times
+10^{10} L_{K\odot}$, which is a similar range to our galaxies."
+ All but two of our sources lie below the relation., All but two of our sources lie below the relation.
+ Therefore the galaxies appear superficially underluminous in point sources., Therefore the galaxies appear superficially underluminous in point sources.
+ To assess the effect of instrumental effects and projection of intracluster medium on the detectability of point sources. we compared from Perseus 11278 (Ly=40.55xTUA ) the optically brightest galaxy detected here. against the available observation of 7720 1," To assess the effect of instrumental effects and projection of intracluster medium on the detectability of point sources, we compared from Perseus 1278 $L_K=40.55 \times 10^{10}L_{K\odot}$ ), the optically brightest galaxy detected here, against the publicly-available observation of 720 )."
+ 7720 lies at a distance of 28 Mpe. is not in a cluster and has a observation of 40 ks length (OBSID 492).," 720 lies at a distance of 28 Mpc, is not in a cluster and has a observation of 40 ks length (OBSID 492)."
+ To account for the effect of distance on the spatial seale. we binned the image of 7720 by a factor of 3. and smoothed it by | aresec to account for the PSF at the position of 11278.," To account for the effect of distance on the spatial scale, we binned the image of 720 by a factor of 3, and smoothed it by 1 arcsec to account for the PSF at the position of 1278."
+ From this image we subtracted a flat background. then multiplied the image by a factor of 2.3 to account for the difference in exposure time and distance between the Perseus and 7720 observations.," From this image we subtracted a flat background, then multiplied the image by a factor of 2.3 to account for the difference in exposure time and distance between the Perseus and 720 observations."
+ We added a flat background component to account for the projected intracluster medium in Perseus., We added a flat background component to account for the projected intracluster medium in Perseus.
+ Finally we generated an image by making a Poisson realization of our model. to account for counting statistics.," Finally we generated an image by making a Poisson realization of our model, to account for counting statistics."
+ We show an image of 11278. 7720 before processing. and the final simulated 7720 image in Fig. 6..," We show an image of 1278, 720 before processing, and the final simulated 720 image in Fig. \ref{fig:fake}."
+ It is clear that most of the halo of point sources (LMXB) spread several aremin around 7720 become undetectable under the observing conditions of 11278., It is clear that most of the halo of point sources (LMXB) spread several arcmin around 720 become undetectable under the observing conditions of 1278.
+ The situation would become worse after the higher absorption to the Perseus cluster is included., The situation would become worse after the higher absorption to the Perseus cluster is included.
+ We only detect in 11278 a point source coincident with the nucleus and a possible LMXB to the SW., We only detect in 1278 a point source coincident with the nucleus and a possible LMXB to the SW.
+ It would be interesting to see whether this last source coincides with a globular cluster there but we have found no published information of globular clusters in our 13 galaxies., It would be interesting to see whether this last source coincides with a globular cluster there but we have found no published information of globular clusters in our 13 galaxies.
+ We conclude that the low X-ray luminosity of most of the 13 galaxies. when compared with nearby ones. is due to our inability to resolve the expected. distributed. population of point sources.," We conclude that the low X-ray luminosity of most of the 13 galaxies, when compared with nearby ones, is due to our inability to resolve the expected, spatially-distributed, population of point sources."
+ ΕΙΝ and NUV photometry of the centres of the detected Perseus cluster galaxies was taken from OConnell et al (2007)., FUV and NUV photometry of the centres of the detected Perseus cluster galaxies was taken from O'Connell et al (2007).
+ We plot FUV and NUV AB magnitudes against the unabsorbed power-law X-ray luminosity in Fig. 7.., We plot FUV and NUV AB magnitudes against the unabsorbed power-law X-ray luminosity in Fig. \ref{fig:uvbandmag}.
+ We see that there is no obvious correlation between UV and X-ray fluxes., We see that there is no obvious correlation between UV and X-ray fluxes.
+ The FUV and NUV magnitudes vary from 19.21.5 and [820.5 respectively. whereas the X-ray luminosity ranges by almost a factor of 100.," The FUV and NUV magnitudes vary from $19-21.5$ and $18-20.5$ respectively, whereas the X-ray luminosity ranges by almost a factor of 100."
+ We found only four published values of radio power for the 13 galaxies., We found only four published values of radio power for the 13 galaxies.
+ These are at 1.4 GHz data and from Miller Owen (2001) and Sijbring (1993)., These are at 1.4 GHz data and from Miller Owen (2001) and Sijbring (1993).
+ The values of the logarithm of radio power in W ! are 21.9. 21.2. 21.5 and 21.2 for CGCG 540-101. NGC 1277. NGC 1278 and NGC 1281. respectively.," The values of the logarithm of radio power in W $^{-1}$ are 21.9, 21.2, 21.5 and 21.2 for CGCG 540-101, NGC 1277, NGC 1278 and NGC 1281, respectively."
+" The black hole mass of each galaxy was calculated using the relation from Marconi Hunt (2003). relating Wey and Lyjy for the ""Group I galaxies: where Lypu is in units of La..."," The black hole mass of each galaxy was calculated using the relation from Marconi Hunt (2003), relating $M_{\rm BH}$ and $L_{K,{\rm
+ bul}}$ for the `Group 1' galaxies: where $L_{K,{\rm bul}}$ is in units of $L_{K\odot}$."
+ We show in Table 5. the total 2MASS K band luminosities (taken from NED) and the derived black hole mass., We show in Table \ref{tab:BHM} the total 2MASS K band luminosities (taken from NED) and the derived black hole mass.
+ The masses vary from 2.4x107 to L1xΙΟM..., The masses vary from $2.4\times10^7$ to $1.1\times10^9\Msun$.
+ The unabsorbed power-law X-ray luminosity was plotted against the black hole mass of the galaxies in Fig. 8.., The unabsorbed power-law X-ray luminosity was plotted against the black hole mass of the galaxies in Fig. \ref{fig:luminosity}.
+ The open diamonds show those sources with a minihalo., The open diamonds show those sources with a minihalo.
+ The distances between the galaxies and the central galaxy. NGC 1275. are also given.," The distances between the galaxies and the central galaxy, NGC 1275, are also given."
+ We see that all the sources are within 250 kpe radius from I1275., We see that all the sources are within 250 kpc radius from 1275.
+ No correlation is found between the X-ray luminosity and the black hole mass., No correlation is found between the X-ray luminosity and the black hole mass.
+ All of the bright early-type galaxies in the region studied are detected in X-rays., All of the bright early-type galaxies in the region studied are detected in X-rays.
+ The X-ray luminosities range from just below, The X-ray luminosities range from just below
+ (RATE. ~LO?—10111. (vanderKlis2000;Remillareetal.2002b).. (vancerNis2000:DiSalvo.Méndez.&vanderElishowever.Migliari.vanderIxlis.&Fender2003).. 2003).. CAbramowiczRemillardetal.2002a)..," \citep[{\it RXTE}; $\sim 10^2-10^3$ \citep{kli00,rem02}. \citep[see, however, Migliari, van der Klis, \& 
+Fender 2003]{kli00,sal03}. \citep{kli00,wij03}. \citep{abr01,abr03,rem02a}."
+The formation of molecular elouds from the diffuse atomic interstellar medium has been the subject of much interest.,The formation of molecular clouds from the diffuse atomic interstellar medium has been the subject of much interest.
+ Proposed mechanisms have been reviewed in the recent literature (Andréetal.2008:Bergin&Tafalla2007;McKeeOstriker 2007).," Proposed mechanisms have been reviewed in the recent literature \citep{ABI09, BT07, MO07}."
+. Despite this interest. most chemical models of cold dense interstellar clouds utilize the simple pseudo-time-dependent approximation. in. which the physical conditions are homogeneous and fixed. beginning from an already dense. cold. and darkened state with all H in H».," Despite this interest, most chemical models of cold dense interstellar clouds utilize the simple pseudo-time-dependent approximation, in which the physical conditions are homogeneous and fixed, beginning from an already dense, cold, and darkened state with all H in $_2$."
+ Moreover. the initial abundances for heavy elements are assumed to be atomic.," Moreover, the initial abundances for heavy elements are assumed to be atomic."
+ Yet. the high abundance of molecular hydrogen in diffuse clouds suggests that molecules can be synthesized as clouds form as well as during their existence.," Yet, the high abundance of molecular hydrogen in diffuse clouds suggests that molecules can be synthesized as clouds form as well as during their existence."
+ One approach to dense cloud formation. that of hydrodynamic shock waves. has been studied by Berginetal.(2004) (hereafter BO4). who showed that the gas-phase molecules Hs and CO are produced early in the formation of the cloud.," One approach to dense cloud formation, that of hydrodynamic shock waves, has been studied by \citet{Bergin04} (hereafter B04), who showed that the gas-phase molecules $_2$ and CO are produced early in the formation of the cloud."
+ In this paper. we revisit the approach of BO4. and include a more complex treatment of the gas-grain chemistry that evolves in tandem with the post-shock physical conditions.," In this paper, we revisit the approach of B04, and include a more complex treatment of the gas-grain chemistry that evolves in tandem with the post-shock physical conditions."
+ We focus on the initial production of major solid phase and gaseous molecules in the post-shock material as the extinction gradually increases., We focus on the initial production of major solid phase and gaseous molecules in the post-shock material as the extinction gradually increases.
+ Prior investigations have considered some aspects of the problem we discuss here., Prior investigations have considered some aspects of the problem we discuss here.
+ Early dynamie models of dense and diffuse clouds were reviewed by Williams (1988).., Early dynamic models of dense and diffuse clouds were reviewed by \citet{Wi88}. .
+ Cyclic models involving shocks were studied by Nejad&Williams (1992)., Cyclic models involving shocks were studied by \citet{NW92}.
+ Our approach is similar to the constant-pressure collapse model of PineaudesForétsetal.(1991).. although that paper did not include grain surface chemistry.," Our approach is similar to the constant-pressure collapse model of \citet{Pineau91}, although that paper did not include grain surface chemistry."
+ Such chemistry was included by Ruffle&Herbst(2001) in a pseudo-time-dependent model of the formation of mantle ices using an earlier version of our gas-grain network., Such chemistry was included by \citet{RH01III} in a pseudo-time-dependent model of the formation of mantle ices using an earlier version of our gas-grain network.
+ These authors were mainly interested in explaining the formation of CO» ice. which had been under-produced in their previous models.," These authors were mainly interested in explaining the formation of $_{2}$ ice, which had been under-produced in their previous models."
+ A pseudo-time-dependent approach with surface chemistry was also used to study the water ice distribution at various conditions throughout the Taurus dark cloud (Nguyenetal.2002)., A pseudo-time-dependent approach with surface chemistry was also used to study the water ice distribution at various conditions throughout the Taurus dark cloud \citep{Nguyen02}.
+. We have used four shock models based on the results presented in BO4; the shock speed. initial total hydrogen density.and," We have used four shock models based on the results presented in B04; the shock speed, initial total hydrogen density,and"
+lig.,Fig.
+ 2 depicts the evolution of Lagrangian radii in. the unsegereeated: clusters., \ref{fig:radlnom} depicts the evolution of Lagrangian radii in the unsegregated clusters.
+ Ehe. evolution of Lagrangian raclit is determined. by three competing processes. two-hocky relaxation which decreases the core radii before core collapse and increases all radii alter core collapse. as well as gas expulsion and mass loss by stellar evolution which increase the cluster radii.," The evolution of Lagrangian radii is determined by three competing processes, two-body relaxation which decreases the core radii before core collapse and increases all radii after core collapse, as well as gas expulsion and mass loss by stellar evolution which increase the cluster radii."
+ The relative importance of the cilferent processes depends on the initial half-mass radius and number of cluster stars., The relative importance of the different processes depends on the initial half-mass radius and number of cluster stars.
+ In the most compact clusters with ry=0.033 pc. the relaxation time is short enough that they undergo core collapse before. gas expulsion sets in.," In the most compact clusters with $r_h=0.033$ pc, the relaxation time is short enough that they undergo core collapse before gas expulsion sets in."
+ Thev later expand due to gas expulsion and mass loss from massive stars in the core and can expand by up to a factor 10 within 10 Mer., They later expand due to gas expulsion and mass loss from massive stars in the core and can expand by up to a factor 10 within 10 Myr.
+" Interestingly. the final radius reached after LO Myr of evolution is the same for clusters starting with ry,=0.033 pe and ry,=0.10 pc."," Interestingly, the final radius reached after 10 Myr of evolution is the same for clusters starting with $r_h=0.033$ pc and $r_h=0.10$ pc."
+ Hence the initial radius cannot be determined. from the one after 10 Myr or very compact clusters., Hence the initial radius cannot be determined from the one after 10 Myr for very compact clusters.
+ In more extended. clusters. the initial relaxation time is too long to allow significant mass segregation before gas expulsion sets in. so these clusters »ecome Nass segreeated and eo into core collapse only on a much longer timescale of several Myr.," In more extended clusters, the initial relaxation time is too long to allow significant mass segregation before gas expulsion sets in, so these clusters become mass segregated and go into core collapse only on a much longer timescale of several Myr."
+" This is most clearly visible for the No—101 stars; ry,=1d pe cluster which ias the longest relaxation time of all clusters studied. and or which the core is contracting after J—5 Myr."," This is most clearly visible for the $N=10^4$ stars, $r_h=1$ pc cluster which has the longest relaxation time of all clusters studied and for which the core is contracting after $T=5$ Myr."
+ More compact clusters and. clusters with lower particle numbers rave developed some degree of mass segregation so that the expansion due to mass-loss outweighs contraction due to wo-body relaxation and the clusters always expand., More compact clusters and clusters with lower particle numbers have developed some degree of mass segregation so that the expansion due to mass-loss outweighs contraction due to two-body relaxation and the clusters always expand.
+ lig., Fig.
+ 3 depicts the evolution of. projected. half-light radius of bright stars in the unsegregated clusters., \ref{fig:radbnom} depicts the evolution of projected half-light radius of bright stars in the unsegregated clusters.
+ We have calculated the Lagrangian radii [rom the positions of all stars with masses larger than 3 M. that have not vet. turned into compact remnants (mainly. black holes given the short calculation times of our runs)., We have calculated the Lagrangian radii from the positions of all stars with masses larger than 3 $_\odot$ that have not yet turned into compact remnants (mainly black holes given the short calculation times of our runs).
+ Phe mass limit of 3 M. was chosen to allow a comparison with observations since in distant star clusters stars fainter than this often cannot be observed and contribute only a small fraction of the total cluster light., The mass limit of 3 $_\odot$ was chosen to allow a comparison with observations since in distant star clusters stars fainter than this often cannot be observed and contribute only a small fraction of the total cluster light.
+ Filled triangles show the age and half-light radii of open clusters from Portegics Zwart et al. (, Filled triangles show the age and half-light radii of open clusters from Portegies Zwart et al. (
+"2010). which are clusters in the mass range 10! M.<MecLO"" M.","2010), which are clusters in the mass range $10^4$ $_\odot < M_C < 10^5$ $_\odot$."
+ Filled circles show the location of additional open clusters with masses in the range lO? M.«Ale<10! M. and which contain massive stars according το Weidner et al. (, Filled circles show the location of additional open clusters with masses in the range $10^3$ $_\odot < M_C < 10^4$ $_\odot$ and which contain massive stars according to Weidner et al. (
+2010).,2010).
+ The parameters of the clusters shown can be found in Table 2., The parameters of the clusters shown can be found in Table 2.
+" lt can be seen that Lagrangian radii of bright stars also mostly expand. except for clusters starting with half-mass radii ry,=0.33 pe for which mass segregation causes a slight decrease at later times."," It can be seen that Lagrangian radii of bright stars also mostly expand, except for clusters starting with half-mass radii $r_h=0.33$ pc for which mass segregation causes a slight decrease at later times."
+" The most compact clusters all reach radii of about rj,=0.2 pe at T—10 Myr.", The most compact clusters all reach radii of about $r_{hp}=0.2$ pc at T=10 Myr.
+ This radius therefore. represents a minimum radius for star clusters and there are indeed. no observed. clusters smaller than it., This radius therefore represents a minimum radius for star clusters and there are indeed no observed clusters smaller than it.
+ Simulations starting with hall-mass racdi ry0.1 pe lead to racii significantly smaller than those of observed clusters., Simulations starting with half-mass radii $r_h \le 0.1$ pc lead to radii significantly smaller than those of observed clusters.
+ For the unsegregated clusters discussed here the best mate-— with observations is achieved for models starting with radii in the range 0.33 pe «ry1 pc., For the unsegregated clusters discussed here the best match with observations is achieved for models starting with half-mass radii in the range $0.33$ pc $< r_h < 1$ pc.
+ We next discuss the number of stellar. collisions in the unseerceated clusters., We next discuss the number of stellar collisions in the unsegregated clusters.
+ According to. PortegicsZwart&AleAlitlan (2002).. the inspiral time of a star of mass mo»Gn is given hy where {τε is the hall-mass relaxation time of the cluster which we have taken from Spitzer(1987). and where mj is the mean stellar mass.," According to \citet{spzmm02}, the inspiral time of a star of mass $m> \langle m \rangle$ is given by where $T_{rh}$ is the half-mass relaxation time of the cluster which we have taken from \citet{s87} and where $\langle m \rangle$ is the mean stellar mass."
+ Significant. number of collisions are only to be expected. after massive stars have spiraled. into, Significant number of collisions are only to be expected after massive stars have spiraled into
+We do not test other IMFs. since a Salpeter IMF is a valic assumption for explaining most of the properties of early-type galaxies (see Renzini 2005).,"We do not test other IMFs, since a Salpeter IMF is a valid assumption for explaining most of the properties of early-type galaxies (see Renzini 2005)."
+ Furthermore. it has beer already shown by Nagashima et al. (," Furthermore, it has been already shown by Nagashima et al. ("
+2005) that a change in the IMF is not enough (in the context of hierarchical mergers) i1 reproducing the correct o/ Fe-mass relation.,2005) that a change in the IMF is not enough (in the context of hierarchical mergers) in reproducing the correct $\alpha/Fe$ -mass relation.
+ We expect that a flattening of the IMF regardless of the galaxy mass leads to a overall shift towards higher values of the predicted [a/Fe ratio. but it does not affect the slope of the «/Fe-mass relation unless one finds à good reason to make the IMF flatter as the galactic mass increase.," We expect that a flattening of the IMF regardless of the galaxy mass leads to a overall shift towards higher values of the predicted $\alpha$ /Fe] ratio, but it does not affect the slope of the $\alpha/Fe$ -mass relation unless one finds a good reason to make the IMF flatter as the galactic mass increase."
+ We will investigate this possibility 11 future work., We will investigate this possibility in future work.
+ This solution appears contrived. however. because itimplies that the single building block should know in advance its destiny m order to self-assign a suitable IMF.," This solution appears contrived, however, because it implies that the single building block should know in advance its destiny in order to self-assign a suitable IMF."
+ In Eq. (, In Eq. (
+7) € represents the efticieney of mass-loadimng during the triggering of a galactic wind by SNII explosions.,7) $\epsilon^{-1}$ represents the efficiency of mass–loading during the triggering of a galactic wind by SNII explosions.
+ Decreasing € produces more feedback. heating more cold gas. ejecting more hot gas from halos. and thus reducing the amount of gas that can potentially form stars.," Decreasing $\epsilon$ produces more feedback, heating more cold gas, ejecting more hot gas from halos, and thus reducing the amount of gas that can potentially form stars."
+ In this case we find that the predicted stellar masses are smaller than the fiducial case., In this case we find that the predicted stellar masses are smaller than the fiducial case.
+ The galaxies look slightly more a-enhaced as expected since the SF process is strongly-disfavored by the SNe explosions., The galaxies look slightly more $\alpha$ -enhanced as expected since the SF process is strongly-disfavored by the SNe explosions.
+ However. this is not a viable solution for the «/Fe-mass relation problem. since the high mass—loading also implies a very low metal content in the stars.," However, this is not a viable solution for the $\alpha/Fe$ -mass relation problem, since the high mass–loading also implies a very low metal content in the stars."
+ The predicted MMR is offset downwards by at least 0.5 dex from the observational one., The predicted MMR is offset downwards by at least 0.5 dex from the observational one.
+ On the other hand. if we switch the SN feedback off. we tend to shghtly worsen the a/Fe-mass relation. whereas the agreement for the MMR improves.," On the other hand, if we switch the SN feedback off, we tend to slightly worsen the $\alpha/Fe$ -mass relation, whereas the agreement for the MMR improves."
+ Following the above line of thoughts. we further modify GallCS by introducing SNIa contribution into Eq. 7..," Following the above line of thoughts, we further modify GalICS by introducing SNIa contribution into Eq. \ref{feed}."
+ Since a elements and Fe are still ejected at the same rate. this change has the same effects of decreasing e*..," Since $\alpha$ elements and Fe are still ejected at the same rate, this change has the same effects of decreasing $\epsilon$."
+ A 0.1 dex increase in the final |a/Fe] ratios can be achieved when a differential wind is invoked. namely if we assume that twice more Fe than O can be ejected in the hot phase due to SNIa explosions.," A 0.1 dex increase in the final $[\alpha/Fe]$ ratios can be achieved when a differential wind is invoked, namely if we assume that twice more Fe than O can be ejected in the hot phase due to SNIa explosions."
+ Again. given the nature of such a mechanism. neither the slope of the predicted a/Fe-mass relation can be steepened nor its scatter reduced.," Again, given the nature of such a mechanism, neither the slope of the predicted $\alpha/Fe$ -mass relation can be steepened nor its scatter reduced."
+ Further investigation will tell us 1f a change in the SNe feedback. namely by allowing them to quench the star formation as in monolithic models (e.g. Pipino et al..," Further investigation will tell us if a change in the SNe feedback, namely by allowing them to quench the star formation as in monolithic models (e.g. Pipino et al.,"
+ 2008b). might be the required galactic-scale source of feedback.," 2008b), might be the required galactic-scale source of feedback."
+ We identify the SF recipe as one of the prescriptions where one can improve upon., We identify the SF recipe as one of the prescriptions where one can improve upon.
+ In fact. Pipino et al. (," In fact, Pipino et al. ("
+2008b) started from the heuristic approach of PMO4. who required the SF efficiency to increase as a function of galactic mass. and showed that the a/Fe-mass relation can be explained by implementing a physically motivated value for the SF efhciency.,"2008b) started from the heuristic approach of PM04, who required the SF efficiency to increase as a function of galactic mass, and showed that the $\alpha/Fe$ -mass relation can be explained by implementing a physically motivated value for the SF efficiency."
+ To explain the higher star formation efficiency in the most massive galaxies. they appeal to massive black holes-triggered SF: a short (10°—10’ yr) super-Eddington phase can provide the accelerated triggering of associated star formation.," To explain the higher star formation efficiency in the most massive galaxies, they appeal to massive black holes-triggered SF: a short $10^6-10^7$ yr) super-Eddington phase can provide the accelerated triggering of associated star formation."
+ The SMBH grows mostly in the initial super-Eddington phase while most of the spheroid stars grow during the succeeding Eddington phase. until the SN-driven wind quenches SE.," The SMBH grows mostly in the initial super-Eddington phase while most of the spheroid stars grow during the succeeding Eddington phase, until the SN-driven wind quenches SF."
+ According to Pipino et al. (, According to Pipino et al. (
+2008b) models. the galaxy is fully assembled on a time-scale of 0.3-0.5 Gyr.,"2008b) models, the galaxy is fully assembled on a time-scale of 0.3-0.5 Gyr."
+ This time-scale is long enough. however. to allow the SMBH to complete its growth in order to reproduce the Magorrian relation.," This time-scale is long enough, however, to allow the SMBH to complete its growth in order to reproduce the Magorrian relation."
+ The fact that GalICS already turns gas into stars at the maximum possible rate during the merger-induced star-burst phase and its quite low mass/space resolution hamper us from a direct implementation of the above recipe., The fact that GalICS already turns gas into stars at the maximum possible rate during the merger-induced star-burst phase and its quite low mass/space resolution hamper us from a direct implementation of the above recipe.
+ Moreover. the fact that stars born the dise and can be transferred to the bulge of the same galaxy because of instabilities. is a possibility not taken into account in Silk (2005).," Moreover, the fact that stars born the disc and can be transferred to the bulge of the same galaxy because of instabilities, is a possibility not taken into account in Silk (2005)."
+ This scenario will be tested in the forthcoming version of GallCS., This scenario will be tested in the forthcoming version of GalICS.
+ We implemented a detailed treatment for the chemical evolution of H. He. O and Fe in GallCS. a semi-analytical model for galaxy formation which successfully reproduces basic low- and high-redshift galaxy properties.," We implemented a detailed treatment for the chemical evolution of H, He, O and Fe in GalICS, a semi-analytical model for galaxy formation which successfully reproduces basic low- and high-redshift galaxy properties."
+ The contribution of supernovae (both type Ia and II) as well as low- and intermediate-mass stars to chemical feedback are taken into account., The contribution of supernovae (both type Ia and II) as well as low- and intermediate-mass stars to chemical feedback are taken into account.
+ The model predictions are compared to the most recent observational results by Thomas et al. (, The model predictions are compared to the most recent observational results by Thomas et al. (
+2008).,2008).
+ We find that this chemically improved GallCS does not produce the observed mass- and c- [o/Fe] relations., We find that this chemically improved GalICS does not produce the observed mass- and $\sigma$ $\alpha$ /Fe] relations.
+ The slope is too shallow and scatter too large. in particular in the low and intermediate mass range.," The slope is too shallow and scatter too large, in particular in the low and intermediate mass range."
+ The model shows significant improvement at the highest masses and velocity dispersions. where the predicted [a/Fe] ratios are now marginally consistent with observed values.," The model shows significant improvement at the highest masses and velocity dispersions, where the predicted $\alpha$ /Fe] ratios are now marginally consistent with observed values."
+ Moreover. an excess of low-mass ellipticals with too high a [c/Fe] ratio is predicted.," Moreover, an excess of low-mass ellipticals with too high a $\alpha$ /Fe] ratio is predicted."
+ We show that this result comes from the implementation of AGN (plus halo) quenching of the star formation in massive haloes., We show that this result comes from the implementation of AGN (plus halo) quenching of the star formation in massive haloes.
+ A thorough exploration of the parameter space shows that the failure of reproducing the mass- and o-[a/Fe] relations can partly be attributed to the way in which star formation and feedback are currently modelled., A thorough exploration of the parameter space shows that the failure of reproducing the mass- and $\sigma$ $\alpha$ /Fe] relations can partly be attributed to the way in which star formation and feedback are currently modelled.
+" The merger process 1s responsible for a part of the scatter,", The merger process is responsible for a part of the scatter.
+ We suggest that the next generation of semi-analytical model should feature feedback (either stellar of from AGN) mechanisms linked to. sirgle galaxies and not only to the halo. especially in the low and intermediate mass range.," We suggest that the next generation of semi-analytical model should feature feedback (either stellar of from AGN) mechanisms linked to single galaxies and not only to the halo, especially in the low and intermediate mass range."
+ Furthermore. a drawback of the is the fact that the MAR cannot be fit simultaneously.," Furthermore, a drawback of the is the fact that the MMR cannot be fit simultaneously."
+ Both effects can be explained by the fact that the model ts still lacking a sort of formation for all its spheroids. which is needed in order to reproduce the a/Fe-mass relation and the MMR at the same," Both effects can be explained by the fact that the model is still lacking a sort of formation for all its spheroids, which is needed in order to reproduce the $\alpha/Fe$ -mass relation and the MMR at the same"
+ , 
+active regions (through an averaging over longitude ancl uses a grid not fine enough to resolve individual sunspots. one already introduces some drastic averaging.,"active regions through an averaging over longitude and uses a grid not fine enough to resolve individual sunspots, one already introduces some drastic averaging."
+ Introducing an a-coellicient concentrated near the surface instead of using double rings may not be such a big step after Chat., Introducing an $\alpha$ -coefficient concentrated near the surface instead of using double rings may not be such a big step after that.
+ Let us end with a comment on what we mean bv à Dabcock-Leighton model. since this term often creates some confusion.," Let us end with a comment on what we mean by a Babcock-Leighton model, since this term often creates some confusion."
+ Babcock (1961) and Leighton (1969) emphasized the surface process of poloidal field generation [rom Ulled active regions.in contrast to the usual mean field MIID where the poloidal field is supposed to be produced in (he interior region of turbulence (see. for exaniple. Choudliui 1993).," Babcock (1961) and Leighton (1969) emphasized the surface process of poloidal field generation from tilted active regions—in contrast to the usual mean field MHD where the poloidal field is supposed to be produced in the interior region of turbulence (see, for example, Choudhuri 1998)."
+ Ilence any. dvnamo model in which the poloidal field is generated in a (thin laver near (he solar surface should be called a Dabcock-Leighton model., Hence any dynamo model in which the poloidal field is generated in a thin layer near the solar surface should be called a Babcock-Leighton model.
+ Dikpati Charbonneau (1999) also use the term in this sense., Dikpati Charbonneau (1999) also use the term in this sense.
+ Durnev (1995. 1996. 1997) [ollowed Leighton (1969) more closely in incorporating the Babcock-Leighton idea (through the double ring method.," Durney (1995, 1996, 1997) followed Leighton (1969) more closely in incorporating the Babcock-Leighton idea through the double ring method."
+ Introducing a phenomenological a-coefficient concentrated near the surface is another wav of representing the Dabcock-Leighton idea., Introducing a phenomenological $\alpha$ -coefficient concentrated near the surface is another way of representing the Babcock-Leighton idea.
+ One should. however. be careful not to interpret this a-coellicient in (he wav il is interpreted in the mean field ALID.," One should, however, be careful not to interpret this $\alpha$ -coefficient in the way it is interpreted in the mean field MHD."
+ For example. the a-coefficient here is not obviously related to the average helicity of turbulence as in mean field MIID (see. for example. Choudhuri 1998. 8116.5).," For example, the $\alpha$ -coefficient here is not obviously related to the average helicity of turbulence as in mean field MHD (see, for example, Choudhuri 1998, 16.5)."
+ This coellicient merely provides a phenomenological description of the production of poloidal field from the decay of tilted active regions. which is obvious in the Formulation of the BL model by Stix (1974).," This coefficient merely provides a phenomenological description of the production of poloidal field from the decay of tilted active regions, which is obvious in the formulation of the BL model by Stix (1974)."
+ Wang Sheelev (1991) also referred to this process as an “a-elfect™ in exactly the same sense as us. even though thev never used the svmbol a in their actual ecquations!," Wang Sheeley (1991) also referred to this process as an $\alpha$ -effect” in exactly the same sense as us, even though they never used the symbol $\alpha$ in their actual equations!"
+x0.,$x>0$.
+ Wewiladoptx=Oasourstandardassum tHel as well., We will adopt $x=0$ as our standard assumption and consider the two extreme cases $x=\pm1$ as well.
+" We adopt a binary fraction fii, of as our standard assumption. Le. for every single star with mass M, there is exactly one binary system of which the primary mass is Mj and the orbital period and mass ratio are within the ranges specified above."," We adopt a binary fraction $f_{\rm bin}$ of as our standard assumption, i.e. for every single star with mass $M_1$ there is exactly one binary system of which the primary mass is $M_1$ and the orbital period and mass ratio are within the ranges specified above."
+ For comparison. ? derives a binary fraction of for O stars in clusters and associations.," For comparison, \citet{Mason+98} derives a binary fraction of for O stars in clusters and associations."
+ ? find fractions between14-80%.. whereas ? infer a fraction higher than70%.," \citet{Garcia+Mermilliod01} find fractions between, whereas \cite{Kobulnicky+Fryer07} infer a fraction higher than."
+. On the other hand. Pols et al (1991) adopted a fraction of of binaries with B-type primaries. P<IO yyr and 0.25«g< 1. Besides our standard assumption we will consider binary fractions of μμ=25% and100%.," On the other hand, Pols et al (1991) adopted a fraction of of binaries with B-type primaries, $P < 10$ yr and $0.25 < q < 1$ Besides our standard assumption we will consider binary fractions of $f_{\rm bin} = 25\%$ and."
+. We note that exploring the effect of different binary fractions partially covers the uncertainties in the assumed upper limits for the orbital pertod and mass ratio that we assumed above for the initial parameter distributions., We note that exploring the effect of different binary fractions partially covers the uncertainties in the assumed upper limits for the orbital period and mass ratio that we assumed above for the initial parameter distributions.
+ The observed rate of type Hb supernovae is expressed with respect to the rate of all types of core collapse supernova., The observed rate of type IIb supernovae is expressed with respect to the rate of all types of core collapse supernova.
+ To compute this number we must consider that a significant fraction of the companion stars end their lives as core collapse supernovae., To compute this number we must consider that a significant fraction of the companion stars end their lives as core collapse supernovae.
+ The precise fraction will depend on the initial mass ptionandconsidertherwatennhistebubiess amount of mass lost from the secondary and the fraction that is actually acereted by the secondary.," The precise fraction will depend on the initial mass ratio distribution, the amount of mass lost from the secondary and the fraction that is actually accreted by the secondary."
+ For clarity. we will just consider the two extreme cases in which all or none of the companion stars result in a core collapse supernova.," For clarity, we will just consider the two extreme cases in which all or none of the companion stars result in a core collapse supernova."
+ Under the standard assumptions described above. we find that the number of type IIb supernovae over the number of core collapse supernovae formed via stable mass transfer às predicted from our models is roughly 0.6%.. see Table 5.. about a factor five lower than the observed rate (Sect. 2)).," Under the standard assumptions described above, we find that the number of type IIb supernovae over the number of core collapse supernovae formed via stable mass transfer as predicted from our models is roughly 0.6, see Table \ref{tab:rates}, about a factor five lower than the observed rate (Sect. \ref{sec:obs}) )."
+ The rate is larger. but only by a factor two. than the rate predicted from single stars.0.3%.. under the assumption that all stars are single (see Section 4)).," The rate is larger, but only by a factor two, than the rate predicted from single stars, under the assumption that all stars are single (see Section \ref{sec:single}) )."
+" Our predicited rate is compatible with the rate found by ? for the ""stripped? supernovae.", Our predicited rate is compatible with the rate found by \cite{Podsiadlowski92} for the 'stripped' supernovae.
+ In contrast to what some previous authors have suggested (e.g.2).. we emphasize that not only for the single star channel but also for the binary channel one needs to fine-tune the initial parameters of the system.," In contrast to what some previous authors have suggested \citep[e.g.][]{Podslad93}, we emphasize that not only for the single star channel but also for the binary channel one needs to fine-tune the initial parameters of the system."
+For single stars the initial mass must be,For single stars the initial mass must be
+scale of the theory for different dSph salaxies. when fitting detailed dynamical models to the data (Lokas 2001). or that AZ/L ratios remain higher than those of stellar populations. making dark matter uecessary even uuder the MOND formulation. e... Sauchez-Saleedo Ternaudez (2007).,"scale of the theory for different dSph galaxies, when fitting detailed dynamical models to the data (Lokas 2001), or that $M/L$ ratios remain higher than those of stellar populations, making dark matter necessary even under the MOND formulation, e.g., Sanchez-Salcedo Hernandez (2007)."
+ Still. this last poiut remains somewhat controversial.," Still, this last point remains somewhat controversial."
+ Aneus (2008) sueeests that the MOND proposal will hold for local dSph galaxies. once the effects of tidal disruption are adequately incorporated.," Angus (2008) suggests that the MOND proposal will hold for local dSph galaxies, once the effects of tidal disruption are adequately incorporated."
+ Reearding corresponding theoretical developiieuts. munerous alternative theories of eravity have recently appeared (Bekeusteiu 2001. Sanders 2005. Sobouti 2007. Bruneton Esposito-Farese 2007. Zhao 2007. Árbev 2008). now mostly erouuded on geometrical extensions of eeneral relativity and field theory. which lead to laws of eravity in the Newtonian limit that iu the large-scale or low-acceleration regime reduce to the MOND prescription fitting formula.," Regarding corresponding theoretical developments, numerous alternative theories of gravity have recently appeared (Bekenstein 2004, Sanders 2005, Sobouti 2007, Bruneton Esposito-Farese 2007, Zhao 2007, Arbey 2008), now mostly grounded on geometrical extensions of general relativity and field theory, which lead to laws of gravity in the Newtonian limit that in the large-scale or low-acceleration regime reduce to the MOND prescription fitting formula."
+ The MOND theory is characterised bv (1) a low acceleration regime where dvuamics mimic the presence of dark iatter. Gi) a high acceleration regime where Newtoulan gravity is recovered. and (11) a somewhat ill-defined transition region.," The MOND theory is characterised by (i) a low acceleration regime where dynamics mimic the presence of dark matter, (ii) a high acceleration regime where Newtonian gravity is recovered, and (iii) a somewhat ill-defined transition region."
+ Variants of MOND are defined by the form of the interpolation fiction. pe). which mediates the transition between the two limit reeiues. see EqAl).," Variants of MOND are defined by the form of the interpolation function, $\mu(x)$, which mediates the transition between the two limit regimes, see Eq.(1)."
+ We show how the MOND prescription with a suitable interpolation function. such as he one used by Dekeusteiu (200) and used by Famaev Binney (2005). cau he written as the addition of both low acceleration nuit MOND and Newtonian contributions to the acceleration ou all scales.," We show how the MOND prescription with a suitable interpolation function, such as the one used by Bekenstein (2004) and used by Famaey Binney (2005), can be written as the addition of both low acceleration limit MOND and Newtonian contributions to the acceleration on all scales."
+ The lavee differences in scales andl magnitudes of the two acceleration terms eusure tha the addition does not spoil the good match with observations on galactic scales., The large differences in scales and magnitudes of the two acceleration terms ensure that the addition does not spoil the good match with observations on galactic scales.
+ We are trying to look for a suitable MOND interpolation fiction for the local dSph ealaxies. but not necessarily one that works on all scales.," We are trying to look for a suitable MOND interpolation function for the local dSph galaxies, but not necessarily one that works on all scales."
+ This. as it is the τοσο of the local dSph galaxies. the one which las proven most difficult for MOND. aud the one where the high-quality receut observations now available allow for a relatively clean test.," This, as it is the regime of the local dSph galaxies, the one which has proven most difficult for MOND, and the one where the high-quality recent observations now available allow for a relatively clean test."
+ We find that this interpolation fiction not ouly vields acceptable AL/L values for all dSph ealaxics without the uecd for any dark iatter in all cases; but also provides a natural explanation for all the scalines seen iu these Calactic satellites.," We find that this interpolation function not only yields acceptable $M/L$ values for all dSph galaxies without the need for any dark matter in all cases, but also provides a natural explanation for all the scalings seen in these Galactic satellites."
+ An interesting correlation between the AL/L£ ratios inferred and the ages of the stellar populations in local dSpli galaxies is found. aud naturally accounted for ee rinodified gravity schemes where stars alone account for —«je eravitational force.," An interesting correlation between the M / L ratios inferred and the ages of the stellar populations in local dSph galaxies is found, and naturally accounted for in modified gravity schemes where stars alone account for the gravitational force."
+ ILlowever. onu solu system scales. Saunders (2006) has shown this particular interpolation function to be oeiconipatible with lianits on the variations of IWeplers constant.," However, on solar system scales, Sanders (2006) has shown this particular interpolation function to be incompatible with limits on the variations of Kepler's constant."
+ Thus. the eood accordance we fiud with observations for dSplvs must be understood as evidence oe1i favour of the particular interpolation function we test. at the weak feld linüt relevant for these systems.," Thus, the good accordance we find with observations for dSph's must be understood as evidence in favour of the particular interpolation function we test, at the weak field limit relevant for these systems."
+ What a full interpolation. function or MOND should be is a more complicated issue., What a full interpolation function for MOND should be is a more complicated issue.
+ One. however. which has to he addressed bearing iu mind the results preseuted here.," One, however, which has to be addressed bearing in mind the results presented here."
+ Section (2) eives a brief suinarv of the interpolatio- function first explored by Dekeusteim (2001) aud givey. equilibrium isothermal configurations., Section (2) gives a brief summary of the interpolation function first explored by Bekenstein (2004) and gives equilibrium isothermal configurations.
+1o These are the- used m section (3) to nodel the loc: dSphl galaxies arca obtain ML ratios. which are fouud to be consistent wit[um those tvpical of stellar populatious.ρα at the ages of tlic different dSphl svstems.," These are then used in section (3) to model the local dSph galaxies and obtain $M/L$ ratios, which are found to be consistent with those typical of stellar populations, at the ages of the different dSph systems."
+ Section (1) iuccles an exploratio- of the most conspicuous scaling relations for local dSp[um ealaxies. finding that Dekeusteiu's interpolation fuuctio- naturally accounts for all features found.," Section (4) includes an exploration of the most conspicuous scaling relations for local dSph galaxies, finding that Bekenstein's interpolation function naturally accounts for all features found."
+ Our couclusious ALC ΚΙΛΑΊπος iu section (5)., Our conclusions are summarised in section (5).
+ Iu terms of the acceleration g felt by a test particle. the MOND proposal is (see ee. Mileroii 2002) where gy is the acceleration assigned bv Newtouiau eravitv. ay an acceleration scale of the theory. aud the interpolating function porr=g/day) is an unspecified function that reduces to unity for large values of its argunient. aud to ifs argument. for simall values of it.," In terms of the acceleration $\boldsymbol{g}$ felt by a test particle, the MOND proposal is (see e.g. Milgrom 2002) where $ \boldsymbol{g}_{N}$ is the acceleration assigned by Newtonian gravity, $a_{0}$ an acceleration scale of the theory, and the interpolating function $\mu(x:=g/a_0)$ is an unspecified function that reduces to unity for large values of its argument, and to its argument, for small values of it."
+ In this wav. Newtomlan eravity is recovered for large values of g. and in the deep MOND regine one obtains Strictly speaking. MOND is a theory. defined ouly at the linut values for the acceleration.," In this way, Newtonian gravity is recovered for large values of $g$, and in the deep MOND regime one obtains Strictly speaking, MOND is a theory defined only at the limit values for the acceleration."
+ The value of the constant αρ las been reasonably determined by calibrating the deep MOND regiae through observatious of rotation velocity curves of large spiral galaxies (e.g. Sanclers MeGanegh 2002)., The value of the constant $a_{0}$ has been reasonably determined by calibrating the deep MOND regime through observations of rotation velocity curves of large spiral galaxies (e.g. Sanders McGaugh 2002).
+ ITowever. the details of the rausition region between the two regimes (i.c. the function j) have woven harder to establish reliably.," However, the details of the transition region between the two regimes (i.e. the function $\mu$ ) have proven harder to establish reliably."
+ This is in part natural. asit isnot a scalar parameter. but a functional dependence hat oue is looking for here.," This is in part natural, as it is not a scalar parameter, but a functional dependence that one is looking for here."
+" A wumber of variants for this ""unction iive been proposed (e.c. Famaey Binney 2005. Zhao 2007. Sanders Noordenueer 2007). but apparently suitable dvuamical svstems in the ransitiou region are hard to come by."," A number of variants for this function have been proposed (e.g. Famaey Binney 2005, Zhao 2007, Sanders Noordermeer 2007), but apparently suitable dynamical systems in the transition region are hard to come by."
+ Astroplivsically. to first approximation one tends to fine either svstenis where 10 dark matter is needed. (stellar aud. plauetary svstenis. elobular clusters. vertical dvuaunics of disk galaxies. ealactic bulges. and elliptical galaxies). or svstenis where dark iatter is nminsswelv doimunaut (dSph Galactic satellites. rotation curves of spiral galaxies. ealaxyv groups and clusters. and cosmological observations}.," Astrophysically, to first approximation one tends to find either systems where no dark matter is needed (stellar and planetary systems, globular clusters, vertical dynamics of disk galaxies, galactic bulges, and elliptical galaxies), or systems where dark matter is massively dominant (dSph Galactic satellites, rotation curves of spiral galaxies, galaxy groups and clusters, and cosmological observations)."
+ Also. the details of any eimpirically inferred transition region. pile).," Also, the details of any empirically inferred transition region, $\mu(x)$ ,"
+"large regions devoid of absorption, even if on average the effects are quite mild.","large regions devoid of absorption, even if on average the effects are quite mild."
+ We also varied the spectral index ns but found small differences to the reference case., We also varied the spectral index $n_s$ but found small differences to the reference case.
+ As can be appreciated from Figure the statistical errors on the observed void fraction are somewhatBl. large and thereby distinguishing between different cosmological model is still very difficult and none of the models presented here could be convincingly ruled out.," As can be appreciated from Figure \ref{fig2}, the statistical errors on the observed void fraction are somewhat large and thereby distinguishing between different cosmological model is still very difficult and none of the models presented here could be convincingly ruled out."
+" Several other parameters could have an impact on the void distribution, in particular the two parameters that are poorly constrained and that describe the thermal state of the IGM: Το and +."," Several other parameters could have an impact on the void distribution, in particular the two parameters that are poorly constrained and that describe the thermal state of the IGM: $T_0$ and $\gamma$."
+ We investigate the effect of a different value of To in Figure , We investigate the effect of a different value of $T_0$ in Figure \ref{fig5}.
+"A colder (hotter) IGM will result in a higher (lower) density[g. of neutral hydrogen and thereby in more absorptions, which, in turn, will decrease (increase) the fraction of large voids."," A colder (hotter) IGM will result in a higher (lower) density of neutral hydrogen and thereby in more absorptions, which, in turn, will decrease (increase) the fraction of large voids."
+" In Figure [6], we also overplot the effect of changing y from the reference value (y= 1.6) to y=1.1 (represented by the dotted line in the Figure), which is a better fit to the data at this redshift (?).."," In Figure \ref{fig5}, we also overplot the effect of changing $\gamma$ from the reference value $\gamma=1.6$ ) to $\gamma=1.1$ (represented by the dotted line in the Figure), which is a better fit to the data at this redshift \citep{bolton07}."
+" A flatter equation of state will not only make underdense regions hotter, but it will also make the slightly overdense regions colder."," A flatter equation of state will not only make underdense regions hotter, but it will also make the slightly overdense regions colder."
+" The overall effect is a slight decrease in the fraction of flux voids, since cold regions embed on average more neutral hydrogen than hot ones."," The overall effect is a slight decrease in the fraction of flux voids, since cold regions embed on average more neutral hydrogen than hot ones."
+ This result shows that by selecting mean-flux-level regions we are sampling regions around mean density (for which the power-law 'equation of state’ could be a fit to the actual , This result shows that by selecting mean–flux–level regions we are sampling regions around mean density (for which the power–law 'equation of state' could be a poor fit to the actual distribution).
+Note that the pooreffect of changing y was distribution).addressed by running a new simulation that results in a low at z=2 (using a modified self-consistentlyversion of the code ykindly provided by J. Bolton) and not by rescaling the gas-temperature relation a posteriori as is sometimes done., Note that the effect of changing $\gamma$ was addressed by running a new simulation that self–consistently results in a low $\gamma$ at $z=2$ (using a modified version of the code kindly provided by J. Bolton) and not by rescaling the gas-temperature relation a posteriori as is sometimes done.
+" Another important input parameter is the overall amplitude of the UVB, parametrized by [_12, for which we have some observational constraints (?).."," Another important input parameter is the overall amplitude of the UVB, parametrized by $\Gamma_{-12}$, for which we have some observational constraints \citep{bolt05}."
+" However, this quantity is degenerate with the value of Teg, so changing [_12 will result in the same change as changing Teg."," However, this quantity is degenerate with the value of $\tau_{\rm eff}$, so changing $\Gamma_{-12}$ will result in the same change as changing $\tau_{\rm eff}$."
+" Having identified voids at z~2, using the mean flux as a threshold, we are now interested in linking the flux properties to the underlying physical properties of the gas density field."," Having identified voids at $z\sim 2$, using the mean flux as a threshold, we are now interested in linking the flux properties to the underlying physical properties of the gas density field."
+" The main problem we have to deal with is that the flux is observed in redshift space, while we want to recover gas properties in real space."," The main problem we have to deal with is that the flux is observed in redshift space, while we want to recover gas properties in real space."
+" The role of peculiar velocities in determining the properties of the transmitted flux has been investigated by several authors (e.g. ?)), however for the purpose of the analysis performed here it produces a shift of the absorption feature from the real space overdensity to the redshift space observed flux."," The role of peculiar velocities in determining the properties of the transmitted flux has been investigated by several authors (e.g. \cite{weinberg}) ), however for the purpose of the analysis performed here it produces a shift of the absorption feature from the real space overdensity to the redshift space observed flux."
+ This increases the scatter since redshift space pixels are not related in a simple manner to the underlying real space density field., This increases the scatter since redshift space pixels are not related in a simple manner to the underlying real space density field.
+" Occasionally, peculiar velocities could be of the order of a few hundred km/s, so the correspondence is non-trivial (see for example ?))."," Occasionally, peculiar velocities could be of the order of a few hundred km/s, so the correspondence is non-trivial (see for example \cite{rauchetal}) )."
+" However, we can rely on hydrodynamical simulations in order to infer, at least on a statistical basis, the physical of the "," However, we can rely on hydrodynamical simulations in order to infer, at least on a statistical basis, the physical properties of the gas."
+"In Π, we show the mean values and propertiesscat"," In Figure \ref{fig6}, we show the mean values and scatter (r.m.s."
+ter gas. in Figurethe real gas-overdensity and temperature (r.m.s. values)corresponding to the spaceflux voids of the previous sections.," values) in the real space gas–overdensity and temperature fields, corresponding to the flux voids of the previous sections."
+" We fields,show the fields both without (diamond", We show the fields both without (diamond
+back aud forth [rom ihe central position τμ=2018.5. Note that this iuplies that we are left with the first and last 23.5-pixel-wide vertical strips where tle discontinulty is repeated.,"back and forth from the central position $x_{\rm raw} =
+2048.5$ Note that this implies that we are left with the first and last 23.5-pixel-wide vertical strips where the discontinuity is repeated."
+ The horizontal feature aoug the ray axis consists of a sinele 2-pixe-wide cdiscontinuity. centered at Uraw=911 pixels [or chip[1] and at yy=1110 pixels for chip[2].," The horizontal feature along the $y_{\rm raw}$ axis consists of a single 2-pixel-wide discontinuity, centered at $y_{\rm raw} = 911$ pixels for chip[1] and at $y_{\rm raw}
+= 1140$ pixels for chip[2]."
+ These discontinuities are located syimnunetrically with respect to the gap between the two chips (horizontal solid red liues in Fig. 1))., These discontinuities are located symmetrically with respect to the gap between the two chips (horizontal solid red lines in Fig. \ref{fig:res_before}) ).
+ A careful look at Fig., A careful look at Fig.
+ 2 reveals hints of finer cdiscontiuuities on |oth astrometry (top panel) ancl flat field (bottom patel). but at much lower amplituces (a few thorsanidtlis of a pixe and ~0.15). however it’s harc| to assess tlieir significauce.," \ref{fig:fetta} reveals hints of finer discontinuities on both astrometry (top panel) and flat field (bottom panel), but at much lower amplitudes (a few thousandths of a pixel and $\sim0.1$ ), however it's hard to assess their significance."
+" Qur aun here was to find a simple correction of the GD folloyviug the basic priuciple: ""we see a systematic error. aud we empirically find a correction for it."," Our aim here was to find a simple correction of the GD following the basic principle: “we see a systematic error, and we empirically find a correction for it”."
+ We tlerelore decided to keep a polynomial of the hird order. as we did in Paper L to remove most o “the GD (clewu to the ~1 uias level). aud theu to use a single ook-up table (one for each chip and filter) to correct. all the stnaller-scale positiollal systematic errors.," We therefore decided to keep a polynomial of the third order, as we did in Paper I, to remove most of the GD (down to the $\sim1$ mas level), and then to use a single look-up table (one for each chip and filter) to correct all the smaller-scale positional systematic errors."
+ This approach was able to )'ovide accuracies down to 0.008 pixel (0:3) mas). as we will show in the next section.," This approach was able to provide accuracies down to 0.008 pixel $\sim0.3$ mas), as we will show in the next section."
+ We set up the |ook-up table as follows., We set up the look-up table as follows.
+ First we defined τρις boiuia‘ies alongside the lithographic feature cliscontiuuiies (red solid lines iu Fig. 1))., First we defined $x_{\rm raw}$ boundaries alongside the lithographic feature discontinuities (red solid lines in Fig. \ref{fig:res_before}) ).
+ We sbdivided eacl of the 675-:May pixelwide regious into six 112.5-pixel-wide sub-regious. for a total o £36 such subdivisions plus two regions at the left aud righte edgese of each chip.," We subdivided each of the $x_{\rm raw}$ -pixel-wide regions into six 112.5-pixel-wide sub-regions, for a total of 36 such subdivisions plus two 23.5-pixel-wide regions at the left and right edges of each chip."
+ To 1jalutaln a similar s:uupliug aloug the tray axis. we defined 15 sub-divisiois.," To maintain a similar sampling along the $y_{\rm raw}$ axis, we defined 18 sub-divisions."
+ As the horizontal component of the lithographic feature clivices tlie two chips iu wo parts of 91] aud 1110 pixels tall. we nace & subclivisious for the tray pixel oue (113.575 pixels each) aud 10 for the other (111 pixels each).," As the horizontal component of the lithographic feature divides the two chips in two parts of 911 and 1140 pixels tall, we made 8 subdivisions for the $y_{\rm raw}$ -pixel one (113.875 pixels each) and 10 for the other (114 pixels each)."
+ At the enc we produced au array of 36x15 alios-square cells. plus two 23.5-pixel-wice srips at the short edges of eacl chip. each made of 15 rectaugular cells (grey dashed lines in Figs. 1)).," At the end we produced an array of $36\times18$ almost-square cells, plus two 23.5-pixel-wide strips at the short edges of each chip, each made of 18 rectangular cells (grey dashed lines in Figs. \ref{fig:res_before}) )."
+ Cell dimeusOLS were ultimately dictated by the necessity to have enough erid-poitts to finely sample the GD :uid to have an aclequate ituber of stars within each cell to robustly measure the value of the eric poiuts in the look-ip tabe., Cell dimensions were ultimately dictated by the necessity to have enough grid-points to finely sample the GD and to have an adequate number of stars within each cell to robustly measure the value of the grid points in the look-up table.
+ We always had more thau 30+) stars in each cell to coustrain the vaue of the table. even for tLe filters with the fewest iunber of well-exposed stars.," We always had more than 30 stars in each cell to constrain the value of the table, even for the filters with the fewest number of well-exposed stars."
+ Typically we hac well over 100 stars per ce, Typically we had well over 100 stars per cell.
+ Figure κ)+) slows an exainple of the geometry. adopted for the ook-up tables., Figure \ref{fig:scheme} shows an example of the geometry adopted for the look-up tables.
+ Thick solid lines mark detector ecges. dashed lines identify lithographic ciscontiuuiies. while dotted lines hiehllight cell borders.," Thick solid lines mark detector edges, dashed lines identify lithographic discontinuities, while dotted lines highlight cell borders."
+ We used stars within each cell to compute 30-clippecL inedian positional residuals dr aud dy. which are assigne to the correspoucing grid point (open cicles).," We used stars within each cell to compute $3\sigma$ -clipped median positional residuals $\overline{\delta x}$ and $\overline{\delta y}$, which are assigned to the corresponding grid point (open circles)."
+ When a cell adjoius eitler detector edges or lithograghic discoutinuities. the eric poiut is dis]jxaced to the edge ofthe cell. as shown.," When a cell adjoins either detector edges or lithographic discontinuities, the grid point is displaced to the edge of the cell, as shown."
+ We use a semiciree to indicate when a grid poiut correspoids to a ciscoutinuity., We use a semicircle to indicate when a grid point corresponds to a discontinuity.
+ For any giveu location on the chip. the look-up table correction is given by a bi-linear itterpolatiou," For any given location on the chip, the look-up table correction is given by a bi-linear interpolation"
+Jdy instead of 1dy.,$J-K$ instead of $H-K$.
+ Mere 247 objects have photometry in the two surveys. 208 without effects of saturation (J211. A 10.5)," Here 247 objects have photometry in the two surveys, 208 without effects of saturation $J>11$, $K>10.5$ )."
+ Again. there are 2 objects with diy> ο... he same ones as in the left panel.," Again, there are 2 objects with $dK>0.5$ mag, the same ones as in the left panel."
+ From the samples shown in the figure. 63 have mid-infrared excess due to the presence of disks based on he IRAC colours. prm]- jon]20.45 minag. pon]- pan]οἱ15 Πας).," From the samples shown in the figure, 63 have mid-infrared excess due to the presence of disks based on the IRAC colours $\mu m$ $\mu m$ $>0.45$ mag, $\mu m$ $\mu m$ $>0.15$ mag)."
+ Amongthe8.0 mislofthe 3bighlyr euriableobject tho sources 1. 1," Among them is 1 of the 2 highly variable objects, the one listed in Table \ref{var}. ."
+" hesecondone(o 577. DI=|32701'53,7"". SIMDAD name CI* 1C 348 LRL 172). las µπ]- pon]0-44 mmagand pimi]- jm] =0.01 mmag."," The second one $\alpha = 03^h44^m22.57^s$ , $\delta = +32^o01'53.7''$, SIMBAD name Cl* IC 348 LRL 72), has $\mu m$ $\mu m$ $=0.44$ mag and $\mu m$ $\mu m$ $=0.01$ mag."
+ Lalsohasstronge rccssal2lqnmn3.6 with µηπ]- —LS8mmeag.iwhichismorethanmosltobjectie snthe[scts grmESS," It also has strong excess at $\,\mu m$ with $\mu m$ $\mu m$ $=4.8$ mag, which is more than most object with disks in IC348."
+" Ann errors.then πμ.twiceThe Hes uscofpeatheobjectsBay\iu material,which bulitmighthe SNgun. a so-called.24i] transition disk."," Thus, this source clearly has circumstellar material, but it might have an inner hole in the disk and therefore little excess at $\,\mu m$, a so-called transition disk."
+" Indeed. the object was previously classified as a ""transition disk candidate (7) and is an interesting target. for further monitoring."," Indeed, the object was previously classified as a 'transition disk candidate' \citep{2010ApJ...708.1107M} and is an interesting target for further monitoring."
+ Since the variability for this source is mostly seen in the Ix-band. it is likely to originate in the clisk.," Since the variability for this source is mostly seen in the K-band, it is likely to originate in the disk."
+ The fraction hiehly variable sources. as defined in this paper. it thus 1/208ol for the total sample (1%)) and 1/63 for objects with disks (254).," The fraction of highly variable sources, as defined in this paper, it thus 1/208 for the total sample ) and 1/63 for objects with disks )."
+ 1n 1€348 the plots indicate that there are little problems with the band correction. most likelv because the extinction is on average lower than in p-Oph and the ONC.," In IC348 the plots indicate that there are little problems with the band correction, most likely because the extinction is on average lower than in $\rho$ -Oph and the ONC."
+ Therefore it ds feasible to provide quantitative limits on the variability of the general population., Therefore it is feasible to provide quantitative limits on the variability of the general population.
+ The median dilference between PALASS and. UIXIDSS photometry is 0.05. 0.04. mmag in J. HL. and [x for the entire sample ancl 0.07. 0.06. 0.05 for 8. IW wich dinlli," The median difference between 2MASS and UKIDSS photometry is 0.05, 0.04, mag in J, H, and K for the entire sample and 0.07, 0.06, 0.05 for the sources with disks."
+s These numbers are still alfected. by the observational errors. which are dominated by the PALASS uncertainties.," These numbers are still affected by the observational errors, which are dominated by the 2MASS uncertainties."
+ The UIXIDSS errors are typically by a factor of 10 lower Hhdisthan 3l5., The UKIDSS errors are typically by a factor of 10 lower than the 2MASS errors.
+ [roctions De. 1C 0 Hus. as COHE as CLEOES (i.c. variable on a 2e level) in J}. 1L. Ix are. for all objects and56%...52%... lor the ones with disks.," The fractions of objects for which the offsets are more then twice as large as the 2MASS errors (i.e. variable on a $\sigma$ level) in J, H, K are, for all objects and, for the ones with disks."
+ The typical lo binomial uncertainties in these fractions are for all objects and for the stars with disks., The typical $1\sigma$ binomial uncertainties in these fractions are for all objects and for the stars with disks.
+ For these variable sources. the median dillerence between 2NLASS and UIXIDSS after subtracting the2ALASS error is 0.05. 0.06. mamas for all objects and 0.07. 0.09. 006 for those with clisks.," For these variable sources, the median difference between 2MASS and UKIDSS after subtracting the2MASS error is 0.05, 0.06, mag for all objects and 0.07, 0.09, 0.06 for those with disks."
+this object is very high (even higher than the lastest L dwarl elu 1). the large difference in the observed degree of polarization Irom this object and DENIS-P J0255-4700 cannot be explainecdl as the mean particle size in the atmosphere of L clwarls of same spectral tvpe should not differ much.,"this object is very high (even higher than the fastest L dwarf Kelu 1), the large difference in the observed degree of polarization from this object and DENIS-P J0255-4700 cannot be explained as the mean particle size in the atmosphere of L dwarfs of same spectral type should not differ much."
+ Further theoretical investigation on the distribution and location of condensates in the coolest L clwarls are needed before modeling their degree of polarization., Further theoretical investigation on the distribution and location of condensates in the coolest L dwarfs are needed before modeling their degree of polarization.
+ ]lence. in the present work we discuss linear polarization of L dwarls of spectral (vpe ranging from LO to L6.," Hence, in the present work we discuss linear polarization of L dwarfs of spectral type ranging from L0 to L6."
+ Figure 5 shows the degree of linear polarization al A=0.850 jam for objects with different spectral types but fixed surface gravity LO? em 7., Figure 5 shows the degree of linear polarization at $\lambda=0.850$ $\mu$ m for objects with different spectral types but fixed surface gravity $^{5}$ cm $^{-2}$.
+ It is found that the observed degree of polarization of several L clwarls can be fitted if (he mean diameter of grain is taken as 1.4 jum and with the polvtropic index n—1.0., It is found that the observed degree of polarization of several L dwarfs can be fitted if the mean diameter of grain is taken as 1.4 $\mu m$ and with the polytropic index n=1.0.
+ The figure shows that the change in polarization with the spectral (ype is not linear but overall the degree of polarization decreases as one moves [rom L1 to LG., The figure shows that the change in polarization with the spectral type is not linear but overall the degree of polarization decreases as one moves from L1 to L6.
+ The degree of polarization decreases substantially [or objects hotter than L1 because condensation is not favored at such high temperature., The degree of polarization decreases substantially for objects hotter than L1 because condensation is not favored at such high temperature.
+ The location of the cloud. base shilts to deeper region of the atmosphere as the effective temperature decreases and hence should cause an increase in (he polarization (figure 3)., The location of the cloud base shifts to deeper region of the atmosphere as the effective temperature decreases and hence should cause an increase in the polarization (figure 3).
+ ILowever. the total amount ol condensing material is conserved.," However, the total amount of condensing material is conserved."
+ Therefore. as the scale height of the cloud laver becomes smaller. (he particle number density becomes higher vielding into higher polarization.," Therefore, as the scale height of the cloud layer becomes smaller, the particle number density becomes higher yielding into higher polarization."
+ Devond Ll object. the temperature becomes too high to favor condensation aud the polarization [alls rapidly to zero.," Beyond L1 object, the temperature becomes too high to favor condensation and the polarization falls rapidly to zero."
+ ZapateroOsorioetal(2004) detected the polarization in the Johnson R and E-band fillers centered on 0.641 and 0.850 pam respectively with the passband of these filters as 0.158 and 0.15 yam., \citet{oso04} detected the polarization in the Johnson R and I-band filters centered on 0.641 and 0.850 $\mu m$ respectively with the passband of these filters as 0.158 and 0.15 $\mu m$.
+ Menard οἱ al. (, Menard et al. (
+2002) detected the polarization in the Bessel I-handl filler on 0.768 yom with the passband as 0.138 jam.,2002) detected the polarization in the Bessel I-band filter on 0.768 $\mu m$ with the passband as 0.138 $\mu m$.
+ We have calculated the degree of polarization al the central wavelengths., We have calculated the degree of polarization at the central wavelengths.
+ The change in degree of polarization over (he spread of wavelength should bv ancl large be absorbed in the error bars., The change in degree of polarization over the spread of wavelength should by and large be absorbed in the error bars.
+ Both 2421ÀSS JL7O7+43 and 2MASS J1412--16 belong to the spectral tvpe of L0.5., Both 2MASS J1707+43 and 2MASS J1412+16 belong to the spectral type of L0.5.
+ PATASS J14124-16 having projected rotational velocity 16.4 + shows degree of polarization 7250.19 while 2M1ASS JL707+43 shows degree of polarization 0.2340.06 but its rotational velocily is not known., 2MASS J1412+16 having projected rotational velocity 16.4 $^{-1}$ shows degree of polarization $\pm$ 0.19 while 2MASS J1707+43 shows degree of polarization $\pm$ 0.06 but its rotational velocity is not known.
+ Figure 5 shows that the observed. data of these two objects can well be explained if the rotational velocity of 2421ASS J17072-43 is much less than 15 !., Figure 5 shows that the observed data of these two objects can well be explained if the rotational velocity of 2MASS J1707+43 is much less than 15 $^{-1}$.
+ The model assumes (he mean diameter of grain is 1.4. yam and the polvtropie index n—1.0., The model assumes the mean diameter of grain is 1.4 $\mu m$ and the polytropic index n=1.0.
+" The same model can explain the observed polarization of 2\TASS J1507-16 (L5 with projected rotational velocity 27.2 1) if its effective temperature is slightly higher then that given by the spectral (wpe - T,;; polynomial formula.", The same model can explain the observed polarization of 2MASS J1507-16 (L5 with projected rotational velocity 27.2 $^{-1}$ ) if its effective temperature is slightly higher then that given by the spectral type - $_{eff}$ polynomial formula.
+ We put it as L5.5 to show that if the effeclive temperature of this object corresponds (o spectral type L5.5 in the polvnomial, We put it as L5.5 to show that if the effective temperature of this object corresponds to spectral type L5.5 in the polynomial
+"where Rea, is the Schwarzschild radius of the central black hole). which produce very redshifted. line emission (with observed. photon energies of 4keV. or less).","where ${\rm R}_{\rm Sch}$ is the Schwarzschild radius of the central black hole), which produce very redshifted line emission (with observed photon energies of $\sim 4\keV$ or less)."
+ ‘Thus. whilst GR enhancement. effects. might be important for understanding the broadest line known. they cannot be relevant to strong iron lines from typical objects.," Thus, whilst GR enhancement effects might be important for understanding the broadest line known, they cannot be relevant to strong iron lines from typical objects."
+ 1n thisfeffer we note that any relative motion between the X-ray source and the aceretion disk. will also allect (and usually enhance) the EW of the iron line through the ellects of special relativistic (SIt) aberration and. Doppler shifts.," In this, we note that any relative motion between the X-ray source and the accretion disk will also affect (and usually enhance) the EW of the iron line through the effects of special relativistic (SR) aberration and Doppler shifts."
+ Such relative cisk/corona motion will naturally occur if the accretion disk and the corona are not rigidly coupled together., Such relative disk/corona motion will naturally occur if the accretion disk and the corona are not rigidly coupled together.
+ For example. sone authors treat the disk-corona as an independent. slim accretion clisk.," For example, some authors treat the disk-corona as an independent, slim accretion disk."
+ Due to the subsequent sub-Weplerian motion of the corona. the disk and. corona will be in relative motion at any given radius.," Due to the subsequent sub-Keplerian motion of the corona, the disk and corona will be in relative motion at any given radius."
+ However. the corona may well be tightly. coupled to the aceretion disk bv magnetic fields (which will force the disk and corona to accrete together).," However, the corona may well be tightly coupled to the accretion disk by magnetic fields (which will force the disk and corona to accrete together)."
+ Even in this circumstance. the SR. effects discussed. here may be important.," Even in this circumstance, the SR effects discussed here may be important."
+ Field. Rogers (1993: hereafter. E1493). have. argued. that magnetic: instabilities and reconnection events in a disk corona could. produce shock waves and/or the streaming of relativistic particles along the magnetic field. lines., Field Rogers (1993; hereafter FR93) have argued that magnetic instabilities and reconnection events in a disk corona could produce shock waves and/or the streaming of relativistic particles along the magnetic field lines.
+ Thus. the plasma which is instantancously responsible for the N-rav emission. might well be in bulk motion relative to the disk.," Thus, the plasma which is instantaneously responsible for the X-ray emission might well be in bulk motion relative to the disk."
+ Such arguments eain qualitative support by drawing an analogy with the solar corona and solar Lares., Such arguments gain qualitative support by drawing an analogy with the solar corona and solar flares.
+ For convenience. we set the speed of light to unity. e=1. throughout this work.," For convenience, we set the speed of light to unity, $c=1$, throughout this work."
+ Most of the SR effects relevant to the EW of the iron line can be studied within a scenario in which the iron line originates from a stationary slab of cold matter., Most of the SR effects relevant to the EW of the iron line can be studied within a scenario in which the iron line originates from a stationary slab of cold matter.
+ Suppose we have a semi-infinite slab of cold gas filling the half-space z«0., Suppose we have a semi-infinite slab of cold gas filling the half-space $z<0$.
+ Let this cold slab be stationary with respect to a distant observer who views it at an inclination 7 with respect to the upward normal to the slab., Let this cold slab be stationary with respect to a distant observer who views it at an inclination $i$ with respect to the upward normal to the slab.
+ Furthermore. suppose this cold. slab is illuminated by an X-ray source which is at some distance / above the face of the slab.," Furthermore, suppose this cold slab is illuminated by an X-ray source which is at some distance $h$ above the face of the slab."
+ We make the following assumptions about the X-ray source: We further assume that the source is moving at velocity e relative to the slab., We make the following assumptions about the X-ray source: We further assume that the source is moving at velocity $v$ relative to the slab.
+ We define a to be the angle this velocity. vector makes with the downward. normal (thus a=0 corresponds to motion directly towards the slab)., We define $\alpha$ to be the angle this velocity vector makes with the downward normal (thus $\alpha=0$ corresponds to motion directly towards the slab).
+ We also define 3 to be the azimuthal direction of the source in the slab plane relative to some reference line on the slab., We also define $\beta$ to be the azimuthal direction of the source in the slab plane relative to some reference line on the slab.
+ We choose the projection of the observers line of sight to be this reference line., We choose the projection of the observers line of sight to be this reference line.
+ Finally. we construct a standard: 2-cl polar co-ordinate svstem (r6) on the slab taking the point that is (instantaneous) below the source to be the origin. and using same reference direction used to define 23.," Finally, we construct a standard 2-d polar co-ordinate system $(r,\phi)$ on the slab taking the point that is (instantaneous) below the source to be the origin, and using same reference direction used to define $\beta$."
+ From simple vector. geometry. we can calculate two other important angles.," From simple vector geometry, we can calculate two other important angles."
+ First. for a given point on the slab (1r.0). the angle that the source velocity makes with the line joining the source with that point. €. is given by Secondly. the angle between the observers line of sight and the source velocity. 5. is The question we wish to address is this: how does the relative slab-source motion influence the strength of the iron Iluorescence line that will result from the illumination.," First, for a given point on the slab $(r,\phi)$, the angle that the source velocity makes with the line joining the source with that point, $\theta$, is given by Secondly, the angle between the observers line of sight and the source velocity, $\eta$, is The question we wish to address is this: how does the relative slab-source motion influence the strength of the iron fluorescence line that will result from the illumination."
+ There will be two relevant elfects., There will be two relevant effects.
+ First. the Doppler shifts and aberration of the source emission. will influence both the number ancl average direction of primary photons that fall above the iron photoelectric threshold.," First, the Doppler shifts and aberration of the source emission will influence both the number and average direction of primary photons that fall above the iron photoelectric threshold."
+ This will change the line photon emission rate as compared. with the static source case., This will change the line photon emission rate as compared with the static source case.
+ Secondly. the fact that the primary emission. sullers relativistic aberration whereas the line emission does not will allect the observed ratio of these two emissions and. hence. the EW of the iron line.," Secondly, the fact that the primary emission suffers relativistic aberration whereas the line emission does not will affect the observed ratio of these two emissions and, hence, the EW of the iron line."
+ We will now discuss these two elfects in turn., We will now discuss these two effects in turn.
+ There are two distinct ingredients involved. in determining he absolute emission rate of Duorescent iron line photons., There are two distinct ingredients involved in determining the absolute emission rate of fluorescent iron line photons.
+ The most obvious one is the number of illuminating photons with energv £cfas where £u=τ]Κον ds the whotoelectric threshold for neutral iron.," The most obvious one is the number of illuminating photons with energy $E>E_{\rm th}$, where $E_{\rm th}=7.1\keV$ is the photoelectric threshold for neutral iron."
+ Only these incident οποίος can eject one of the WKeshell electrons. from. iron and. thus. initiate the radiative cascade within the atom hat results in a Ke line photon being emitted.," Only these incident photons can eject one of the K-shell electrons from iron and, thus, initiate the radiative cascade within the atom that results in a $\alpha$ line photon being emitted."
+ For a ixed illuminating spectrum. the iron line strength is simply xoportional to the normalization of that spectrum.," For a fixed illuminating spectrum, the iron line strength is simply proportional to the normalization of that spectrum."
+ The second ingredient in determining the absolute iron ine emission is the geometry of the illumination. (Dasko 1978: George Fabian. 1991)., The second ingredient in determining the absolute iron line emission is the geometry of the illumination (Basko 1978; George Fabian 1991).
+ Consider. normally incident οποίος with energv ££Li., Consider normally incident photons with energy $E>E_{\rm th}$.
+ On average. such photons raverse a distance corresponding to unity optical depth xior to being photoclectrically absorbed.," On average, such photons traverse a distance corresponding to unity optical depth prior to being photoelectrically absorbed."
+ The resulting iron luorescence photons have to travel through at least the same optical depth of material in order to escape the slab., The resulting iron fluorescence photons have to travel through at least the same optical depth of material in order to escape the slab.
+ Some raction of these iron line photons will be absorbed in this »ocess (either by Ix-shell photoionization of low-Z metals or L-shell photoionization of iron)., Some fraction of these iron line photons will be absorbed in this process (either by K-shell photoionization of $Z$ metals or L-shell photoionization of iron).
+ Now. consider photons that," Now, consider photons that"
+increase the size of the rregion sample. and to extend their gradient study to other atoms such as neon and sulfur in the near future. for à more complete comparison with the PNe. with the goal of obtaining firmer observational constraints for the chemical evolutionary nodels.,"increase the size of the region sample, and to extend their gradient study to other atoms such as neon and sulfur in the near future, for a more complete comparison with the PNe, with the goal of obtaining firmer observational constraints for the chemical evolutionary models."
+For the infrared selected galaxies in the [RAS 2 Jv sample. the mean value of q is z2.35 (Yun.ReddyaudCoudou2001).,"For the infrared selected galaxies in the $IRAS$ 2 Jy sample, the mean value of $q$ is $\simeq2.35$ \citep{yun01}."
+. Caven the FIR huuinosity aud radio flux deusitvof 112032|1707. we estimate yo=l.?7.," Given the FIR luminosity and radio flux density of 12032+1707, we estimate $q=1.77$."
+ This value appears low compared with the mean of 2.35 and could iuplv the preseuce of an ACN., This value appears low compared with the mean of 2.35 and could imply the presence of an AGN.
+ This appears contradictory to the non-detection of a compact radio AGN core., This appears contradictory to the non-detection of a compact radio AGN core.
+ However. radio-excess objects iu the FRAS 2 Jy sample are defined as those objects having a radio bhunünositv that is greater than 5 times luger that predicted bv the radio-FIR correlation.," However, radio-excess objects in the $IRAS$ 2 Jy sample are defined as those objects having a radio luminosity that is greater than 5 times larger that predicted by the radio-FIR correlation."
+ Equivaleutlv. this means objects for which qzl1.61.," Equivalently, this means objects for which $\leq$ 1.64."
+ Furthermore. 112032|1707 has a very large infrared Iunuinositv (log(Lig/L..)= 12.57). placing the ealaxy at the high tail of the ULIRC huuiositv distribution (NimaudSaunders 1998).," Furthermore, 12032+1707 has a very large infrared luminosity $(L_{\rm IR}/L_{\odot})=12.57$ ), placing the galaxy at the high tail of the ULIRG luminosity distribution \citep{kim98}."
+. For the most hunuinous sources in the ZRAS 2 Jv sample. scattering from the radio-FIR correlation is significantly larger than for the weaker sources with weaker FIR flux densities.," For the most luminous sources in the $IRAS$ 2 Jy sample, scattering from the radio-FIR correlation is significantly larger than for the weaker sources with weaker FIR flux densities."
+ Therefore. the relatively low q-value for our source is not sufficient to indicate the presence of an AGN component.," Therefore, the relatively low $q$ -value for our source is not sufficient to indicate the presence of an AGN component."
+ Other evidence that 112032|1707. is starburst-powered comes froi its relatively cool IR. color. Fosjau/foujnn=0.18 (simandSaunders1998).," Other evidence that 12032+1707 is starburst-powered comes from its relatively cool IR color, $F_{\rm
+ 25\mu m}/F_{\rm 60\mu m}=0.18$ \citep{kim98}."
+" It is oulv for warn colors (Τοπ)ζουμ>0.25) that au AGN conrponeut is required to explain the resulting dust teniperatures,", It is only for warm colors $F_{\rm 25\mu m}/F_{\rm 60\mu m}>0.25$ ) that an AGN component is required to explain the resulting dust temperatures.
+ The zero-iutensitv line width of the OIL maser chussion in 112032|1707 is around 2000i. which makes this source the lost of oue of the broadest OII imiegeduaser lines.," The zero-intensity line width of the OH maser emission in 12032+1707 is around 2000, which makes this source the host of one of the broadest OH megamaser lines."
+ As ientioued in the introduction. this velocity range could be due to a combination of effects such as disk rotation aud gas flows.," As mentioned in the introduction, this velocity range could be due to a combination of effects such as disk rotation and gas flows."
+ Due to the limited SNR in our data. dt ds difficult. to clotermine unainbieuouxslv the cause of the masers in 112032|1707.," Due to the limited SNR in our data, it is difficult to determine unambiguously the cause of the masers in 12032+1707."
+ Below we discuss a few possible scenarios that could agree with our observatious., Below we discuss a few possible scenarios that could agree with our observations.
+ It has been suggested that very broad ΟΠ mceamascr lines might be a combination of maser chussion occuring in of the merging unclei. as is the case for the OT meceamaser prototype 2220 (Diamondetal.1989).," It has been suggested that very broad OH megamaser lines might be a combination of maser emission occurring in of the merging nuclei, as is the case for the OH megamaser prototype 220 \citep{diamond89}."
+. We can rule out that maser cussion in 112032|L707 is the combined emissiou from thesouthern aud northern nuclei. since it is clear that all maser Cluission occurs only in the northeri uucleus refolispatial) ).," We can rule out that maser emission in 12032+1707 is the combined emission from thesouthern and northern nuclei, since it is clear that all maser emission occurs only in the northern nucleus \\ref{ohspatial}) )."
+ We can not however exclude the possibility that the northern nucleus itself is made up of two more Closcly interacting ealaxies., We can not however exclude the possibility that the northern nucleus itself is made up of two more closely interacting galaxies.
+ In 2220 the uuclei have a projected separation of only 300 pe. which would correspond to an angular resolution of 86 mas at the distance of 112032|1707. much smaller than the near-infrared secius limit of reported for the NIR image (Nimctal. 2002)..," In 220 the nuclei have a projected separation of only 300 pc, which would correspond to an angular resolution of 86 mas at the distance of 12032+1707, much smaller than the near-infrared seeing limit of reported for the NIR image \citep{kim02}. ."
+ This issue requires further iuvestigation, This issue requires further investigation
+"of ρω is reasonably well approximated by pj,x/.τι in agreement with Eq. 34..","of $p_{max}$ is reasonably well approximated by $p_{max} \propto t^{-1/2}$, in agreement with Eq. \ref{eq:emax1t},"
+" since £,() is roughly. constant (see solid curve on the right panels of Figs. 2--", since $\xi_c(t)$ is roughly constant (see solid curve on the right panels of Figs. \ref{fig:B1n01}- \ref{fig:B5n01}-
+3--4 and diseussion below)., \ref{fig:B1n003} and discussion below).
+ A crucial ingredient in ealeulating the maximum energy αἱ a given (ime is the strength of the magnetic field., A crucial ingredient in calculating the maximum energy at a given time is the strength of the magnetic field.
+ The magnetic field intensitv in (he downstream plasma is plotted in (he right panels (dot-dashed line) for the three cases considered here., The magnetic field intensity in the downstream plasma is plotted in the right panels (dot-dashed line) for the three cases considered here.
+ The twpical values are between afew—LOG ab late times and ~30—005€ al the beginning of the Sedov phase.," The typical values are between $\sim {\rm a\,few -}10\mu G$ at late times and $\sim 30-100\mu G$ at the beginning of the Sedov phase."
+ After the first few thousand vears. the dependence on time is not far [rom 97sx|HU. as would result from Eq. 32.. ," After the first few thousand years, the dependence on time is not far from $\delta B_2 \propto t^{-3/10}$, as would result from Eq. \ref{eq:delB1t}, ,"
+"using the additional information on (he approximate constancy οἱ £.(1) and fey, (dashecl line in (he left panels of Figs.", using the additional information on the approximate constancy of $\xi_c(t)$ and $R_{sub}$ (dashed line in the left panels of Figs.
+ 2. to 4))., \ref{fig:B1n01} to \ref{fig:B1n003}) ).
+ Despite these resemblances. the situation to which the plots refer is more complicated (han that described by Eq. 32..," Despite these resemblances, the situation to which the plots refer is more complicated than that described by Eq. \ref{eq:delB1t},"
+" where a number of effects have been neglected. first among these the presence of a precursor which evolves with (ime (/2,4; is changing as can be seen from the dot-dashed curve on the left of Figs."," where a number of effects have been neglected, first among these the presence of a precursor which evolves with time $R_{tot}$ is changing as can be seen from the dot-dashed curve on the left of Figs."
+ 2. to 4)) and the time varving adiabatic compression it entails., \ref{fig:B1n01} to \ref{fig:B1n003}) ) and the time varying adiabatic compression it entails.
+ The right panels also show the acceleration efficiency. (solid line) ancl normalized escape flix (dashed lines)., The right panels also show the acceleration efficiency (solid line) and normalized escape flux (dashed lines).
+ One should notice that even when the acceleration efficiency is verv high. of order ~50—GO%. the escape energy [Iux never exceeds ~3056.," One should notice that even when the acceleration efficiency is very high, of order $\sim 50-60\%$, the escape energy flux never exceeds $\sim 30\%$."
+ As discussed above. this latter quantity should be the one that is more directly related to Che cosmic rav energetics in the Galaxy. at least at the highest energies. while the former is more relevant [or thegeneration of secondary radiation due to cosmic ray interactions in the acceleration region.," As discussed above, this latter quantity should be the one that is more directly related to the cosmic ray energetics in the Galaxy, at least at the highest energies, while the former is more relevant for thegeneration of secondary radiation due to cosmic ray interactions in the acceleration region."
+corrected [OU] luminosities SFR=ὃν10ο [M PWW dp ,"corrected ] luminosities ${\rm SFR} = 5 \times 10^{-34} L_{\rm [O\,II]}$ $_{\odot}$ $^{-1}$ $^{-1}$ ]."
+Gallagher et al. (, Gallagher et al. (
+1989) for blue ealaxies aud Cowie ct al. (,1989) for blue galaxies and Cowie et al. (
+1997) for rapidly star-forming ealaxies found a considerably lower couversion factor of 1.0«10MEAT AIMS L/W. For high redshitdt ealaxies. Thompson Djorgovski (1991) used a factor of 1.2«10.7! for the [On| liue ο,"1997) for rapidly star-forming galaxies found a considerably lower conversion factor of $1.0 \times 10^{-34} L_{\rm [O\,II]}$ $_{\odot}$ $^{-1}$ /W. For high redshifdt galaxies, Thompson Djorgovski (1991) used a factor of $1.2 \times 10^{-34}$ for the ] line emission."
+ For galaxies seen in the light of Πα at ;=0.25. the 1]] line fluxes are measured with the CADIS filter at Ae= l65nnu. (," For galaxies seen in the light of $\alpha$ at $z=0.25$, the ] line fluxes are measured with the CADIS filter at $\lambda_{\rm c}=465$ nm. ("
+see top of Table 1).,see top of Table 1).
+ In 77 the line fluxes are plotted for the galaxies with (9/Vyou>1.," In 7 the line fluxes are plotted for the galaxies with $(S/N)_{\rm [O\,II]}>1$."
+ The average flux ratio results oein FOH]/F(IIn)~0.9. in good agreement with the values cited above.," The average flux ratio results in $F{\rm [O\,II]}/F{\rm (H\alpha)} \sim 0.9$, in good agreement with the values cited above."
+ This ratio is roughly in the mean of the ratios derived by IKeunicutt (1998) and Sullivan ct al. (, This ratio is roughly in the mean of the ratios derived by Kennicutt (1998) and Sullivan et al. (
+2000).,2000).
+ Although there is a difference of roughly (στ between the epochs for 2=1.20 aud :=125. the line ratio may uot have changed significauth.," Although there is a difference of roughly Gyr between the epochs for $z=1.20$ and $z=0.25$, the line ratio may not have changed significantly."
+ Based on the CFRS ealaxy sample with redshifts in the range 0.5<2< 1.0. Carollo Lilly (2001) found a remarkable similarity of metallicities to that for local galaxies.," Based on the CFRS galaxy sample with redshifts in the range $0.5<z<1.0$ , Carollo Lilly (2001) found a remarkable similarity of metallicities to that for local galaxies."
+ Also. Pettini et al. (," Also, Pettini et al. ("
+1998) and Teplitz et al. (,1998) and Teplitz et al. (
+2000) found the line spectra at high redshift to be very similar to those observed iu local enussion line galaxies.,2000) found the line spectra at high redshift to be very similar to those observed in local emission line galaxies.
+ The calibration used iu the present study for the [Ou] emission is therefore Rosa-Couzalez et al. (, The calibration used in the present study for the ] emission is therefore Rosa-Gonzalez et al. (
+2002) derived. SER. estimators using the Calzetti extinction curve aud accounting for the underlying stellar Balmer absorption.,2002) derived SFR estimators using the Calzetti extinction curve and accounting for the underlying stellar Balmer absorption.
+ Applving their calibration to our non-extiuction corrected 11] galaxy sample we would cud up with about two times higher star forming rates., Applying their calibration to our non-extinction corrected ] galaxy sample we would end up with about two times higher star forming rates.
+bv scattering. until it escapes through the central funnels parallel to the disc axis.,"by scattering, until it escapes through the central funnels parallel to the disc axis."
+ The collimation results from the fact that the funnel is tall and thin. and has scattering walls.," The collimation results from the fact that the funnel is tall and thin, and has scattering walls."
+ To apply the formalism leading to (2)) we identify the blackbody radius Ras R~Row=ry. with p=τη/Ε. where r1.," To apply the formalism leading to \ref{puch}) ) we identify the blackbody radius $R$ as $R \sim R_{\rm sph} = rR_s$, with $r = 27\dot
+m\bar r/4$, where $\bar r \sim 1$."
+" The blackbocy luminosity emitted. by the disc within oj, is the intrinsic luminosity £.", The blackbody luminosity emitted by the disc within $R_{\rm sph}$ is the intrinsic luminosity $L$.
+ This dilfuses inwarels and is collimated bv the funnels., This diffuses inwards and is collimated by the funnels.
+ For the beaming factor 6 we consider a simple ονΠοσα. funnel around the central disc axis., For the beaming factor $b$ we consider a simple cylindrical funnel around the central disc axis.
+ ΙΓ the typical cvlindrical radius of the funnel is fy. and its height is flyXRoy. the halfangle over which radiation escapes is 6csin5Rof/llyRof/dfy.," If the typical cylindrical radius of the funnel is $R_0$, and its height is $H_0 \gg R_0$, the half–angle over which radiation escapes is $\theta \simeq \sin^{-1} R_0/H_0 \simeq
+R_0/H_0$."
+ Vhen the beaming fraction b is simply the total fractional area of the two funnels. ic. b~(1cos@)—RE/2HG.," Then the beaming fraction $b$ is simply the total fractional area of the two funnels, i.e. $b \sim (1 - \cos\theta) \sim
+R_0^2/2H_0^2$."
+ Close to the disc plane. the structure of the central region of the disc wind (and thus the funnel radius £5) is independent of the value of IEddington ratio m lat Large A. since all such disces have the same central accretion rate behaviour A202)cΑμ IR.," Close to the disc plane, the structure of the central region of the disc wind (and thus the funnel radius $R_0$ ) is independent of the value of Eddington ratio $\dot m >1$ at large $R$ , since all such discs have the same central accretion rate behaviour $\dot M(R) \simeq \dot M_E(R/R_s)$ ."
+ We expect that Ry~AR with Ac1. as Py sets the lengthscales in this region.," We expect that $R_0 \sim \lambda
+R_s$ with $\lambda > 1$, as $R_s$ sets the lengthscales in this region."
+ However the funnel height ff sensitive to imi; or equivalentIy fé= rusas ab points [ar [rom the clise plane the wind How pattern is set by i. which is equivalent to saving that the large.scale How. pattern is self.similar aud scaled by ον ," However the funnel height $H_0$ sensitive to $\dot
+m$, or equivalently $R_{\rm sph} = rR_s$, as at points far from the disc plane the wind flow pattern is set by $\dot m$, which is equivalent to saying that the large–scale flow pattern is self–similar and scaled by $r$ ."
+In. particular this requires fy~pisaxXpor. where p«1. so that finally where «re stands for the climensionless combination We see the the observational requirement (5)) implies Afpo~58. so that the funnel height is only a few. percent of Roy. de. quoLow10 ?.," In particular this requires $H_0 \sim
+\mu R_{\rm sph} \propto \mu r$, where $\mu < 1$, so that finally where $x$ stands for the dimensionless combination We see the the observational requirement \ref{num}) ) implies $\lambda/\mu \sim 58$, so that the funnel height is only a few percent of $R_{\rm sph}$, i.e. $\mu \sim {\rm few} \times 10^{-2}$ ."
+ We get finally TheΤΙ reasoning& of ththis paragrapherapl assumes1 that mi is larelarge enough that the two scales Ry ane Loy are very dilferent., We get finally The reasoning of this paragraph assumes that $\dot m$ is large enough that the two scales $R_0$ and $R_{\rm sph}$ are very different.
+ The scaling of b with mis clearly more complex for smaller m., The scaling of $b$ with $\dot m$ is clearly more complex for smaller $\dot m$.
+ In particular. unless {10Ry. which requires. &5.eE U7. one would formally get b>L.," In particular, unless $H_0 > R_0$, which requires $\dot m
+> 8.5x^{-1/2}$ , one would formally get $b >1$."
+ The last Section dealt. with those ULXs (the majority) for which soft components are seen. and obey the LugxLt relation.," The last Section dealt with those ULXs (the majority) for which soft components are seen, and obey the $L_{\rm soft} \propto T^{-4}$ relation."
+ We noted above that WP show that the remaining (nonpowerlaw) ULXs follow opposed correlations £L—17 or the mediumenergy Xravs., We noted above that KP show that the remaining (non–power–law) ULXs follow opposed correlations $L \sim T^4$ for the medium–energy X–rays.
+ Here the normalization differs for cach individual svstem., Here the normalization differs for each individual system.
+ The luminosities of he two groups dilfer sharply: the powerlawsoft.excess systems have. inferred. luminosities permanently above 351075 while the nonpowerlaw systems are »vmanentlv below 107 t+.," The luminosities of the two groups differ sharply: the power��law–soft–excess systems have inferred luminosities permanently above $3\times 10^{39}$ $^{-1}$, while the non–power–law systems are permanently below $10^{39}$ $^{-1}$."
+ Le seenmis likely that these inler systems correspond to ΙΙ aceretion on o black holes with masses <1OAL..," It seems likely that these fainter systems correspond to –Eddington accretion on to black holes with masses $\ga
+10\msun$."
+ It is clear that. such systems must exist. and that there is no reason to expect the collimation leading to the opposite ZugxZ3 behaviour of the bright ULXs.," It is clear that such systems must exist, and that there is no reason to expect the collimation leading to the opposite $L_{\rm
+ soft} \propto T^{-4}$ behaviour of the bright ULXs."
+ The normalizations of the ἐν~T correlations are then fixed bv the system inclinations and the inner disc radii., The normalizations of the $L \sim T^4$ correlations are then fixed by the system inclinations and the inner disc radii.
+ Phe latter are indeed. of order a few Schwarzschild radii for black holes of ~10M.., The latter are indeed of order a few Schwarzschild radii for black holes of $\sim 10\msun$.
+ We can now check whether the inferred behaviour of the beamine factor leads to sensible parameters for observed bright ULXs., We can now check whether the inferred behaviour of the beaming factor leads to sensible parameters for observed bright ULXs.
+ Although the relation (8)) for 6 was derived using the inferred blackbody disc emission. its gcomoetrical nature and the fact that electron scattering is independent of photon οποίον makes it. probable that it holds for all forms of ULX luminosity. ancl indeed even in ULXs where no blackbocly disc component has been identified. provided only that these correspond to superLEcldington accretion.," Although the relation \ref{beam}) ) for $b$ was derived using the inferred blackbody disc emission, its geometrical nature and the fact that electron scattering is independent of photon energy makes it probable that it holds for all forms of ULX luminosity, and indeed even in ULXs where no blackbody disc component has been identified, provided only that these correspond to super–Eddington accretion."
+ In particular we can use the bxii7 scaling in considering mediumenergv Xravs. which are generally assumed. to carey most of thebolometric luminosity ofa ULX.," In particular we can use the $b \propto \dot m^{-2}$ scaling in considering medium–energy X–rays, which are generally assumed to carry most of thebolometric luminosity ofa ULX."
+ The effect of beaming is tocause an observer to infer a spherical luminosity(cf Shakura Sunvaev. 1973: Degelman et aL.," The effect of beaming is tocause an observer to infer a spherical luminosity(cf Shakura Sunyaev, 1973; Begelman et al.,"
+ 2006)., 2006).
+ Eliminating 6 using (6)) (or (5)) gives We can reexpress this as, Eliminating $b$ using \ref{b}) ) (or \ref{num}) ) gives We can re–express this as
+identical with that of CTG.,identical with that of CTG.
+ For the line at 5105À.. their adopted foggf value is smaller than that of CTC by 0.06 dex.," For the line at 5105, their adopted $log~gf$ value is smaller than that of CTG by 0.06 dex."
+ Thus from this line alone. Chey would derive a larger abundance of Cu than we derive in (he present investigation: this amounts to 0.03 ces. averaging over the two lines.," Thus from this line alone, they would derive a larger abundance of Cu than we derive in the present investigation; this amounts to 0.03 dex, averaging over the two lines."
+ Although we do not apply this “correction” in the comparison ol results. we note that 5105 is analvzed in this paper only for a few e@iants in Md. NGC 362. NGC 288. MIS and all stars in MI5. (," Although we do not apply this “correction” in the comparison of results, we note that 5105 is analyzed in this paper only for a few giants in M4, NGC 362, NGC 288, M13 and all stars in M15. ("
+5) Adopted models.,5) Adopted models.
+ Mishenina et al..," Mishenina et al.,"
+ following Mishenina Ixovtvukh. adopt ]urucz models with convective overshoot turned on (Ixurucz1993).," following Mishenina Kovtyukh, adopt Kurucz models with convective overshoot turned on \citep{KurCD1993}."
+. The differences induced in [Fe/H] among metal-poor giants by switching from MARCS to Iurucz models has been explored by Kraft Ivans (2002)., The differences induced in [Fe/H] among metal-poor giants by switching from MARCS to Kurucz models has been explored by Kraft Ivans (2002).
+ Among intermecliately metal-poor stars (-L.7« |Fe/HM| <-— 1.2). the effect on [Fe/I]] (from Fe I) is negligible (< 0.01 dex). but rises to 0.07 to 0.10 dex among relatively metal-rich ancl metal-poor giants.," Among intermediately metal-poor stars $-1.7~<$ [Fe/H] $<~-1.2$ ), the effect on [Fe/H] (from Fe I) is negligible $\lesssim$ 0.01 dex), but rises to 0.07 to 0.10 dex among relatively metal-rich and metal-poor giants."
+ Fortunately. of the 21 halo field giants considered here. 11 have also been studied by CTG (Ixraftetal.1992:Shetrone1996).," Fortunately, of the 21 halo field giants considered here, 11 have also been studied by CTG \citep{Kraft1992, Shetrone1996}."
+. We can then empirically determine the offsets in Try. logg. and |Fe/1] which must be applied to bring the field and cluster giants onto the same abundance svstem.," We can then empirically determine the offsets in $_{eff}$, $log~g$, and [Fe/H] which must be applied to bring the field and cluster giants onto the same abundance system."
+" In the sense CTG minus Mishenina Ixovtvukh. we find the following mean dillerences: AT),=εν(089.5) Alogg=+0.28(o0.20): A /IJ=+0.05(o=0.10)."," In the sense CTG minus Mishenina Kovtyukh, we find the following mean differences: $\Delta$ $_{eff}=+47K ~(\sigma=89.5K)$; $\Delta~log~g=+0.28 ~(\sigma=0.20)$; $\Delta$ $=+0.05 ~(\sigma=0.10)$."
+" The offsets in T,ει and |Fe/I] are small (GE significant). but the offset in /ogg is rather large."," The offsets in $_{eff}$ and [Fe/H] are small (if significant), but the offset in $log~g$ is rather large."
+ Fortunately. the values of [Cu/Fe| are not very sensitive to variations in /ogg.," Fortunately, the values of [Cu/Fe] are not very sensitive to variations in $log~g$."
+ In addition. and more importantly. the offset is just about what one would expect owing to the dilferences in the loggf. values for Fe H.," In addition, and more importantly, the offset is just about what one would expect owing to the differences in the $log~gf$ values for Fe II."
+ If |Fe/1I] from Fe IL is brought into agreement with Ρο from Fe I. then a shift of 0.11 in the Fe II {ουgf values corresponds to a shift in fogg of 0.20 dex (vans et al.," If [Fe/H] from Fe II is brought into agreement with [Fe/H] from Fe I, then a shift of 0.11 in the Fe II $log~gf$ values corresponds to a shift in $log~g$ of 0.20 dex (Ivans et al."
+ 2001. Table 3). close to thevalue of 0.28 cited above.," 2001, Table 3), close to thevalue of 0.28 cited above."
+" The |Fe/H] offset cited above is also explicable in terms of the offset in T, rj.", The [Fe/H] offset cited above is also explicable in terms of the offset in $_{eff}$ .
+ Again. according," Again, according"
+"higli-: Lyra :~2 to.~1 hieher-: :G 2=6). selectedτ~6 |i/'-baud |[27729/25,70: NTU B89,τι”? /-dropout +6. /-—dropouts 7/-dropout /-dropouts (QuOQ.h)=(0.3.0.7.0.7)unless 2” ","$z$ $\alpha$ $z\sim2$ $z\sim4$ $z$ $z\sim6$ $z\geq6$ $z\sim6$ $i'$ $+27\degr 29'25.''9$ $\sim876$ $^2$ $i'$ $z\simeq 6$ $i'$ $i'$ $i'$ $(\Omega_m, \Omega_{\Lambda}, h)=(0.3,0.7,0.7)$ $2''$ "
+"The contribution to cosmological mass density. would. be OQ,h=Mimi,/04 ev where the sum is over neutrino types.",The contribution to cosmological mass density would be $\Omega_\nu h^2 = \sum{m_{\nu_i}}/94$ ev where the sum is over neutrino types.
+ The neutrino oscillation experiments do not. provide information on the actual masses of neutrino species but on the square of the mass dillerences., The neutrino oscillation experiments do not provide information on the actual masses of neutrino species but on the square of the mass differences.
+ These are small. such that. the largest mass cdillerence. suggested: bv the atmospheric oscillations. is Am=0.05 ev.," These are small, such that, the largest mass difference, suggested by the atmospheric oscillations, is $\Delta m\approx
+0.05$ ev."
+" Wom, then the three active neutrino types would have no significant cosmological mass density (=107) and could not contribute to the mass budget. of any bound astronomical object.", If then the three active neutrino types would have no significant cosmological mass density $\approx 10^{-3}$ ) and could not contribute to the mass budget of any bound astronomical object.
+ But another possibility is that iAm and that the masses of alb three active types are nearly equal., But another possibility is that $m_\nu >>\Delta m$ and that the masses of all three active types are nearly equal.
+ In this case an upper limit to the masses is provided. by an experimental limit on the mass of the electron neutrino. bc. 2.2 ev at cL. (," In this case an upper limit to the masses is provided by an experimental limit on the mass of the electron neutrino, i.e. 2.2 ev at c.l. ("
+Groom et al.,Groom et al.
+ 2000) [fit were the case that the electron neutrino mass were near 2 ev. then neutrinos would constitute a significantfraction of the cosmic density (ο20.13 for h=0.7).," 2000) If it were the case that the electron neutrino mass were near 2 ev, then neutrinos would constitute a significantfraction of the cosmic density $\Omega_\nu \approx 0.13$ for h=0.7)."
+ llowever. neutrinos of this mass could not. contribute to the mass budget of galaxies.," However, neutrinos of this mass could not contribute to the mass budget of galaxies."
+ This follows [from a classic argument by ‘Tremaine Gunn (1979) based. upon conservation of the phase space density of the neutrino Iuid., This follows from a classic argument by Tremaine Gunn (1979) based upon conservation of the phase space density of the neutrino fluid.
+ Relativistic neutrinos are created with à maximum. phase space density of (2rf)° per type including anti-neutrinos his is a factor of two less than the absolute limit implied by uantum. mechanical degeneracy)., Relativistic neutrinos are created with a maximum phase space density of $(2\pi\hbar)^{-3}$ per type including anti-neutrinos (this is a factor of two less than the absolute limit implied by quantum mechanical degeneracy).
+ In. subsequent evolution of the neutrino Duid involving gravitational instability and collapse. the final phase space density cannot. exceed. this value.," In subsequent evolution of the neutrino fluid involving gravitational instability and collapse, the final phase space density cannot exceed this value."
+" This provides a relation between the final density of neutrino dark matter and the velocity. dispersion of the system: with three types Equivalently. for virialized svstems. this may be written as a relation between the cllective core racdus of a dark halo and its velocity dispersion: this is. roughly. Clearly for objects with the required velocity dispersion of galaxy halos (25m/s<a,«200 km/s). neutrinos in the mass range of one to two ον could not. possible cluster on sub-Mpe scales."," This provides a relation between the final density of neutrino dark matter and the velocity dispersion of the system; with three types Equivalently, for virialized systems, this may be written as a relation between the effective core radius of a dark halo and its velocity dispersion; this is, roughly, Clearly for objects with the required velocity dispersion of galaxy halos $25\,{\rm km/s}<\sigma_r<200\, {\rm km/s}$ ), neutrinos in the mass range of one to two ev could not possible cluster on sub-Mpc scales."
+ Llowever. it obviously would be possible for such neutrinos to contribute to the mass budget of large clusters.," However, it obviously would be possible for such neutrinos to contribute to the mass budget of large clusters."
+ SS shows the the fitted matter density of the dark cores of the clusters of galaxies listed. in ‘Table 1 vs. the eas temperature of the clusters., 8 shows the the fitted matter density of the dark cores of the clusters of galaxies listed in Table 1 vs. the gas temperature of the clusters.
+" The solid. points are the clusters with significant mass discrepancies (M4/M,7 1.5) and the open points the clusters with lower discrepancies (AliΛΙ< 1.5).", The solid points are the clusters with significant mass discrepancies $M_d/M_g>1.5$ ) and the open points the clusters with lower discrepancies $M_d/M_g<1.5$ ).
+ The solid and dashed lines show the relation between maximum. neutrino density and the temperature implied by 113. for neutrinos of ον and 1 ev respectively.," The solid and dashed lines show the relation between maximum neutrino density and the temperature implied by 13, for neutrinos of 2 ev and 1 ev respectively."
+ For neutrinos of 2 ev. this relation appears to form an upper envelope for the clusters. with the largest. discrepancies.," For neutrinos of 2 ev, this relation appears to form an upper envelope for the clusters with the largest discrepancies."
+ However. due to the extreme sensitivity of this density limit on neutrino mass 113). neutrinos with mass as low as 1 ον could not comprise the dark component.," However, due to the extreme sensitivity of this density limit on neutrino mass 13), neutrinos with mass as low as 1 ev could not comprise the dark component."
+ Ll have caleulated the structure of a. self-gravitating system supported by the degenerate pressure of neutrinos (a neutrino star)., I have calculated the structure of a self-gravitating system supported by the degenerate pressure of neutrinos (a neutrino star).
+" In the non-relativistic case there is à pressure-density relation. of the form p.=Ape’? where AO=5.5107(m,ery5 and density is in &/enmr'."," In the non-relativistic case there is a pressure-density relation of the form $p_\nu = K{\rho_\nu}^{5/3}$ where $K = 5.5\times 10^{32}(m_\nu/
+2\, ev)^{-{8\over 3}}$ , and density is in $^3$ ."
+ Inserting this relation in 11. LE determine the run of density in," Inserting this relation in 1, I determine the run of density in"
+angular separation of from high mass bright stars and 700 K from lower mass stars.,angular separation of from high mass bright stars and 700 K from lower mass stars.
+" Adding a second filter pair considerably improves these results by 200 K, while adding a third pair confirms these limiting values."," Adding a second filter pair considerably improves these results by 200 K, while adding a third pair confirms these limiting values."
+" With the considered data analysis methods and according to the evolutionary models from ? for the COND atmosphere models, we can estimate that in a very young system of 10 Myr, we should be able to characterize a planet of 1 with H2H3 at separations larger than around a low mass star (MO at 10 pc) where the star-planet contrast is favorable, but only further than around a high mass star (FO at 10 pc) where the contrast difference is larger."," With the considered data analysis methods and according to the evolutionary models from \citet{baraffe2003} for the COND atmosphere models, we can estimate that in a very young system of 10 Myr, we should be able to characterize a planet of 1 with H2H3 at separations larger than around a low mass star (M0 at 10 pc) where the star-planet contrast is favorable, but only further than around a high mass star (F0 at 10 pc) where the contrast difference is larger."
+" With two filter pairs, the limit would be around a low mass star and around high mass star."," With two filter pairs, the limit would be around a low mass star and around a high mass star."
+" For older systems, only planets of a few amasses of Jupiter could be characterized."," For older systems, only planets of a few masses of Jupiter could be characterized."
+" At 100 Myr, a Jupiter mass planet would remain out of reach for characterization with H2H3 filters around a high mass star, and only at separations larger than around a low mass star."," At 100 Myr, a Jupiter mass planet would remain out of reach for characterization with H2H3 filters around a high mass star, and only at separations larger than around a low mass star."
+" At this age, the limits of 700 K and 500 K which can be reached at small angular separation around high and low mass stars would respectively correspond to planets with masses of ~6.5 and ~3Μπιρ.."," At this age, the limits of 700 K and 500 K which can be reached at small angular separation around high and low mass stars would respectively correspond to planets with masses of $\sim$ 6.5 and $\sim$ 3."
+ Using improved signal extraction methods providing more accurate photometry of the companion would certainly push down those limits., Using improved signal extraction methods providing more accurate photometry of the companion would certainly push down those limits.
+ NCs are the cases where several models correspond to the flux measurements in all filter pairs with which they are detected., NCs are the cases where several models correspond to the flux measurements in all filter pairs with which they are detected.
+" From these remaining models, it is possible to determine if a combination of(T.y;; g)) is more represented than others, making this combination the most probable values of andg."," From these remaining models, it is possible to determine if a combination of; ) is more represented than others, making this combination the most probable values of and."
+". If several combinations are counted an equal number of times, an average value and an error can be determined for the values of andg."," If several combinations are counted an equal number of times, an average value and an error can be determined for the values of and."
+". In any case, the error is at least equal to the steps in the grids of models."," In any case, the error is at least equal to the steps in the grids of models."
+ The estimation of the most probable values for and has been performed for all non-uniquely characterized models at all simulated angular separations and magnitudes., The estimation of the most probable values for and has been performed for all non-uniquely characterized models at all simulated angular separations and magnitudes.
+ Figure 6 shows an histogram of the errors on and when using H2H3 filter pair for high mass (FO at 10 pc) and low mass (MO at 10 pc) stars., Figure \ref{fig:best_sequence_nochar} shows an histogram of the errors on and when using H2H3 filter pair for high mass (F0 at 10 pc) and low mass (M0 at 10 pc) stars.
+ NCs are mostly dominated by errors on the determination ofg.., NCs are mostly dominated by errors on the determination of.
+ In particular we see that around a low mass star where the contrast is more, In particular we see that around a low mass star where the contrast is more
+sample.,sample.
+" Notably, when we fit two gaussians (one for the disk and halo) to the Vg and [Fe/H] distributions, about 45% of the spectroscopic sample are halo stars (consistent with Spagnaetal.(2010) within ~5% errors), while halo stars make up only 8% of the complete sample."," Notably, when we fit two gaussians (one for the disk and halo) to the $V_{\phi}$ and [Fe/H] distributions, about $45\%$ of the spectroscopic sample are halo stars (consistent with \citet{Spagna2010} within $\sim 5\%$ errors), while halo stars make up only $8\%$ of the complete sample."
+" And when we compare the two panels in Figure A we are using exactly the same volume and exactly the same measurements: the only difference between the samples is that the spectroscopic sample includes only ~2% of all the stars, with the selection probability about 10 times higher for halo stars than for disk stars."," And when we compare the two panels in Figure \ref{f:spagna} we are using exactly the same volume and exactly the same measurements: the only difference between the samples is that the spectroscopic sample includes only $\sim 2\%$ of all the stars, with the selection probability about 10 times higher for halo stars than for disk stars."
+so-called survival function (or survivor function). S(/). of a given sample including right censored cata (ic. data with lower limits).,"so-called survival function (or survivor function), $S(t)$, of a given sample including right censored data (i.e., data with lower limits)."
+ A survival function is related to the DE bv the following relation: S(/)21 F(). To construct this estimator. we first. mathematically formulate the current problem according to Feigelson&Nel-son(1985) .," A survival function is related to the DF by the following relation: $S(t) = 1 - F(t)$ To construct this estimator, we first mathematically formulate the current problem according to \citet{FN85}."
+ selectedSince most of the astronomical observational data are with a certain detection limit. usually a sample includes upper limits. instead. of lower limits.," Since most of the astronomical observational data are selected with a certain detection limit, usually a sample includes upper limits, instead of lower limits."
+ 1n statistical terminology. lower-limiut data are called. 7Lleft-censored”.," In statistical terminology, lower-limit data are called ``left-censored''."
+ Let 2E.1E...T£ denote measurements. where the superscript £ means “left”.," Let $T_1^L, T_2^L, \dots, T_n^L$ denote measurements, where the superscript $L$ means “left”."
+ LODE<<ot). where A; is its upper limit. then we cannot obtain the exact value of d; but the upper limit24).," If $T_i^L < A_i$, where $A_i$ is its upper limit, then we cannot obtain the exact value of $T_i$ but the upper limit$A_i$."
+ Phe information available is a combination of {λα=(05.1) and. [45.0]. [or TLmo and 1<ib. respectively.," The information available is a combination of $\{X_i^L, \delta_i^L\} = \{T_i, 1\}$ and $\{A_i, 0\}$ for $T_i^L \ge A_i$ and $T_i^L < A_i$, respectively."
+ Suppose that 74 js à sample crawn from a clistribution PR)=DP<(0 unionswhere PE) denotes the probability of an event , Suppose that $T_i^L$ is a random sample drawn from a distribution $F^L(t) = P(T_i \le t)$ where $P(E)$ denotes the probability of an event $E$.
+bolt.nn» are mutually independent. identically clistributed ο to as LID) ancl indepencent of the (ruc measurementscensorship.FE Έλτον this statistical model is called. random ," If $\{A_i\}_{i = 1, 2, \dots, n}$ are mutually independent, identically distributed (referred to as IID) and independent of the true measurements$\{T_i^L\}_{i = 1, 2, \dots, n}$ , this statistical model is called random censorship."
+We can easily. transform. left-censored data to right-censored. ones by setting a constant Al: 4;—M TEN;=Al Nb. €;—M APLδι=OF.," We can easily transform left-censored data to right-censored ones by setting a constant $M$; $T_i = M - T_i^L$ $X_i = M - X_i^L$ , $C_i = M - A_i^L$, $\delta_i = \delta_i^L$."
+ μμ.» represent right-censored. data.," Then, $\{X_i,\delta_i\}_{i = 1, 2, \dots, n}$ represent right-censored data."
+" The IxM estimator has the following form: where rf, denote distinct. ordered observed. values in which ties are identified. and à;=3xo,ο d;=E:[etr= rik "," The KM estimator has the following form: where $x'_{(j)}$ denote distinct, ordered observed values in which ties are identified, and $n_j = \#\left\{ k; x_k \ge x'_{(j)} \right\}$, $d_j = \#\left\{ k; x_k = x'_{(j)} \right\}$ ."
+A detailed but comprehensive derivation is found in. e.g.. Feigelson&Nelson(1985) or Ixalblleisch&Prentice (2002).," A detailed but comprehensive derivation is found in, e.g., \citet{FN85} or \citet{KP02}."
+. Once we get the WAL estimate of the survival function. S(D. we ean convert it into the estimate of the DE. (0). as Then. In this analysis. we set. AZ=0. Le. we made a flip of a slen sd.," Once we get the KM estimate of the survival function, $\hat{S}(t)$, we can convert it into the estimate of the DF, $\hat{F}(t)$, as Then, In this analysis, we set $M=0$, i.e., we made a flip of a sign $t \rightarrow -t$."
+ In order to see if these two DES are different. we should perform a statistical test.," In order to see if these two DFs are different, we should perform a statistical test."
+ Though there are a few methods to do so. we adopt the Mantel-HLlaenszel. (ALLL) two-sample test (also often referred to as the logrank test. but the latter denotes some variations)," Though there are a few methods to do so, we adopt the Mantel-Haenszel (MH) two-sample test (also often referred to as the logrank test, but the latter denotes some variations)."
+ (20one derivation is shown in. e.g.. Ixalbllieisch&Prent," The derivation is shown in, e.g., \citet{KP02}."
+"ice"" The MET test. makes use of the values as follows: we sort all the galaxy sample along with the(minus logarithmic) luminosity. logLig;- "," The MH test makes use of the values as follows: first we sort all the galaxy sample along with the (minus logarithmic) luminosity, $- \log {L_{\rm PAH}}_{,j}$."
+"‘Then. we define mean numbers of galaxies at each luminosity logLyx; for both samples. mi;—dj where 4 denotes the label of the two samples. Πο and. ER""disk. and njConan|ona; We also define variance (Fk and /( again denote el and. disk"") audcovariance! Cer]=ο."," Then, we define mean numbers of galaxies at each luminosity $- \log {L_{\rm PAH}}_{,j}$ for both samples, $m_{k,j} \equiv d_{j} \frac{n_{k,j}}{n_j}$ where $k$ denotes the label of the two samples, “ell” and “disk”, and $n_j \equiv n_{{\rm ell},j} + n_{{\rm disk},j}$ We also define variance $k$ and $\ell$ again denote “ell” and “disk”), and; $v_{k\ell,j} = -v_{kk,j}$."
+ Consider a deviation vector (X.Aa.κ»Mists)= ) ," Consider a deviation vector $\left(\sum_j \Delta_{{\rm ell},j}, 
+\sum_j \Delta_{{\rm disk},j} \right) = 
+\left( \sum_j d_{{\rm ell},j} - m_{{\rm ell},j} , 
+\sum_j d_{{\rm disk},j} - m_{{\rm disk},j} , \right)$ ."
+Since the dof is 2e»f=1 in the current problem. MAaj=Nai.," Since the dof is $2 - 1 = 1$ in the current problem, $\sum_j \Delta_{{\rm ell},j} = -\sum_j \Delta_{{\rm disk},j}$."
+" Then. simply we can use a statistic Nafen,"," Then, simply we can use a statistic ${\Delta_{k}}^2/v_{kk}$."
+ M the two samples are crawn from the same DE (this is the null hypothesis to be. tested). this statistic should. obey the vOdistribution with dof—1.," If the two samples are drawn from the same DF (this is the null hypothesis to be tested), this statistic should obey the $\chi^2$ -distribution with $\mbox{dof}=1$ ."
+ From the sample of e and b. wehave 47=0.694 and {951.respectively.," From the sample of $a$ and $b$ , wehave $\chi^2= 0.694$ and 0.951,respectively."
+ Probabilities of having these values under the x7-distribution are 0.405 and 0.329., Probabilities of having these values under the $\chi^2$ -distribution are 0.405 and 0.329.
+ Phen. if we set the confidence limit of 0.05. we cannot reject the null hypothesis. te. the two DES are not significantly dillerent.," Then, if we set the confidence limit of 0.05, we cannot reject the null hypothesis, i.e., the two DFs are not significantly different."
+The initial delay in the light curve larger radii progenitors is due to the shock taking longer to arrive al the surface.,The initial delay in the light curve larger radii progenitors is due to the shock taking longer to arrive at the surface.
+ The shock wave can be identified in figure 2 by the deviation ol the long dashed lines and the solid lines for each of the three models., The shock wave can be identified in figure 2 by the deviation of the long dashed lines and the solid lines for each of the three models.
+ The time the shock wave spends inside (he largest radius progenitor is 8 davs compared to 2 hours for the smallest radius., The time the shock wave spends inside the largest radius progenitor is 8 days compared to 2 hours for the smallest radius.
+ It ean be seen in Figs., It can be seen in Figs.
+ 2a. and 26 that the shock in model A has already reached the surface by ( = 4.2 hours. while both models D and C have shocks moving down the density eradient near 10. AI. ancl 7.5 M. respectively.," 2a, and 2c that the shock in model A has already reached the surface by t = 4.2 hours, while both models B and C have shocks moving down the density gradient near 10 $_\odot$ and 7.5 $_\odot$, respectively."
+ In Figure 2c the temperature profiles clearly show a forward shock and a reverse shock forming an mesa in the middle., In Figure 2c the temperature profiles clearly show a forward shock and a reverse shock forming an mesa in the middle.
+ In the largest radii model (he deposited shock energv will not do as much PdV work. since it is already in an expanded state when the shock arrives. and the material will slay hotter longer.," In the largest radii model the deposited shock energy will not do as much PdV work, since it is already in an expanded state when the shock arrives, and the material will stay hotter longer."
+ Once the shock wave emerges for models with increasingly larger radii the radiation has a smaller diffusion time and the photon flux is higher accounting for the progressively broader initial peak and plateau., Once the shock wave emerges for models with increasingly larger radii the radiation has a smaller diffusion time and the photon flux is higher accounting for the progressively broader initial peak and plateau.
+ A comparison of the diffusion times between the three models clarifies this: The ratio of model A to D is 10. model A to C is 100 using diffusion time scale at constant opacity e 3&8 M/(dx The initial peak of model C has a 4 day rise time compared to the almost instantaneous rise time of models A and D (see Fig.," A comparison of the diffusion times between the three models clarifies this; The ratio of model A to B is 10, model A to C is 100 using diffusion time scale at constant opacity $\sim$ $\kappa$ $\pi$ The initial peak of model C has a 4 day rise time compared to the almost instantaneous rise time of models A and B (see Fig."
+ Ia. Ib. le. 2d).," 1a, 1b, 1c, 2d)."
+ This ean be explained by the behavior of the shock wave in low. constant density material.," This can be explained by the behavior of the shock wave in low, constant density material."
+" As the shock nears the surface of model C it becomes smoothed out and looses its ""shock"" definition and more mass has a higher temperature. eiving a broad peak to the photosperic temperature evolution."," As the shock nears the surface of model C it becomes smoothed out and looses it's “shock” definition and more mass has a higher temperature, giving a broad peak to the photosperic temperature evolution."
+ This is unlike (he smaller radii models which have a sharper. more defined shock and (hus achieves a higher temperature. but then expands and cools The difference between the radii can explain the maximum LC Iuninosity even though the photospheric temperature in the largest radii model is lower.," This is unlike the smaller radii models which have a sharper, more defined shock and thus achieves a higher temperature, but then expands and cools The difference between the radii can explain the maximum LC luminosity even though the photospheric temperature in the largest radii model is lower."
+ The ratio of the largest to smallest radius model is 100. while the ratio of the breakout temperature is only about 2.5 (Fig.," The ratio of the largest to smallest radius model is 100, while the ratio of the breakout temperature is only about 2.5 (Fig."
+ Ib)., 1b).
+ This leads to a factor of about 256 brighter in luminosity (6 in magnitude. Figs.," This leads to a factor of about 256 brighter in luminosity (6 in magnitude, Figs."
+ la. le. Id) for the larger radius progenitor.," 1a, 1c, 1d) for the larger radius progenitor."
+ For smaller stus much more of the internal enerev deposited bv the shock is transformed into PAV work in order to unbind the star., For smaller stars much more of the internal energy deposited by the shock is transformed into PdV work in order to unbind the star.
+ The initial burst is much brighter. seen as à spike. but [alls rapidly with the expansion.," The initial burst is much brighter, seen as a spike, but falls rapidly with the expansion."
+ At late times (he liminosity profiles are a useful diagnostic for the gamma-ray deposition., At late times the luminosity profiles are a useful diagnostic for the gamma-ray deposition.
+ All three models are shown at ( = 94 davs (Model A long dash. B dotted. C small dash). and have indistinguishable luminosity profiles at ( = 189 (solid lines). ancl 379 (clark solid," All three models are shown at t = 94 days (Model A long dash, B dotted, C small dash), and have indistinguishable luminosity profiles at t = 189 (solid lines), and 379 (dark solid"
+kinematic determination.,kinematic determination.
+ The radial velocities used are the weighted. averages from all previous measurements studies of NGC 5128 and are described completely in Ix. Wooclley et al. (, The radial velocities used are the weighted averages from all previous measurements studies of NGC 5128 and are described completely in K. Woodley et al. (
+2006. in preparation).,"2006, in preparation)."
+" The GC's were assigned weights in the kinematic fitting by considering two components to their associated uncertainties: the individual observational uncertainty i vc,. and (he random velocity component of the GCS evident bv the large dispersion in the confirmed GC velocities."," The GCs were assigned weights in the kinematic fitting by considering two components to their associated uncertainties: the individual observational uncertainty in $v_r$, and the random velocity component of the GCS evident by the large dispersion in the confirmed GC velocities."
+ In almost all cases. this latter term dominates. leaving the clusters with verv similar weights.," In almost all cases, this latter term dominates, leaving the clusters with very similar weights."
+" The fitted values of QF and ©, were determined lor projected galactocentric radial bins with inner and outer radii listed in column 1 of Table 1..", The fitted values of $\Omega R$ and $\Theta_o$ were determined for projected galactocentric radial bins with inner and outer radii listed in column 1 of Table \ref{tab:GC}.
+ Successive columns give the mean radius and the outer radius in kpc assuming a distance of 3.9 Alpe. the number of GCs. the rotation amplitude in kms !. the rotation axis of the GCs in each circular radial bin in degrees East of North. the velocity dispersion in km |. the projected mass correction factor. the pressure supported component of mass. the rotationally supported component of mass. and the total mass (see Section 3. for the mass discussion).," Successive columns give the mean radius and the outer radius in kpc assuming a distance of 3.9 Mpc, the number of GCs, the rotation amplitude in km $^{-1}$, the rotation axis of the GCs in each circular radial bin in degrees East of North, the velocity dispersion in km $^{-1}$, the projected mass correction factor, the pressure supported component of mass, the rotationally supported component of mass, and the total mass (see Section \ref{mass} for the mass discussion)."
+ The radial bin 0-5 kpc has a different. rotation solution [rom bins 5-10. 10-15. 15-25. 25-50. and 0-50 kpc.," The radial bin 0-5 kpc has a different rotation solution from bins 5-10, 10-15, 15-25, 25-50, and 0-50 kpc."
+ In bin 0-5 kpc. these GC's nominally rotate around a similar rotation axis as (he remaining bins. but in an opposing direction.," In bin 0-5 kpc, these GCs nominally rotate around a similar rotation axis as the remaining bins, but in an opposing direction."
+ This result may. be real for this innermost radial bin since it contains 55Hr) clusters distributed evenly in angle over (hat portion of the skv (see Figure 2))., This result may be real for this innermost radial bin since it contains 55 clusters distributed evenly in angle over that portion of the sky (see Figure \ref{fig:kinGC}) ).
+ However. the rotation zunplitude is scarcely different [rom zero. and the svstem is clearly. dominated by internal random motion there.," However, the rotation amplitude is scarcely different from zero, and the system is clearly dominated by internal random motion there."
+ Fig., Fig.
+ 2 clisplavs the sine curve fits. defined by Equation 1.., \ref{fig:kinGC} displays the sine curve fits defined by Equation \ref{eqn:kin}.
+ Devond a galactocentrie distance of 15 kpc. the observational bias of the detections (most of the known outer clusters lie along the isophotal major axis of the galaxy. which is where previous searches have mainly concentrated) is evident bv a clear lack of known GC's near the photometric minor axis of O=119 and 299+5° (Dufouretal.1979).," Beyond a galactocentric distance of 15 kpc, the observational bias of the detections (most of the known outer clusters lie along the isophotal major axis of the galaxy, which is where previous searches have mainly concentrated) is evident by a clear lack of known GCs near the photometric minor axis of $\Theta = 119$ and $299\pm5^o$ \citep{dufour79}."
+. This bias leads to higher uncertainties in the determined kinematic quantities in (hese outer regions., This bias leads to higher uncertainties in the determined kinematic quantities in these outer regions.
+ Nevertheless. it is notable that the kinematies of the metal-poor and metal-rich populations of GCs in NGC 5128 do not show anv extreme differences in the outer-halo region and 1 (reat (hem as a combined sample in the Following discussion.," Nevertheless, it is notable that the kinematics of the metal-poor and metal-rich populations of GCs in NGC 5128 do not show any extreme differences in the outer-halo region and I treat them as a combined sample in the following discussion."
+ Full details of the kinematic solutions of the metallicity sub-populations are discussed in Ix. Woodley et al. (, Full details of the kinematic solutions of the metallicity sub-populations are discussed in K. Woodley et al. (
+2006. in preparation).,"2006, in preparation)."
+ The velocity dispersion of the GCS ranges from 105 to 166 km s... and the entire svstem as a whole has a mean dispersion of 12226 km !.," The velocity dispersion of the GCS ranges from 105 to 166 km $^{-1}$ , and the entire system as a whole has a mean dispersion of $123\pm6$ km $^{-1}$."
+ These results show striking resemblances to those of Peng.Ford.&Freeman(2004): rom an earlier ancl smaller sample. (μον [ind the velocity dispersion of the GCs in NGC 5128 to range from 100—150 km ! out to a projected radius of 50 kpe. with thelargest dispersion at 20 kpe (see their Figure 20).," These results show striking resemblances to those of \cite{pff04III}: from an earlier and smaller sample, they find the velocity dispersion of the GCs in NGC 5128 to range from $100-150$ km $^{-1}$ out to a projected radius of 50 kpc, with thelargest dispersion at 20 kpc (see their Figure 20)."
+The uncertainty in Eq. (,The uncertainty in Eq. (
+1) is the value of EWi4. which may be different for individual sources.,"1) is the value of $EW_{\rm ion}$, which may be different for individual sources."
+" We estimate the value of EW,., from the Boroson&Green(1992) sample.", We estimate the value of $EW_{\rm ion}$ from the \citet{bg92} sample.
+ This sample contains all 87 PG quasars (z«0.5) with high quality optical spectra., This sample contains all 87 PG quasars $z<0.5$ ) with high quality optical spectra.
+" The average value of EWy,; is 100 for 70 radio-quiet quasars.", The average value of $EW_{{\rm H}\beta}$ is 100 for 70 radio-quiet quasars.
+ Unlike blazars. the optical continuum of radio-quiet quasar has not been contaminated by the beamed synchrotron emission from the jet.," Unlike blazars, the optical continuum of radio-quiet quasar has not been contaminated by the beamed synchrotron emission from the jet."
+" We will take EW,=100 and use the broad emission line H. to estimate the tonizing continuum luminosity of BL Lae objects.", We will take $EW_{\rm ion}=100$ and use the broad emission line ${\rm H}\beta$ to estimate the ionizing continuum luminosity of BL Lac objects.
+ The BL Lae objects seem to follow the statistical behavior of quasars in the correlations between radio and broad line emission (Celotti.Padovani.&Ghisellini1997;Cao&Jiang 1999).," The BL Lac objects seem to follow the statistical behavior of quasars in the correlations between radio and broad line emission \citep{c97,cj99}."
+. It may imply that the properties of BLRs in BL Lae objects are not significantly different from that in quasars., It may imply that the properties of BLRs in BL Lac objects are not significantly different from that in quasars.
+" If the Ευ of BL Lac objects deviates from that of quasars systematically. then the estimated black hole mass could be modified with EW,,, (see further discussion in Sect."," If the $EW_{\rm ion}$ of BL Lac objects deviates from that of quasars systematically, then the estimated black hole mass could be modified with $EW_{\rm ion}$ (see further discussion in Sect."
+ 5)., 5).
+ For a standard thin disk. the tonizing continuum luminosity 18 mainly determined by the central black hole mass and accretion rate.," For a standard thin disk, the ionizing continuum luminosity is mainly determined by the central black hole mass and accretion rate."
+ For a low aceretion rate. the accretion flow will transit from standard thin disk type to the ADAF type (Narayan1995;Yi 1996). ," For a low accretion rate, the accretion flow will transit from standard thin disk type to the ADAF type \citep{ny95,y96}. ."
+The standard thin accretion disks are thought to be in most quasars (Koratkar&Blaes1999)., The standard thin accretion disks are thought to be in most quasars \citep{kb99}.
+. In thiscase. the luminosity at optical wavelength can be related with the disk luminosity by Ly77GAL jon(S100 Ad) (Kaspietal. 2000)..," In thiscase, the luminosity at optical wavelength can be related with the disk luminosity by $L_{\rm d}\simeq 9\lambda L_{\lambda, {\rm ion}}$ (5100 ) \citep{kas00}. ."
+" The central black hole mass My, is then estimated by My,~1079(La/|. where m=M/Mygu."," The central black hole mass $M_{\rm bh}$ is then estimated by $M_{\rm bh}\simeq 10^{-38} (L_{\rm d}/{\dot m}) M_\odot$, where $\dot m=\dot M/\dot M_{\rm Edd}$."
+ The ionizing luminosity {νο can be inferred from Lys. and we can finally estimate the black hole mass from the broad line luminosity Lys. if 71 is known.," The ionizing luminosity $L_{\lambda,\rm ion}$ can be inferred from $L_{\rm H\beta}$, and we can finally estimate the black hole mass from the broad line luminosity $L_{{\rm H}\beta}$, if $\dot m$ is known."
+ For m=1. the lower limit of the black hole mass is available.," For $\dot m=1$, the lower limit of the black hole mass is available."
+ The transition of the aceretion flow from the thin disk type to the ADAF type occurs while #7 decreases to a value below Moy (Narayan&Υ1995;Yi1996).," The transition of the accretion flow from the thin disk type to the ADAF type occurs while $\dot m$ decreases to a value below ${\dot
+m}_{\rm crit}$ \citep{ny95,y96}."
+. One possible mechanism for the transition might be the evaporation of the disk (Meyer&2000:Meyer-Hofmeister&Meyer 2001).," One possible mechanism for the transition might be the evaporation of the disk \citep{mm94,l99,rc00,mm01}."
+.Gu&Lu(2000) suggested that the transition can be triggered by the thermal instability of a radiation pressure-dominated standard aceretion disk., \citet{gl00} suggested that the transition can be triggered by the thermal instability of a radiation pressure-dominated standard accretion disk.
+ The spectrum of an ADAF: L(Myg.i.0.7) can. be calculated if the parameters Mjy. 7. 0. and the fraction of the magnetic pressure ./ are specified.," The spectrum of an ADAF: $L_{\lambda}(M_{\rm bh}, {\dot m}, \alpha, \beta)$ can be calculated if the parameters $M_{\rm bh}$, $\dot m$, $\alpha$, and the fraction of the magnetic pressure $\beta$ are specified."
+ The parameter ./ is defined as Pr2(1—Dpro., The parameter $\beta$ is defined as $p_{\rm m}=(1-\beta)p_{\rm tot}$.
+ We can calculate the spectra of ADAFs using the approach proposed by Mahadevan(1997)., We can calculate the spectra of ADAFs using the approach proposed by \citet{m97}.
+. For the fixed black hole mass My. the optical luminosity Li increases with a.," For the fixed black hole mass $M_{\rm bh}$, the optical luminosity $L_{\lambda}$ increases with $\alpha$."
+ For AGNs. the viscosity o could be as high as |. as suggested by Narayan(1996).," For AGNs, the viscosity $\alpha$ could be as high as 1, as suggested by \citet{n96}."
+. In this work. We set o=0.3 to calculate optical luminosity L4 (Yi1996;Choi.Yang.&Yi2001).," In this work, We set $\alpha=0.3$ to calculate optical luminosity $L_{\lambda}$ \citep{y96,c01}."
+. The fraction of magnetic pressure ./20.5 for equipartition cases., The fraction of magnetic pressure $\beta=0.5$ for equipartition cases.
+ In fact. our numerical results show that the maximaloptical luminosity ZA5 always requires ./—0.5. if all other parameters are fixed.," In fact, our numerical results show that the maximaloptical luminosity $L_{\lambda}^{\rm max}$ always requires $\beta=0.5$, if all other parameters are fixed."
+ The accretion rate m<m;c0.2807. should be satisfied for an ADAF (Mahadevan1997)., The accretion rate $\dot m<\dot m_{\rm crit} \simeq 0.28\alpha^2$ should be satisfied for an ADAF \citep{m97}.
+. So. varying the accretion rate 77. we can find the maximal optical luminosity at Aj numerically for à=0.3 and ./20.5 (see Fig.," So, varying the accretion rate $\dot m$, we can find the maximal optical luminosity at $\lambda_0$ numerically for $\alpha=0.3$ and $\beta=0.5$ (see Fig."
+ 1)., 1).
+ We plot the maximal value of ALAVA(Ào) varying with black hole mass Mpy in Fig., We plot the maximal value of $\lambda L_{\lambda}^{\rm max}(\lambda_0)$ varying with black hole mass $M_{\rm bh}$ in Fig.
+ 2., 2.
+" We can obtain a lower limit on the mass of the black hole from Ly,; using the relation ALT—My, plotted in Fig.", We can obtain a lower limit on the mass of the black hole from $L_{{\rm H}\beta}$ using the relation $\lambda L_{\lambda}^{\rm max}-M_{\rm bh}$ plotted in Fig.
+ 2., 2.
+ For BL Lae objects. most accretion power may be carried by strong jets and the accretion flows may probably be described by the adiabatic inflow-outflow solutions (ADIOSs) (Blandford.&Begelman1999).," For BL Lac objects, most accretion power may be carried by strong jets and the accretion flows may probably be described by the adiabatic inflow-outflow solutions (ADIOSs) \citep{bb99}."
+. In this case. the flow is fainter than the pure ADAF considered here.," In this case, the flow is fainter than the pure ADAF considered here."
+ The maximal AZA(Ao) is therefore still valid for an ADIOS., The maximal $\lambda L_{\lambda}^{\rm max}(\lambda_0)$ is therefore still valid for an ADIOS.
+ It will not affect our estimate of the black hole mass limits., It will not affect our estimate of the black hole mass limits.
+ Donatoetal.(2001) complied a sample of blazars observed in the X-ray band. which includes almost all HBLs. IBLs. and many LBLs.," \citet{d01} complied a sample of blazars observed in the X-ray band, which includes almost all HBLs, IBLs, and many LBLs."
+ We collect all BL Lae objects in their sample and | Jy+S4+S5 catalogues (Stickel.Meisenheimer.&Kühr1994:Stickel.&Kühr1994.1996) as our start sample.," We collect all BL Lac objects in their sample and 1 Jy+S4+S5 catalogues \citep{s94,sk94,sk96} as our start sample."
+ This sample includes most identified LBLs and HBLs., This sample includes most identified LBLs and HBLs.
+" We search the literature for all sources with broad emission line fluxes, and this leads to a sample of 23 sources (listed in Table 1)."," We search the literature for all sources with broad emission line fluxes, and this leads to a sample of 23 sources (listed in Table 1)."
+ Most of them are LBLs. except three HBLs: 06514428. Mkn 421. and Mkn 501.," Most of them are LBLs, except three HBLs: $+$ 428, Mkn 421, and Mkn 501."
+ The sources with relatively strong broad emission lines (EW>10 A)) are considered as genuine HPQs by Veron-Cetty&Veron(2000) and are therefore not included in our sample., The sources with relatively strong broad emission lines $EW>10$ ) are considered as genuine HPQs by \citet{vv00} and are therefore not included in our sample.
+The source 3C 279 (1253-055) ts classified as a BL Lac object in Veron-Cetty&Veron(2000) due to its small broad line equivalent width. though this source is usually regarded as a quasar in other literature.,"The source 3C 279 $-$ 055) is classified as a BL Lac object in \citet{vv00} due to its small broad line equivalent width, though this source is usually regarded as a quasar in other literature."
+ Many sources in this sample have detected or Ha line emission instead of H7., Many sources in this sample have detected or $\alpha$ line emission instead of $\beta$.
+ We use line ratios reported by Francis(1991) to estimate H./ line luminosity {η from the line luminosity of Ha orII., We use line ratios reported by \citet{f91} to estimate $\beta$ line luminosity $L_{{\rm H}\beta}$ from the line luminosity of $\alpha$ or.
+. The broad emission line luminosity Ly. are listed in Table 1., The broad emission line luminosity $L_{{\rm H}\beta}$ are listed in Table 1.
+ The distribution of the line luminosity Ly. is plotted in Fig., The distribution of the line luminosity $L_{{\rm H}\beta}$ is plotted in Fig.
+ 3., 3.
+ We use Eq. (, We use Eq. (
+1) and the line luminosity Ly; to estimate the ionizing continuum luminosity at 4861A.,1) and the line luminosity $L_{{\rm H}\beta}$ to estimate the ionizing continuum luminosity at 4861.
+". The central black hole masses My,, and Myij» can be estimated in the cases of a thin disk and an ADAF. respectively."," The central black hole masses $M_{\rm bh,1}$ and $M_{\rm bh,2}$ can be estimated in the cases of a thin disk and an ADAF, respectively."
+ The mass derived for the standard thin disk case also depends on accretion rate 7., The mass derived for the standard thin disk case also depends on accretion rate $\dot m$ .
+" The lower limit ofMy,, can be available by setting m= 1. If the transition of the accretion flow from the thin disk type to the ADAF type occurs ata~ aga. we can have an upper limit on Μι setting 71=0.025 for à= 0.3. For ADAFs. Μι. 15 the lower limit. since the maximal optical continuum luminosity iscalculated for the given black hole mass."," The lower limit of$M_{\rm bh,1}$ can be available by setting $\dot m=1$ If the transition of the accretion flow from the thin disk type to the ADAF type occurs at $\dot m\sim \dot m_{\rm crit}$ , we can have an upper limit on $M_{\rm bh,1}$ setting $\dot m=0.025$ for $\alpha=0.3$ For ADAFs, $M_{\rm bh,2}$ is the lower limit, since the maximal optical continuum luminosity iscalculated for the given black hole mass."
+ We list the derived black hole masses in Columns (6) and (7) of Table I., We list the derived black hole masses in Columns (6) and (7) of Table 1.
+2011).,.
+. The problems arise when (he exact evolution of the inner lobes of DDRGs is analvzed., The problems arise when the exact evolution of the inner lobes of DDRGs is analyzed.
+ These inner lobes appear rather slim on the high-dinamies radio maps. displaving in particular axial ratios in the range of 9—16. to be compared with the range of 2—5 established for the outer lobes of same svstems. and also for the lobes in ‘regular ERI radio galaxies.," These inner lobes appear rather slim on the high-dynamics radio maps, displaying in particular axial ratios in the range of $9 - 16$, to be compared with the range of $2 - 5$ established for the outer lobes of same systems, and also for the lobes in `regular' II radio galaxies."
+ This indicates that the restarted jets in DDRGs propagate within the environment much rarified wilh respect to the original IGM in which the outer lobes are evolving. as in [act expected in the framework of the jet intermittency scenario. bul vet much denser (han the one provided simply by the injection of a jet plasma to the outer lobes in the previous epoch of the jet aclivily (see.e.e..thedetailedanalysisbySalourisetal.2008).," This indicates that the restarted jets in DDRGs propagate within the environment much rarified with respect to the original IGM in which the outer lobes are evolving, as in fact expected in the framework of the jet intermittency scenario, but yet much denser than the one provided simply by the injection of a jet plasma to the outer lobes in the previous epoch of the jet activity \citep[see, e.g., the detailed analysis by][]{saf08}."
+. Thus required contamination of the remnant cocoons by the additional matter is one of the central questions regarding the nature of DDRGs (seethediscussionin[κακοetal.2000)., Thus required contamination of the remnant cocoons by the additional matter is one of the central questions regarding the nature of DDRGs \citep[see the discussion in][]{kai00}.
+. It is not clear. however. if any entrainment process operates on sulliciently short timescales to enrich substantially the outer cocoons with Che required amount of the cold matter from the IGM during the iletime of a radio source.," It is not clear, however, if any entrainment process operates on sufficiently short timescales to enrich substantially the outer cocoons with the required amount of the cold matter from the IGM during the lifetime of a radio source."
+ For (his reason. Brocksoppetal.(2007.2011). argued (hat racio emission observed [rom the inner structures of DDRGs is not due to the lobes of newly born jets. bul predominantly due to bow shocks induced within the extended lobes by only briefly interrupted relativistic outflows.," For this reason, \citet{bro07,bro11} argued that radio emission observed from the inner structures of DDRGs is not due to the lobes of newly born jets, but predominantly due to bow shocks induced within the extended lobes by only briefly interrupted relativistic outflows."
+ In (his scenario no contamination of (he extended/outer cocoon by Che additional matter would be needed., In this scenario no contamination of the extended/outer cocoon by the additional matter would be needed.
+ In this paper we consider and analvze a possibility Chat the structure of the lareest radio galaxy. known up to date with the total projected linear size of \INIpc. discovered by Machalskietal.(2008) is lormed by a restarted. rather than a primary jet activity.," In this paper we consider and analyze a possibility that the structure of the largest radio galaxy known up to date — with the total projected linear size of $D=4.69$ Mpc, discovered by \citet{mach08} — is formed by a restarted, rather than a primary jet activity."
+ The motivation lor the performed analvsis and new radio observations is due to the fact that the axial ratio lor the lobes in this outstanding example of a eiant. Ryx12. is of the same order or even larger than axial ratios for the inner lobes in all well studied DDRGs. aud therefore very different from the (vpical axial ratios characterizing ‘normal size FILII sources or other known GIGs.," The motivation for the performed analysis and new radio observations is due to the fact that the axial ratio for the lobes in this outstanding example of a giant, $R_{\rm T}\approx 12$, is of the same order or even larger than axial ratios for the inner lobes in all well studied DDRGs, and therefore very different from the typical axial ratios characterizing `normal size' II sources or other known GRGs."
+ Since in the framework of the standard ονπαΙσ]. model for classical doubles the axial ratio of the lobes depends predominantly on the jet. kinetic power relative to the density of the surrounding. must be a powerful source evolving in a particularly low-density environment. as indeed concluded in Machalskietal.(2008).," Since in the framework of the standard dynamical model for classical doubles the axial ratio of the lobes depends predominantly on the jet kinetic power relative to the density of the surrounding, must be a powerful source evolving in a particularly low-density environment, as indeed concluded in \citet{mach08}."
+. The question is if the low-density of the matter surrounding eijant lobes of is due (o a unique location of in a void region. or due to some previous jet activity epoch which took place before the formation of the currently observed lobes. and which rarified the IGM thereby Gin analogy with DDRGs).," The question is if the low-density of the matter surrounding giant lobes of is due to a unique location of in a void region, or due to some previous jet activity epoch which took place before the formation of the currently observed lobes, and which rarified the IGM thereby (in analogy with DDRGs)."
+ In (he former case. may be used as à particularly interesting (due to ils enormous size) probe of a ‘warm-hot phase of the interealactie medium. (hereafter WIIIM) within a filamentary galaxy. distribution far away lrom any group or cluster of galaxies.," In the former case, may be used as a particularly interesting (due to its enormous size) probe of a `warm-hot phase of the intergalactic medium' (hereafter WHIM) within a filamentary galaxy distribution far away from any group or cluster of galaxies."
+ In the latter, In the latter
+only use those galaxies for which there were 2ALASS data available [rom the 2nd incremental data release.,only use those galaxies for which there were 2MASS data available from the 2nd incremental data release.
+ Phe 2\LASS catalogue lists numerous magnitudes of which we used the default ‘Liducial circular magnitudes at 20 mag/seq.aresec'. because of the wide range of sizes and shapes of the galaxies in the comparison sample.," The 2MASS catalogue lists numerous magnitudes of which we used the default `fiducial circular magnitudes at 20 mag/sq.arcsec', because of the wide range of sizes and shapes of the galaxies in the comparison sample."
+ 3oselli et ((1998) eive their own near-H1 photometry., Boselli et (1998) give their own near-IR photometry.
+ We performed. a cross-check of their A-band. photometry between 2ZALASS values and found a fairly good consisteney: apart from some Se-type galaxies 2ALASS isophotal magnitudes were 0.15 mag fainter than the Bosclli values. with a scatter of 0.31 mag.," We performed a cross-check of their $K$ -band photometry between 2MASS values and found a fairly good consistency: apart from some Sc-type galaxies 2MASS isophotal magnitudes were 0.15 mag fainter than the Boselli values, with a scatter of 0.31 mag."
+ That the level of Boselli and 2PATASS magnitudes are so close. and that the sizes of galaxies in the other two samples are not significantly cillerent from the Boselli galaxies. gives us confidence to use PALASS magnitudes for the nearby comparison galaxies. ancl to directly compare them to the ‘total magnitudes! of the ELAIS objects.," That the level of Boselli and 2MASS magnitudes are so close, and that the sizes of galaxies in the other two samples are not significantly different from the Boselli galaxies, gives us confidence to use 2MASS magnitudes for the nearby comparison galaxies, and to directly compare them to the `total magnitudes' of the ELAIS objects."
+ Nevertheless. for consistency. reasons with Dale ancl Roussel samples. we decided to use only. those galaxies from Boselli which had 2\LASS photometry available.," Nevertheless, for consistency reasons with Dale and Roussel samples, we decided to use only those galaxies from Boselli which had 2MASS photometry available."
+ In. the following. with the Dale anc Roussel samples we mean those galaxies from the respective original works for which we found 23LASS magnitudes.," In the following, with the Dale and Roussel samples we mean those galaxies from the respective original works for which we found 2MASS magnitudes."
+ Our Boselli sample means those galaxies from Boselli with 2ALASS photometry. and those galaxies excluded. which already are included in the Roussel sample.," Our Boselli sample means those galaxies from Boselli with 2MASS photometry, and those galaxies excluded, which already are included in the Roussel sample."
+ As for the Z50-data. the standard CAAL Interactive Analysis (CLA) packages were used for pre-processing of," As for the -data, the standard CAM Interactive Analysis (CIA) packages were used for pre-processing of"
+rotators.,rotators.
+ Tence. the value πο~0.215 for the vacunu 907 rotator in Fie.," Hence, the value $V_{\rm
+drop}\sim 0.2 V_0$ for the vacuum $90^\circ$ rotator in Fig."
+ represeuts the quadrupolar contribution to the potential drop., represents the quadrupolar contribution to the potential drop.
+ At conductivity (σέ).~OL the maximal potential drop is close to the orthogonal vacunni rotator value of Vg~0.210 at all inclination augles.," At conductivity $(\sigma/\Omega)^2 \sim 0.04$, the maximal potential drop is close to the orthogonal vacuum rotator value of $V_{\rm drop}\sim 0.2 V_0$ at all inclination angles."
+ Qualitativelv. maeuctosplercs with finite conductivity (6/0)?>0.01 can effectively redistribute the surface charges and maintain the linear relatiouship between the maximal poteutial drop aud the spin-down Iuninositv.," Qualitatively, magnetospheres with finite conductivity $(\sigma/\Omega)^2>0.04$ can effectively redistribute the surface charges and maintain the linear relationship between the maximal potential drop and the spin-down luminosity."
+ Below (oQ3?=0.0L. the surface charge redistribution is incomplete (the monopolar surface charge is concentrated closer to the star in the corotating steady-state solution) aud the curves in Fie.," Below $(\sigma/\Omega)^2
+= 0.04$, the surface charge redistribution is incomplete (the monopolar surface charge is concentrated closer to the star in the corotating steady-state solution) and the curves in Fig."
+ eain a tilt more roimuinisceut of the vacuum solution., gain a tilt more reminiscent of the vacuum solution.
+ The potential drops are uot constant with inclination angle at a eiven conductivity. aud the linear relation (121) no longer holds.," The potential drops are not constant with inclination angle at a given conductivity, and the linear relation \ref{eq:spindown}) ) no longer holds."
+ We have explicitly verified that by the eud of the sinmlatiou these solutions reached steady-state iu the corotating frame., We have explicitly verified that by the end of the simulation these solutions reached steady-state in the corotating frame.
+ Redistribution of surface charee in nearvactin maegnetospheres also affects the spin-down power., Redistribution of surface charge in near-vacuum magnetospheres also affects the spin-down power.
+ Looking back at Fig., Looking back at Fig.
+ 2— aud the/ piecewise linear relations (9)). we see that the final of the spin-down transition between force-free aud vacunu rotators occurs for (0/0)?<0.0L," \ref{dipole} and the piecewise linear relations \ref{eq:spindownsigma}) ), we see that the final of the spin-down transition between force-free and vacuum rotators occurs for $(\sigma/\Omega)^2 < 0.04$."
+ Despite the fact that potential is no longer coustaut with angle for cach value of the conductivity in this regine. the transition in spin-down is still contiuuous and s1100tli.," Despite the fact that potential is no longer constant with angle for each value of the conductivity in this regime, the transition in spin-down is still continuous and smooth."
+ We have presented. ᾱ- continuous one-parameter family of pulsar magnuetosphere solutious that span the range between the vacuun and force-free launits., We have presented a continuous one-parameter family of pulsar magnetosphere solutions that span the range between the vacuum and force-free limits.
+" Each solution of the family is characterized bv the value. of the concuctivity paramcter, 70/Q. which is related to the maximum potential drop that a test particle can experience as it moves along field lines. Va."," Each solution of the family is characterized by the value of the conductivity parameter, $\sigma/\Omega$, which is related to the maximum potential drop that a test particle can experience as it moves along field lines, $V_{\rm drop}$."
+ The zero couductiity limit. 6/0 »0. >Vaopsaes Vields the vacua magnetosphere.," The zero conductivity limit, $\sigma/\Omega\to0$, $V_{\rm drop}\to V_{\rm
+drop,vac}$, yields the vacuum magnetosphere."
+ ItVarop shows a substantial effective poteutial drop that makes up a significant fraction of the rotation-iunduced potential ciffercuce between the pole aud the equator of the star (2060% depeuding ou the inclination)., It shows a substantial effective potential drop that makes up a significant fraction of the rotation-induced potential difference between the pole and the equator of the star $20{-}60\%$ depending on the inclination).
+ The high conductmityv lanit. σο>ore. vields essentiallv an ideal force-free iiagnuetosphere.," The high conductivity limit, $\sigma/\Omega\to\infty$, yields essentially an ideal force-free magnetosphere."
+ Uulike iu vacunuun. abundant maguctospleric charges short out any potential differences along feld lines. leading to vanishing effective. potential drop. Varco5»0.," Unlike in vacuum, abundant magnetospheric charges short out any potential differences along field lines, leading to vanishing effective potential drop, $V_{\rm drop,ff}
+\to 0$."
+ While in the vacuna case nearly all field limes. even those that extend bevoud the πο cvlinder. eventually retur to the star aud are formally closed. im the ideal force-free case a fraction of field lines open up aud reach infinity.," While in the vacuum case nearly all field lines, even those that extend beyond the light cylinder, eventually return to the star and are formally closed, in the ideal force-free case a fraction of field lines open up and reach infinity."
+ Our simulations show that resistive hiel-o pulsars spin down at least 3 times faster than resistive low-o aud vacunu pulsus (for the same value of inclination). in agreement with earlier ideal force-free simmlatious (SOG).," Our simulations show that resistive $\sigma$ pulsars spin down at least $3$ times faster than resistive $\sigma$ and vacuum pulsars (for the same value of inclination), in agreement with earlier ideal force-free simulations (S06)."
+ Qur resistive solutions bridge the eap between the force-ree and vacuuni hits by having imtermecdiate values of major magnetospheric parameters. such as effective )otenfial drop. O«<Viropaae. the fraction of open field lines. aud theVagos spin-down rate.," Our resistive solutions bridge the gap between the force-free and vacuum limits by having intermediate values of major magnetospheric parameters, such as effective potential drop, $0<V_{\rm
+drop}<V_{\rm drop,vac}$, the fraction of open field lines, and the spin-down rate."
+ Before discussing possible observational implications of hese solutions. we note that our approach to modeling resistive pulsar magnetosphneres is quite simplistic: we use a form of Ον Law in which we neelect several ας (Meier2001) auc we assume constant conductivity c/O throughout all space.," Before discussing possible observational implications of these solutions, we note that our approach to modeling resistive pulsar magnetospheres is quite simplistic: we use a form of Ohm's Law in which we neglect several terms \citep{Meier04} and we assume constant conductivity $\sigma/\Omega$ throughout all space."
+ The finite bulk conductivity used in this work can be thought of as not due to collisions. but due to the iusufiicicut plasma supply in the nagnctosplere that cannot short out all the accelerating ficlds.," The finite bulk conductivity used in this work can be thought of as not due to collisions, but due to the insufficient plasma supply in the magnetosphere that cannot short out all the accelerating fields."
+ It is possible that real maguetospheres have a lower (anomalous) conductivity in the current laver aud the current sheet than in the rest of the magnetosphere., It is possible that real magnetospheres have a lower (anomalous) conductivity in the current layer and the current sheet than in the rest of the magnetosphere.
+ We welt also expect a lower value of conductivity ou the open field lines than on the closed field lines., We might also expect a lower value of conductivity on the open field lines than on the closed field lines.
+ Since our nodels do not explicitly mocel the plasima fiuid. refsecintro)) the plasma velocity component along the direction of maguetie field is unconustraimed. which jecessitates the choice of a frame din which to write the Olüuu's Law.," Since our models do not explicitly model the plasma fluid \\ref{sec:intro}) ), the plasma velocity component along the direction of magnetic field is unconstrained, which necessitates the choice of a frame in which to write the Ohm's Law."
+ We choose the so-called iiuinunuu velocity rane. which has a number of attractive properties (see refsecicurrentder)). vet this is nof a covariant choice.," We choose the so-called minimum velocity frame, which has a number of attractive properties (see \\ref{sec:currentder}) ), yet this is not a covariant choice."
+ With these miportaut caveats in mind. let us now for the sake of arguneut take our model seriously as a description of pulsar magnuetosphere aud consider the COMSCGUCLLCCS.," With these important caveats in mind, let us now for the sake of argument take our model seriously as a description of pulsar magnetosphere and consider the consequences."
+ Tuterimittent pulsars offer a unique test bed of pulsar theory (e.g... PSR B1931|21. Ixxiuueretal.2006.. and PSR J18532|0029. Lvue2009:: see also Waneetal.2007.. Zhangetal. 2007.. anc Timokhin 2010)).," Intermittent pulsars offer a unique test bed of pulsar theory (e.g., PSR B1931+24, \citealt{Kramer06}, and PSR J1832+0029, \citealt{Lyne09}; see also \citealt{Wang07}, \citealt{Zhang07}, and \citealt{Timokhin10}) )."
+" Such pulsars switch between an “on. radio-loud. state im which they behave like normal radio pulsars. and an ""off. radio-quiet. state im which they produce no detectable radio enüssiou."," Such pulsars switch between an “on”, radio-loud, state in which they behave like normal radio pulsars, and an “off”, radio-quiet, state in which they produce no detectable radio emission."
+" The two intermittent pulsars. for which published data exist. have quite different duty cvcles: PSR D1951|21. with a period Pzz0.8 s. evcles through the von”""off? sequence of states approximately once a month. whereas PSR J1s832|0029. with a period"," The two intermittent pulsars, for which published data exist, have quite different duty cycles: PSR B1931+24, with a period $P\approx0.8$ s, cycles through the “on”–”off” sequence of states approximately once a month, whereas PSR J1832+0029, with a period"
+In this section we describe simple analytical models for wind emission and. apply them to our sample.,In this section we describe simple analytical models for wind emission and apply them to our sample.
+ Initially we are not interested. in where the wine comes from or what it consists of. rather that a eeneric wind with a particular composition. being ejected at à particular rate from a source at a given distance from us. would. produce radio emission.," Initially we are not interested in where the wind comes from or what it consists of, rather that a generic wind with a particular composition, being ejected at a particular rate from a source at a given distance from us, would produce radio emission."
+ The optically thin lux from a source with these parameters can be calculated and is given by where Aisin M. +. [ids the mean atomic weight. Moy is the wind velocity in km ty is in Hz. 5 is the ratio of electron to ion number density. g is the Gaunt factor. Z is the ionic charge ancl £ is the distance in kpe (Wright Barlow 1975).," The optically thin flux from a source with these parameters can be calculated and is given by where $\dot{M}$ is in $_\odot$ $^{-1}$, $\mu$ is the mean atomic weight, $v_{\rm km/s}$ is the wind velocity in km $^{-1}$, $\nu$ is in Hz, $\gamma$ is the ratio of electron to ion number density, $g$ is the Gaunt factor, $Z$ is the ionic charge and $D$ is the distance in kpc (Wright Barlow 1975)."
+ Phe Gaunt factor here has a slight dependence on temperature and composition and if one does not consider extremely hot winds. then it can be approximated by (Leitherer Robert 1991).," The Gaunt factor here has a slight dependence on temperature and composition and if one does not consider extremely hot winds, then it can be approximated by (Leitherer Robert 1991)."
+ We can now investigate in more detail from where the wind originates. what is its expected composition and what is the expected radio emission from such a region.," We can now investigate in more detail from where the wind originates, what is its expected composition and what is the expected radio emission from such a region."
+ To ereate sustained mass transfer. between the binary components either the is undergoing Roche lobe overflow. the wind is providing the accretion via Boncli-Llovle accretion. or both.," To create sustained mass transfer between the binary components either the is undergoing Roche lobe overflow, the wind is providing the accretion via Bondi-Hoyle accretion, or both."
+ Which scenario occurs depends on the nature of the secondary. and the binary. parameters.," Which scenario occurs depends on the nature of the secondary, and the binary parameters."
+ The svmbiotic svstem AG Dra with a ~1.5 M; giant. à 0.5 ALy white dwarf and a period of 544 days does not normally fill its Roche lobe (the stellar raclius of the Ix-giant is 30 lt; and the Roche lobe size is ~170 Roe).," The symbiotic system AG Dra with a $\sim 1.5$ $_{\sun}$ K-giant, a $\sim 0.5$ $_{\sun}$ white dwarf and a period of 544 days does not normally fill its Roche lobe (the stellar radius of the K-giant is $\sim 30$ $_{\sun}$ and the Roche lobe size is $\sim 170$ $_{\sun}$ )."
+ Since the svstem must be accreting then the logical conclusion is that the secondary is on the Asvniptotic Giant Branch (AGB) ancl feeding the primary by a slow. dense wind.," Since the system must be accreting then the logical conclusion is that the secondary is on the Asymptotic Giant Branch (AGB) and feeding the primary by a slow, dense wind."
+ Out of the four sources considered. here. this is the most likely to be feeding its environment. directly bv a wind.," Out of the four sources considered here, this is the most likely to be feeding its environment directly by a wind."
+" Calculating the expected Uux from. this wind uses à temperature. velocity. outllow and distance of 2 — 150000 IK. e 2 30 km to M = 12LT «10 M4 vr and D = 1.56L.81 kpe from Tomosya ""Tomov (1999). compositions of y=1.4. 7?= land Z=1 from Nussbaumoer Vogel (1987) and Spergel. Giuliani Knapp (1983)."," Calculating the expected flux from this wind uses a temperature, velocity, outflow and distance of $T$ = 000 K, $v$ = 30 km $^{-1}$, $\dot{M}$ = 1.2–1.7 $\times 10^{-8}$ $_{\sun}$ $^{-1}$ and $D$ = 1.56–1.81 kpc from Tomova Tomov (1999), compositions of $\mu = 1.4$, $\gamma = 1$ and $Z = 1$ from Nussbaumer Vogel (1987) and Spergel, Giuliani Knapp (1983)."
+ οσο parameters vield a flux at 5 οκ of Sc = 1226 quls., These parameters yield a flux at 5 GHz of $S_{C}$ = 12–26 $\umu$ Jy.
+ In contrast to ACG) Dra. the three other sources either have a small secondary companion and therefore Roche lobe overllow would be providing the mass-transfer. or as in the case lor V1974 Ce. radio emission would have faced: aud dissipated from the time of the last outburst in 1992.," In contrast to AG Dra, the three other sources either have a small secondary companion and therefore Roche lobe overflow would be providing the mass-transfer, or as in the case for V1974 Cyg, radio emission would have faded and dissipated from the time of the last outburst in 1992."
+ In the previous section it was shown that only the AGB binary AG Dra would have a strong enough secondary wind to emit significant radio emission., In the previous section it was shown that only the AGB binary AG Dra would have a strong enough secondary wind to emit significant radio emission.
+ However. if the WD is raciating sullicient soft. N-ravs then a corona will be formed around the secondary. star. which causes a strong stellar wind to be evaporated (Basko Sunvaey 1973).," However, if the WD is radiating sufficient soft X-rays then a corona will be formed around the secondary star, which causes a strong stellar wind to be evaporated (Basko Sunyaev 1973)."
+" This type of interaction has been investigated. in the context. of the supersoft sources by van Teeseling Ixing (1998. hereafter VEIO9S) and they show that the mass loss rate in the secondary due to irradiation is where © is an irradiation ellicieney factor. re is the raclius of the secondary in solar units. eo is orbital separation in units of 10? m. m» is the mass of the secondary in solar units. Lao is the irradiation luminosity in units of 107 3 (=10"" erg !). i, is an ionisation ellicicney parameter. aud Ad is given in units of 10.""."," This type of interaction has been investigated in the context of the supersoft sources by van Teeseling King (1998, hereafter vTK98) and they show that the mass loss rate in the secondary due to irradiation is where $\phi$ is an irradiation efficiency factor, $r_2$ is the radius of the secondary in solar units, $a_{9}$ is orbital separation in units of $10^{9}$ m, $m_2$ is the mass of the secondary in solar units, $L_{30}$ is the irradiation luminosity in units of $10^{30}$ J $^{-1}$ $\equiv 10^{37}$ erg $^{-1}$ ), $\eta_s$ is an ionisation efficiency parameter, and $\dot{M}_{7}$ is given in units of $10^{-7}$."
+ Dt is reasonable. therefore. o assume that since the supersoft sources produce soft. X-ravs curing quiescence from their sustained nuclear burning. hat the majority (if not all) of the supersoft sources have mass losses which are wind cominatect.," It is reasonable, therefore, to assume that since the supersoft sources produce soft X-rays during quiescence from their sustained nuclear burning, that the majority (if not all) of the supersoft sources have mass losses which are wind dominated."
+ Yo obtain an estimate for the mass. loss [from he secondary due to irradiation from the WD. some assumptions to the &ecometry of the svstem need to be mace.," To obtain an estimate for the mass loss from the secondary due to irradiation from the WD, some assumptions to the geometry of the system need to be made."
+ In equation 3.. the variable © is an efficiency parameter representing the fraction of the companions face which is xing raciated and the fraction of the wind mass escaping he svstem.," In equation \ref{mass-loss-equation}, the variable $\phi$ is an efficiency parameter representing the fraction of the companion's face which is being radiated and the fraction of the wind mass escaping the system."
+ Commentary on the importance of ó is given in Ixniggc. Wine Patterson (2000). and a value of 1. is expected.," Commentary on the importance of $\phi$ is given in Knigge, King Patterson (2000), and a value of 1 is expected."
+ The other unknown parameter in equation 3. is he clliciency of the incident radiation in driving a wind. 5..," The other unknown parameter in equation \ref{mass-loss-equation} is the efficiency of the incident radiation in driving a wind, $\eta_{s}$."
+ For a wide range of temperatures. including the expected emperature of the irradiation. (Lew 10 I). η.c1.," For a wide range of temperatures, including the expected temperature of the irradiation (few $\times
+10^{5}$ K), $\eta_{s} \simeq 1$."
+ We can therefore calculate an estimation for the mass-lossin he wind from the secondary. which is given in table 2..," We can therefore calculate an estimation for the mass-lossin the wind from the secondary, which is given in table \ref{system_parameters}."
+ ‘Together with the mass-doss rate and equation 1.. he Εικ from the wind can be found.," Together with the mass-loss rate and equation \ref{wind_fluxes}, the flux from the wind can be found."
+ Unfortunately. the expected. fluxes are not as precise as one hopes as a large number of important parameters such as composition (yr. ~ and Z). wind temperature ancl velocity (77 anc ry) are unknown.," Unfortunately, the expected fluxes are not as precise as one hopes as a large number of important parameters such as composition $\mu$, $\gamma$ and $Z$ ), wind temperature and velocity $T$ and $v_{\infty}$ ) are unknown."
+ van leeseling Ixing (1998) consider to be the main absorber of the WD radiation. which has implications to the composition and temperature of the resulting wind.," van Teeseling King (1998) consider to be the main absorber of the WD radiation, which has implications to the composition and temperature of the resulting wind."
+ The temperature of the wind will range from the ionisation temperature of to the temperature of the incident radiation. so we use a wind temperature of Z/c5«10% in this paper.," The temperature of the wind will range from the ionisation temperature of to the temperature of the incident radiation, so we use a wind temperature of $T \simeq 5 \times 10^{5}$ in this paper."
+ The composition of a purely wind would eive values for the mean atomic weight (y= 4). the electron to ion number density ratio (=2) and the ionic charge," The composition of a purely wind would give values for the mean atomic weight $(\mu = 4)$ , the electron to ion number density ratio $(\gamma = 2)$ and the ionic charge"
+'efseciEP.. the radial variation of the diffusion coefficieut leads to an additional contribution to the advected flux of planets.,", the radial variation of the diffusion coefficient leads to an additional contribution to the advected flux of planets."
+ Since —OD/OJx—(Y+Lh)J5V for Xoxi./ and since fk=—3/2 in he MMSN aud &>—3/2 in the alpha disk models. this additional term tends to advect planets luwza rdi all the disk models considered.," Since $-\partial D/\partial J\propto -(7+4k)\,J^{6+4k}$ for $\bar{\Sigma}\propto r^k$, and since $k=-3/2$ in the MMSN and $k>-3/2$ in the alpha disk models, this additional term tends to advect planets inward in all the disk models considered."
+" We cau celine a local advection timescale associated with his new contribution. J/(0D/0.). and compare it to the values of /44, and fai defined above."," We can define a local advection timescale associated with this new contribution, $J/(\partial D/\partial J)$, and compare it to the values of $t_{\rm{mig}}$ and $t_{\rm{diff}}$ defined above."
+ We iud hat this adveetion titjescale is everywhere exactly equal to {μμ in the MNMSN. while it sslοιτο than {μή everywhere in our alpha disk moclels (as shown explicitly for the a=0.02 Case by he short-daslied liue in panel IV of Fig. 1)).," We find that this advection timescale is everywhere exactly equal to $t_{\rm{diff}}$ in the MMSN, while it is shorter than $t_{\rm{diff}}$ everywhere in our alpha disk models (as shown explicitly for the $\alpha=0.02$ case by the short-dashed line in panel IV of Fig. \ref{fig:4dprop}) )."
+ Consequently. even in disk regious where cliffSm uominally dominates over migration in otiw models (as measured by the ratio /qir//mis). this additional contribution cdie to racially inhomogeneous diffusion should still cause a rather significant effective advection o ‘planets towards the star.," Consequently, even in disk regions where diffusion nominally dominates over migration in our models (as measured by the ratio $t_{\rm{diff}}/t_{\rm{mig}}$ ), this additional contribution due to radially inhomogeneous diffusion should still cause a rather significant effective advection of planets towards the star."
+ We have studied two differeut types of time-clepeucdent moclels for the diffusive migration of plauets i1 proo-planetary disks. correspoucliig approximately to a contiuuous auc a burst scenario for planet formation.," We have studied two different types of time-dependent models for the diffusive migration of planets in proto-planetary disks, corresponding approximately to a continuous and a burst scenario for planet formation."
+ In. both cases. for the sake of clarity. the source of planets was localized. taking te form of a sharply peaked Gaussian appreximatiug the delta source fuuction of Eq. (13)).," In both cases, for the sake of clarity, the source of planets was localized, taking the form of a sharply peaked Gaussian approximating the delta source function of Eq. \ref{eq:sourceterm}) )."
+ We have verified that our results do not depend sensitively on the detailed shape of the Caussiau source fuction adopted., We have verified that our results do not depend sensitively on the detailed shape of the Gaussian source function adopted.
+ As we shall see. tie. predictions for these rather different scenarios fr plauet formation are broadly consistent. witl each other.," As we shall see, the predictions for these rather different scenarios for planet formation are broadly consistent with each other."
+ Figre 2. shows steady-state cistributiiis Obtained by iunerically solving Eq. (36)), Figure \ref{fig:4steady} shows steady-state distributions obtained by numerically solving Eq. \ref{FPsource}) )
+ with a source teru centered at ry=10AU. steadily produ‘in@ earth-mass planets at a rate A=(10?vr)," with a source term centered at $r_S=10\au$, steadily producing earth-mass planets at a rate $\Lambda=(10^5\yr)^{-1}$."
+ The toplost curve in eacl pauel is the steady-state cistributiou., The topmost curve in each panel is the steady-state distribution.
+ The euclosed curves. ou the other haud. are snapstots of the decaying distribuious at selected time intervals alter the source is shut oll.," The enclosed curves, on the other hand, are snapshots of the decaying distributions at selected time intervals after the source is shut off."
+ We have veriiec that he steady-state cistributious for the MMSN in Panels I and HIE are fully ‘Ousistent with ie analytic solution. Eqs. (15))-(16)).," We have verified that the steady-state distributions for the MMSN in Panels I and III are fully consistent with the analytic solution, Eqs. \ref{eq:mu}) \ref{eq:fss}) ),"
+ lor power-law disks., for power-law disks.
+ A fiducial normalization of e—0.5 for to‘que [luctuations is assumed in all cases except in panel HI. where €—0.02 has been scaled dow1 to exterd the mieration-dominated regime out to 10AU in the MMSN.," A fiducial normalization of $\epsilon=0.5$ for torque fluctuations is assumed in all cases except in panel III, where $\epsilon =0.02$ has been scaled down to extend the migration-dominated regime out to $\sim 10\au$ in the MMSN."
+ Thi=. at 10AU. he disss lu Patels I and IV. are diffusion-doimiuated and the disks in Panels HE aud UI we inigratiou-douinatect.," Thus, at $10\au$, the disks in Panels I and IV are diffusion-dominated and the disks in Panels II and III are migration-dominated."
+ The clistinetn)1 between diffusion-dominated axl inigration-doiminated evolutions is most clearly ---Lrustrated by σοιipariug panels E aud HIE., The distinction between diffusion-dominated and migration-dominated evolutions is most clearly illustrated by comparing panels I and III.
+ While the initial steady-state peaked distribution is eroded xd advected for lle most yart in panel IL. with Bittle diffusiou at large radii. there is siguilicautly inore diffusion a arge radii in panel 1. That distiuction is uot as clear iu the alpha disk models. however.," While the initial steady-state peaked distribution is eroded and advected for the most part in panel III, with little diffusion at large radii, there is significantly more diffusion at large radii in panel I. That distinction is not as clear in the alpha disk models, however."
+ Note tla. iu both panels IE and IV. the iitial steady-state distribution is strongly peaked not at the source ‘actius. |out at radii where iuigraion stalls (as seen from panel HIE in Fie. 1)).," Note that, in both panels II and IV, the initial steady-state distribution is strongly peaked not at the source radius, but at radii where migration stalls (as seen from panel III in Fig. \ref{fig:4dprop}) )."
+ Iu, In
+"stars) is expected to be about of that of sdB stars, and for SPY this equates to about 9 objects.","stars) is expected to be about of that of sdB stars, and for SPY this equates to about 9 objects."
+ They are likely He-deficient sdO stars., They are likely He-deficient sdO stars.
+" There are 13 He-deficient sdO stars in SPY, and 7 ones (likely evolve from sdB stars) of them with masses less than ~0.5 Μο in our study, similar to the 9 predicted by theory."," There are 13 He-deficient sdO stars in SPY, and 7 ones (likely evolve from sdB stars) of them with masses less than $\sim$ 0.5 $M_\odot$ in our study, similar to the 9 predicted by theory."
+" To more easily compare with theoretical results, we presented the mass distributions of sdO stars obtained from Han et al."," To more easily compare with theoretical results, we presented the mass distributions of sdO stars obtained from Han et al."
+" (2003)? in Fig.5, where the dot-dashed line is from CE ejection and the circle-line is from the merger of two WDs."," $^2$ in Fig.5, where the dot-dashed line is from CE ejection and the circle-line is from the merger of two He-WDs."
+ We see that CE ejection can only account for sdO stars with masses lower than ~0.5 Μο., We see that CE ejection can only account for sdO stars with masses lower than $\sim$ 0.5 $M_\odot$.
+" The He-He merger mainly contributes to high-mass sdO stars, but it is also responsible for a small fraction of low-mass sdO stars."," The He-He merger mainly contributes to high-mass sdO stars, but it is also responsible for a small fraction of low-mass sdO stars."
+" In particular, the fraction of very low-mass stars (« 0.4Μο) is smaller among sdO stars than among sdB stars, as shown in Figs.4 and 5."," In particular, the fraction of very low-mass stars $<
+ 0.4M_\odot$ ) is smaller among sdO stars than among sdB stars, as shown in Figs.4 and 5."
+" This difference probably stems from very few sdB stars with masses lower than 0.4Mo being hotter than 40,000K during their EHB stage."," This difference probably stems from very few sdB stars with masses lower than $0.4M_\odot$ being hotter than 40,000K during their post-EHB stage."
+ We know that a binary resulting from CE ejection generally has a short orbital period., We know that a binary resulting from CE ejection generally has a short orbital period.
+" Thus, if low-mass sdO stars are mainly from the CE ejection channel, then a large fraction of these objects should be observed in short orbital period binaries."," Thus, if low-mass sdO stars are mainly from the CE ejection channel, then a large fraction of these objects should be observed in short orbital period binaries."
+" Among 23 sdO stars with masses less than ~0.5 Mo, (if they indeed evolve from sdB stars), we would expect a fraction of binaries similar to that of sdB stars, i.e. (Napiwotzki et al."," Among 23 sdO stars with masses less than $\sim$ 0.5 $M_\odot$, (if they indeed evolve from sdB stars), we would expect a fraction of binaries similar to that of sdB stars, i.e. (Napiwotzki et al."
+ 2004)., 2004).
+" However, only five sdO stars have been identified in binaries from SPY."," However, only five sdO stars have been identified in binaries from SPY."
+" The possible reasons for the low fraction are (a) the He-He merger contributes a small fraction of low-mass sdO stars, which are all singles; (b)low-mass sdO stars from CE ejection generally have lower mass companions (most likely M, see Fig.15 in Han et al."," The possible reasons for the low fraction are (a) the He-He merger contributes a small fraction of low-mass sdO stars, which are all singles; (b)low-mass sdO stars from CE ejection generally have lower mass companions (most likely M, see Fig.15 in Han et al."
+" 2003), and may be identified as singles; (c) some other observational effects, such as random inclination of orbital planes, contribute to a decrease in detection efficiency."," 2003), and may be identified as singles; (c) some other observational effects, such as random inclination of orbital planes, contribute to a decrease in detection efficiency."
+" Moreover, in comparison to the merger channel, CE ejection, which may leave a thin hydrogen-rich envelope after the ejection, often produces hot subdwarfs with surface helium deficiency."," Moreover, in comparison to the merger channel, CE ejection, which may leave a thin hydrogen-rich envelope after the ejection, often produces hot subdwarfs with surface helium deficiency."
+" Thus, if sdO stars are identified in binaries with short orbital periods, they are most likely to be helium deficient, since the merger of He-He WDs can only produce single helium-enriched sdO stars."," Thus, if sdO stars are identified in binaries with short orbital periods, they are most likely to be helium deficient, since the merger of He-He WDs can only produce single helium-enriched sdO stars."
+" The observations confirm this prediction; i.e., four of the five identified binary sdO stars have short orbital periods and are helium deficient."," The observations confirm this prediction; i.e., four of the five identified binary sdO stars have short orbital periods and are helium deficient."
+" An extreme case from the CE ejection channel is that all of the common envelope has been ejected, leading to a naked He core (hot subdwarfs) and a companion."," An extreme case from the CE ejection channel is that all of the common envelope has been ejected, leading to a naked He core (hot subdwarfs) and a companion."
+ This may produce helium-enriched sdO stars with short orbital periods., This may produce helium-enriched sdO stars with short orbital periods.
+" Since this case is rare, the fraction of helium enriched sdO binaries among He-sdO stars (most of them are from the merger channel) is very low at most, Napiwotzki et al."," Since this case is rare, the fraction of helium enriched sdO binaries among He-sdO stars (most of them are from the merger channel) is very low at most, Napiwotzki et al."
+" 2004, see also Heber 2008)."," 2004, see also Heber 2008)."
+ thank an for his/her valuable us improve the , We thank an anonymous referee for his/her valuable comments that helped us to improve the paper.
+"XZ thanks Dr. Richard WePokorny improving anonymousthe refereeEnglish comments10603013,th", XZ thanks Dr. Richard Pokorny for improving the English language of the original manuscript.
+"at helped to was in paper. by the Natural for China under languageNos.10821061,oftheoriginal manuscript.Thiswork partChinese supportedAcademy of Science under ScienceGrant FoundationNos.06YofQ011001, and the GrantYunnan National Science Foundationand (Grant 2007CB815406,No.08YJ041001)."," This work was in part supported by the Natural Science Foundation of China under Grant Nos.10821061,10603013, and 2007CB815406, the Chinese Academy of Science under Grant Nos.06YQ011001, and the Yunnan National Science Foundation (Grant No.08YJ041001)."
+"at helped to was in paper. by the Natural for China under languageNos.10821061,oftheoriginal manuscript.Thiswork partChinese supportedAcademy of Science under ScienceGrant FoundationNos.06YofQ011001, and the GrantYunnan National Science Foundationand (Grant 2007CB815406,No.08YJ041001).t"," This work was in part supported by the Natural Science Foundation of China under Grant Nos.10821061,10603013, and 2007CB815406, the Chinese Academy of Science under Grant Nos.06YQ011001, and the Yunnan National Science Foundation (Grant No.08YJ041001)."
+"at helped to was in paper. by the Natural for China under languageNos.10821061,oftheoriginal manuscript.Thiswork partChinese supportedAcademy of Science under ScienceGrant FoundationNos.06YofQ011001, and the GrantYunnan National Science Foundationand (Grant 2007CB815406,No.08YJ041001).th"," This work was in part supported by the Natural Science Foundation of China under Grant Nos.10821061,10603013, and 2007CB815406, the Chinese Academy of Science under Grant Nos.06YQ011001, and the Yunnan National Science Foundation (Grant No.08YJ041001)."
+"at helped to was in paper. by the Natural for China under languageNos.10821061,oftheoriginal manuscript.Thiswork partChinese supportedAcademy of Science under ScienceGrant FoundationNos.06YofQ011001, and the GrantYunnan National Science Foundationand (Grant 2007CB815406,No.08YJ041001).the"," This work was in part supported by the Natural Science Foundation of China under Grant Nos.10821061,10603013, and 2007CB815406, the Chinese Academy of Science under Grant Nos.06YQ011001, and the Yunnan National Science Foundation (Grant No.08YJ041001)."
+"Seudder, J. D. 19923. ApJ. 398. 299 Scudder. J. D. 1992b. ApJ. 398. 319 Shebalin. J.. Matthaeus; W.. Montgomery. D. 1983. J. Plasma Phys.","Scudder, J. D. 1992a, ApJ, 398, 299 Scudder, J. D. 1992b, ApJ, 398, 319 Shebalin, J., Matthaeus, W., Montgomery, D. 1983, J. Plasma Phys.,"
+" 29. 525 Tu. C. Y.. Marsch, E.. 1990. 1,"," 29, 525 Tu, C. Y., Marsch, E., 1990, J. Geophys."
+ Res..," Res.,"
+" 95. 4337 Tu. C, Y.. Marsch, E.. 1995, Sp."," 95, 4337 Tu, C. Y., Marsch, E., 1995, Sp."
+"Geophys. Set,", Sci.
+" Rev. 72. 1 Tu. C.-Y.. Marsch, E. 1997, Solar Phys.."," Rev., 73, 1 Tu, C.-Y., Marsch, E. 1997, Solar Phys.,"
+" 171. 363 Tu. C.-Y.. Marsch, E. 20012. J. Geophys."," 171, 363 Tu, C.-Y., Marsch, E. 2001a, J. Geophys."
+" Res.,"," Res.,"
+" 106, 8233 Tu. C.-Y. Marsch. E. 20010, A&AA, 368, 1071 Tu. C.-Y. Pu, Z.-Y.."," 106, 8233 Tu, C.-Y., Marsch, E. 2001b, A, 368, 1071 Tu, C.-Y., Pu, Z.-Y.,"
+"galaxy sample was realized, since it does not change the number of galaxies inside the volume element.","galaxy sample was realized, since it does not change the number of galaxies inside the volume element."
+ Given an assumed cosmology we can then calculate the redshift and the two angular positions for each galaxy., Given an assumed cosmology we can then calculate the redshift and the two angular positions for each galaxy.
+" In particular, we assumed a standard )CCDMcosmologycorrespondingtothe followingseto fcosmologicalparaezautareQ -024.,0.76,,0.04,, 0.73,, sigma,-o74,, )."," In particular, we assumed a standard CDM cosmology corresponding to the following set of cosmological parameters, )."
+" At the end of this procedure we obtained a set of real-space galaxy positions theta,; yforNNeat galaxies."," At the end of this procedure we obtained a set of real-space galaxy positions _n, for galaxies."
+ In the previous section we described how to obtain a set of unperturbed galaxy redshifts given a specific density field., In the previous section we described how to obtain a set of unperturbed galaxy redshifts given a specific density field.
+ Here we will describe how to generate the observed photo-z uncertainties., Here we will describe how to generate the observed photo-z uncertainties.
+" Note, that until now we have not specified a specific redshift likelihood."," Note, that until now we have not specified a specific redshift likelihood."
+" The method presented in this work is general, and does not need to make any specific assumptions about the functional shape of the redshift likelihood."," The method presented in this work is general, and does not need to make any specific assumptions about the functional shape of the redshift likelihood."
+ This is due to our sampling approach which permits us to use any desired redshift likelihood., This is due to our sampling approach which permits us to use any desired redshift likelihood.
+" In this work we will assume that the uncertainty in the observed redshifts is due to the observational strategy, i.e.. due to the instrument."," In this work we will assume that the uncertainty in the observed redshifts is due to the observational strategy, i.e.. due to the instrument."
+" The redshift likelihood is therefore independent of the underlying density field Although it is possible to use an arbitrary redshift likelihood, here we will employ a truncated Gaussian distribution given as where is the photo-z dispersion."," The redshift likelihood is therefore independent of the underlying density field Although it is possible to use an arbitrary redshift likelihood, here we will employ a truncated Gaussian distribution given as where is the photo-z dispersion."
+ This simple distribution permits us to test our method and to estimate its general performance., This simple distribution permits us to test our method and to estimate its general performance.
+" Given the true redshift generated via the procedure described in section 5.1 and by specifying a desired redshift dispersionsigma,,, the observed redshift can be generated by drawing random samples from the truncated Gaussian distribution, given in (11)), via rejection sampling."," Given the true redshift generated via the procedure described in section \ref{MOCK_Observation} and by specifying a desired redshift dispersion, the observed redshift can be generated by drawing random samples from the truncated Gaussian distribution, given in equation \ref{eq:redshift_likelihood_b}) ), via rejection sampling."
+ We distort the redshift of each individual galaxy following this recipe., We distort the redshift of each individual galaxy following this recipe.
+" For the test case considered in this work we treat all galaxies with the same sigma,-9o3.. Since we are able to treat each galaxy individually, there is no problem in specifying a different photo-z likelihood for every galaxy."," For the test case considered in this work we treat all galaxies with the same _z=. Since we are able to treat each galaxy individually, there is no problem in specifying a different photo-z likelihood for every galaxy."
+ The possibility of modeling heterogenous data sets with photo-z pdfs that differ from galaxy to galaxy is very interesting since it allows incorporating detailed uncertainty information from photo-z estimators which could take into account color or morphological information., The possibility of modeling heterogenous data sets with photo-z pdfs that differ from galaxy to galaxy is very interesting since it allows incorporating detailed uncertainty information from photo-z estimators which could take into account color or morphological information.
+" It would also be possible to merge data sets with different redshift accuracies, including spectroscopic surveys."," It would also be possible to merge data sets with different redshift accuracies, including spectroscopic surveys."
+ We will explore these promising ideas in future work., We will explore these promising ideas in future work.
+ The process of generating a galaxy distribution with redshift distortions corresponding to the likelihood given in equation (11)) is visualized in figure 2.., The process of generating a galaxy distribution with redshift distortions corresponding to the likelihood given in equation \ref{eq:redshift_likelihood_b}) ) is visualized in figure \ref{fig:mock}.
+" It can be seen that given a redshift dispersion of sigma,-oos Structures along the line of sight are smeared out on a scale of roughly Mpc.", It can be seen that given a redshift dispersion of _z= structures along the line of sight are smeared out on a scale of roughly Mpc.
+ In the following we will apply our method to this mock survey., In the following we will apply our method to this mock survey.
+ In this section we will apply our method to our mock survey in order to estimate its performance., In this section we will apply our method to our mock survey in order to estimate its performance.
+" In particular, we focus on the convergence behavior of the sampler which determines the efficiency of the method in a realistic setting."," In particular, we focus on the convergence behavior of the sampler which determines the efficiency of the method in a realistic setting."
+ The Metropolis Hastings Sampler is designed to have the target, The Metropolis Hastings Sampler is designed to have the target
+We analyzed all the light curves using the program by Djurasevicetal.(1992a).. generalized for the case of an overcontact configuration (Djurasevicetal.1998).,"We analyzed all the light curves using the program by \citet{djur92a}, generalized for the case of an overcontact configuration \citep{djur98}."
+. The program is based on the Roche model and the principles described by Wilson&Devinney(1971)., The program is based on the Roche model and the principles described by \citet{wilanddev71}.
+. The system parameters are estimated by applying an inverse-problem method (Djurasevicetal.1992b).. based on the modified Marquardt(1963) algorithm.," The system parameters are estimated by applying an inverse-problem method \citep{djur92b}, based on the modified \citet{marquardt} algorithm."
+ The underlying binary system model allows various system configurations from detached to overcontact. including active spot regions on the components.," The underlying binary system model allows various system configurations from detached to overcontact, including active spot regions on the components."
+ The latest version of the model and the solving procedure is described in detail by Djurasevieetal.(20056)., The latest version of the model and the solving procedure is described in detail by \citet{djura08}.
+ The fitting was done for a limited set of parameters. while the following parameters were fixed based on information from independent analysis of the radial velocity curves and plausible assumptions about the systems.," The fitting was done for a limited set of parameters, while the following parameters were fixed based on information from independent analysis of the radial velocity curves and plausible assumptions about the systems."
+ The limb darkening in the latest version of the model follows the nonlinear approximation given by Claret(2000) and the limb-darkening coefficients for the corresponding passbands were interpolated from Claret's tables according to the current values of Tey and log(g) in each iteration., The limb darkening in the latest version of the model follows the nonlinear approximation given by \citet{claret00} and the limb-darkening coefficients for the corresponding passbands were interpolated from Claret's tables according to the current values of ${\rm T_{eff}}$ and $\log(g)$ in each iteration.
+" The results of the analysis are given for each star in a table (see Tables 5.. 6.. 7.. and 8)). where # is the total number of the B. V. and R observations; X(O-Cy - the final sum of squares of residuals between observed (LCO) and synthetic (LCC) light-curves: c - the standard deviation of the residuals: g=ΠΟΠΗ - the mass ratio of the components: Ty. - the temperatures of the hotter and the cooler component: Phe Anc. fno - the gravity-darkening exponents. albedos and nonsynchronous rotation coefficients of the components: A,. 4. A. and q, - the temperature coefficients. angular dimensions. longitudes. and latitudes (1n arc degrees) of the spots (1f present in the model): Εις - the filling factors for the critical Roche lobe of the hotter and cooler component: i[] - the orbit inclination (in are degrees): Li/(LiΕις+L3) - the third light contribution: Qi. Q;, and Q,,, - the dimensionless surface potentials of the system components and of the inner and outer contact surfaces: fowl] - the degree of overcontact: Ry. - the polar radit of the components in units of separation: Li/(Li+Lo - the luminosity of the hotter star: Nl,[M&S]. ινΓι] - the stellar masses and mean radi of the components m solar units: loggp. - the logarithm (base 10) of the system components effective gravity: Mi - the absolute bolometric magnitudes of the components: and ej;4[R..] - the orbital semi-major axis in units of solar radius."," The results of the analysis are given for each star in a table (see Tables \ref{TabQXAnd}, \ref{TabRWCom}, \ref{TabMRDel}, and \ref{TabBD}) ), where $n$ is the total number of the B, V, and R observations; ${\rm \Sigma (O-C)^2}$ - the final sum of squares of residuals between observed (LCO) and synthetic (LCC) light-curves; ${\rm \sigma}$ - the standard deviation of the residuals; $q=m{\rm _c}/m{\rm _h}$ - the mass ratio of the components; ${\rm T_{\rm h,c}}$ - the temperatures of the hotter and the cooler component; $\beta_{\rm h,c}$, ${\rm A_{\rm h,c}}$, $f_{\rm h,c}$ - the gravity-darkening exponents, albedos and nonsynchronous rotation coefficients of the components; ${\rm A_{s}}$, ${\rm \theta_{s}}$, ${\rm \lambda_{s}}$, and ${\rm \varphi_{s}}$ - the temperature coefficients, angular dimensions, longitudes, and latitudes (in arc degrees) of the spots (if present in the model); ${\rm F_{h,c}}$ - the filling factors for the critical Roche lobe of the hotter and cooler component; $i \ [^\circ]$ - the orbit inclination (in arc degrees); ${\rm L_3/(L_h+L_c+L_3)}$ - the third light contribution; $\Omega_{\rm h,c}$, $\Omega_{in}$ and $\Omega_{out}$ - the dimensionless surface potentials of the system components and of the inner and outer contact surfaces; $f{\rm _{over}}[\%]$ - the degree of overcontact; ${\rm R_{h,c}}$ - the polar radii of the components in units of separation; ${\rm L_h/(L_h+L_c)}$ - the luminosity of the hotter star; $\cal M_{\rm h,c} {\rm [M_{\odot}]}$, $\cal R_{\rm h,c} {\rm [R_{\odot}]}$ - the stellar masses and mean radii of the components in solar units; ${\rm log} \ g_{\rm h,c}$ - the logarithm (base 10) of the system components effective gravity; $M^{\rm h,c}_{\rm bol}$ - the absolute bolometric magnitudes of the components; and $a_{\rm orb} {\rm [R_{\odot}]}$ - the orbital semi-major axis in units of solar radius."
+ Although the mass ratio for each system was fixed when solving the inverse problem to the value adopted from the relevant spectroscopic study. the influence of the mass ratio," Although the mass ratio for each system was fixed when solving the inverse problem to the value adopted from the relevant spectroscopic study, the influence of the mass ratio"
+ellicient Goldreich Nicholson viscosity.,efficient Goldreich Nicholson viscosity.
+ While this gives a firmer theoretical foundation for the quadratic prescription it does not help with the observational problem of insufficient dissipation in the case of tides and stellar pulsations., While this gives a firmer theoretical foundation for the quadratic prescription it does not help with the observational problem of insufficient dissipation in the case of tides and stellar pulsations.
+ A possible resolution of this problem is suggested by the fact that the successful applications of the (wo prescriptions correspond to very dillerent perturbation periods., A possible resolution of this problem is suggested by the fact that the successful applications of the two prescriptions correspond to very different perturbation periods.
+ The Zahn(1966.1989) scaling seems to work well lor periods of order days. and the Goldveich and collaborators quadratic scaling seems (o apply for periods of order minutes.," The \citet{Zahn_66, Zahn_89} scaling seems to work well for periods of order days, and the Goldreich and collaborators quadratic scaling seems to apply for periods of order minutes."
+ This distinction is important because. in stars wilh surface convection. Ixolmogorov scaling predicts (that the eddies with turnover times of several minutes would have (vpical sizes that are very small compared to the local pressure scale height and any other external length scales.," This distinction is important because, in stars with surface convection, Kolmogorov scaling predicts that the eddies with turnover times of several minutes would have typical sizes that are very small compared to the local pressure scale height and any other external length scales."
+ On (he other hand turnover times of days correspond to eddies wilh typical sizes comparable or larger than the local pressure scale height., On the other hand turnover times of days correspond to eddies with typical sizes comparable or larger than the local pressure scale height.
+ In this case Kolmogorov scaling is not expected to apply., In this case Kolmogorov scaling is not expected to apply.
+ The flow seen in 2D and 3D simulations of stellar convection is very different [rom a Kolmogorov cascade (Solia&Chan1984:SteinNorcllund1939:Malagolietal.1990).," The flow seen in 2D and 3D simulations of stellar convection is very different from a Kolmogorov cascade \citep{Sofia_Chan_84,
+Stein_Nordlund_89, Malagoli_Cattaneo_Brummell_90}."
+. There are (wo important disünctons., There are two important distinctions.
+ The first is that the velocity power spectirunm is much flatter in the simulations than Ixolmogorov. aud so one expects to find a slower loss of dissipation efficiency wilh increased. [requency. of the external shear. assuming (hat (he dissipation is dominated by (he resonant edcdies. as long as the external shear has a period Chat corresponds to eddy (πόνος times too lone to fall in the inertial subrange of ]xolmogorov. turbulence.," The first is that the velocity power spectrum is much flatter in the simulations than Kolmogorov, and so one expects to find a slower loss of dissipation efficiency with increased frequency of the external shear, assuming that the dissipation is dominated by the resonant eddies, as long as the external shear has a period that corresponds to eddy turnover times too long to fall in the inertial subrange of Kolmogorov turbulence."
+ The second is that the flow is no longer isotropic ancl hence the effeclive viscosity should be a Censor. rather than a scalar.," The second is that the flow is no longer isotropic and hence the effective viscosity should be a tensor, rather than a scalar."
+ As a first attempt to explore this possibility. Penevοἱal.(2007.200580). adapted the Goodman&Oh(1997) perturbative caleulation to the Robinsonetal.(2003) numerical model of stellar convection and found an asymmetric effective viscosity that scaled. linearly," As a first attempt to explore this possibility \citet{Penev_Sasselov_Robinson_Demarque_07, Penev_Sasselov_Robinson_Demarque_08}
+ adapted the \citet{Goodman_Oh_97} perturbative calculation to the \citet{Robinson_et_al_03} numerical model of stellar convection and found an asymmetric effective viscosity that scaled linearly"
+"of optical companions from Poisson statisties is 36νο, and 36 for the DIRG. and C'S. respectively.","of optical companions from Poisson statistics is 36, and 36 for the BIRG, and CS, respectively."
+ If optical companions are subtracted. [ppys lis 2JC. and for the DIRG ancl the CS. respectively.," If optical companions are subtracted, $f_{phys}$ is $\approx$, and for the BIRG and the CS, respectively."
+ These results show that there is not a significant excess of companions between Dright Ηνο galaxies with respect to non-active galaxies. if all companion galaxies with 5ApeSDeZLOA aare taken into account.," These results show that there is not a significant excess of companions between Bright IRAS galaxies with respect to non-active galaxies, if all companion galaxies with $5 Kpc \simlt D_C \simlt 10 Kpc$ are taken into account."
+" However. this result should be viewed with caution since fj,>ρω."," However, this result should be viewed with caution since $f_{opt} \gg f_{phys}$."
+ A statistical approach is not appropriate in this companion size range., A statistical approach is not appropriate in this companion size range.
+ Ànv inter-sample dillerence can be proved as significant only if [ως iis estimated from redshift measurements lor all companion galaxies., Any inter-sample difference can be proved as significant only if $f_{phys}$ is estimated from redshift measurements for all companion galaxies.
+ OL 87 BIRG galaxies. 22 has at least. one companion of diameter D>104Ape wwithin a search radius 3 Dy. against £29 of the 90 objects of the CS.," Of 87 BIRG galaxies, $\approx$ has at least one companion of diameter $D \geq
+10Kpc$ within a search radius 3 $D_S$, against $\approx$ 29 of the 90 objects of the CS."
+ The expected number of optical companions from Poisson statistics is20%. and for the BIRG and CS respectively.," The expected number of optical companions from Poisson statistics is, and for the BIRG and CS respectively."
+" ο iis subtracted. fj, i8 2 and for the BIRG and the CS. respectively."," If $f_{opt}$ is subtracted, $f_{phys}$ is $\approx$ and for the BIRG and the CS, respectively."
+ These results show an excess of large companions (De>10 pe) in the Bright IRAS galaxies with respect to non-active galaxies., These results show an excess of large companions $D_C \geq 10 Kpc$ ) in the Bright IRAS galaxies with respect to non-active galaxies.
+ A X? test gives a confidence level for this result of99., A $\chi^2$ test gives a confidence level for this result of.
+9%. The search radiuscases was taken as 250 Ixpc of projected linear distance. bevond which we assumed a “non detection.," The search radius was taken as 250 Kpc of projected linear distance, beyond which we assumed a “non detection."""
+ The left hand side of Fig., The left hand side of Fig.
+ 2. presents three panels with the cumulative distribution of (he nearest companion (without. correction [or optical companions) up to a projected linear distance (d) of 140 Ixpc., \ref{fig02} presents three panels with the cumulative distribution of the nearest companion (without correction for optical companions) up to a projected linear distance $d_p$ ) of 140 Kpc.
+" The upper panel shows the cumulative distribution of companions with diameter in (he range 5ApezDz,10Ape.companions."," The upper panel shows the cumulative distribution of companions with diameter in the range $5 Kpc \simgt D \simgt 10 Kpc$,."
+ The middle panel shows the cumulative distribution of companions with diameter D>I0Apc. and the lower panel shows the same distribution for companions with D> 20 pc.," The middle panel shows the cumulative distribution of companions with diameter $D\geq 10Kpc$, and the lower panel shows the same distribution for companions with $D\geq$ 20 Kpc."
+" The error bars on the control samples Lequencies were sel wilh a 7""bootstrap technique (ElronanclTibshirani.1993) bw randomly re-sampline 1e control sample galaxies into a large number (3000) of pseucdo-control samples (i. e.. we »ult 3000 pseudo-control samples of 90 randomlv-selected. galaxies)."," The error bars on the control samples frequencies were set with a “bootstrap"" technique \citep{boot} by randomly re-sampling the control sample galaxies into a large number (3000) of pseudo-control samples (i. e., we built 3000 pseudo-control samples of 90 randomly-selected galaxies)."
+ The uncertaintv on je companion frequency was set as equal (o twice (he standard deviation measured [rom 1e. distribution of 3000 companion frequencies computed for each pseudo control sample., The uncertainty on the companion frequency was set as equal to twice the standard deviation measured from the distribution of 3000 companion frequencies computed for each pseudo control sample.
+ ων-omparing the environments of BIRG galaxies aud control sample galaxies. it is found that jere is a statistically signilicant excess of bright companions (DeZ104v pe) in the inlrared Mnitters.," Comparing the environments of BIRG galaxies and control sample galaxies, it is found that there is a statistically significant excess of bright companions $D_C \simgt
+10 Kpc$ ) in the infrared emitters."
+ For companion diameters 5ApexDe<LOA the samples show no significant, For companion diameters $5 Kpc \leq D_C \leq 10 Kpc$ the samples show no significant
+data.,data.
+ The 0.000173 Iz frequency is a beat. frequency with the orbital period of the satellite (96 minutes): the 0.00519 Ην frequency is its μήνα harmonic., The 0.000173 Hz frequency is a beat frequency with the orbital period of the satellite (96 minutes); the 0.00519 Hz frequency is its third harmonic.
+ In a simulation of the dataset we have reproduced this behavior., In a simulation of the dataset we have reproduced this behavior.
+ The pulse profile of the PSPC observation is shown in Figure 4., The pulse profile of the PSPC observation is shown in Figure 4.
+ There is no evidence in the Ift for harmonics higher than the fundamental., There is no evidence in the fft for harmonics higher than the fundamental.
+ The pulsed fraction is 3622556., The pulsed fraction is $\pm$.
+.. This disagrees with the value derived from the CILANDIA detection., This disagrees with the value derived from the CHANDRA detection.
+ This discrepancy can be accounted [or by the different spectral response of the (wo detectors and a difference between the unpulsed ancl pulsed spectrum., This discrepancy can be accounted for by the different spectral response of the two detectors and a difference between the unpulsed and pulsed spectrum.
+ Relative to the ACIS-I detector the PSPC weights lower energv photons significantly more than higher energy. photons., Relative to the ACIS-I detector the PSPC weights lower energy photons significantly more than higher energy photons.
+ For example. the ratio of effective area for the PSPC at 0.8 keV to its effective area al 2.0 keV is ~ 3.," For example, the ratio of effective area for the PSPC at 0.8 keV to its effective area at 2.0 keV is $\sim$ 3."
+ The same ratio lor the ACIS-I detector is ~0.8., The same ratio for the ACIS-I detector is $\sim$ 0.8.
+ Further we lind that the pulsed source spectrum is somewhat solter (han ils unpulsed emission., Further we find that the pulsed source spectrum is somewhat softer than its unpulsed emission.
+ Therelore because of the PSPC's enerev response il will detect the pulsed part of source emission more efficiently than the unpulsed part., Therefore because of the PSPC's energy response it will detect the pulsed part of source emission more efficiently than the unpulsed part.
+ This qualitatively accounts lor the higher pulsed fraction as determined by the PSPC., This qualitatively accounts for the higher pulsed fraction as determined by the PSPC.
+ We have calculated the sensitivitv of the remaining non-CILANDRA observations using this value of the pulsed fraction., We have calculated the sensitivity of the remaining non-CHANDRA observations using this value of the pulsed fraction.
+ We find that none of the other non-CILANDRA observations listed in Table 1 would be sensitive to this pulsation., We find that none of the other non-CHANDRA observations listed in Table 1 would be sensitive to this pulsation.
+ We have searched [or a possible radio/optical counterpart to CNOU 121134 and find no radio counterpart., We have searched for a possible radio/optical counterpart to CXOU J0110043.1-721134 and find no radio counterpart.
+ There is an 13th magnitude star (801020208965 in the Guide Star Catalogue-II available on-line through /cal?]/2714ssRAI99.3) which is located 1.2 arc seconds [rom the source., There is an 18th magnitude star (S01020208968 in the Guide Star Catalogue-II available on-line through sRM99.3) which is located 1.2 arc seconds from the source.
+ The position of CNOU JOLLO0L3.1-721134 has an uncertainlv of 0.6 are seconds. therefore it is unlikely to be associated with this star.," The position of CXOU J0110043.1-721134 has an uncertainly of 0.6 arc seconds, therefore it is unlikely to be associated with this star."
+ Nevertheless it remains a possibility pending refinement of the X-ray source position., Nevertheless it remains a possibility pending refinement of the X-ray source position.
+ From the density. of stars in the Guide Star Catalogue within 30 arc, From the density of stars in the Guide Star Catalogue within 30 arc
+ 3.10).,$3.1\sigma$ ).
+ Moreover. the rise at the end of the curve between ~S.0Sl has τς2.741.0 hrs (3.30).," Moreover, the rise at the end of the curve between $\sim 8.0-8.1$ has $\tau = 2.7\pm 1.0$ hrs $3.3\sigma$ )."
+ Lt is possible to estimate the minimum Doppler factor (eg. Dondi Chisellini 1995. Mattox et al.1," It is possible to estimate the minimum Doppler factor $\delta$ (e.g. Dondi Ghisellini 1995, Mattox et al.,"
+997)7.. although during this pointed observation it was not possible to perform multiwavelength observations. because the blazar was too close to the Sun.," although during this pointed observation it was not possible to perform multiwavelength observations, because the blazar was too close to the Sun."
+ Therefore. we adopted. for the X-rav [lux and spectrum. the values measured. during the December 2009 outburst. when 3€ 454.8 reached: similar ~ rav fluxes: a=0.4 and £Agv~30 μον (Bonnoli ct al.," Therefore, we adopted for the X-ray flux and spectrum, the values measured during the December 2009 outburst, when 3C 454.3 reached similar $\gamma-$ ray fluxes: $\alpha=0.4$ and $F_{\rm 1~keV}\sim 30$ $\mu$ Jy (Bonnoli et al."
+ 2010)., 2010).
+ We obtain ὃ2:14. for 1 GeV 5 rays and 7=2.7 hours. a high value for a lower limit. but not unlikely (0f Chisellini et al.," We obtain $\delta \geq 14$, for 1 GeV $\gamma$ rays and $\tau = 2.7$ hours, a high value for a lower limit, but not unlikely (cf Ghisellini et al."
+ 2010)., 2010).
+ We noted a few cases of sudden drop in the flux with ~30 min time scale on ~6.16.2 and ~7.8 April., We noted a few cases of sudden drop in the flux with $\sim 30$ min time scale on $\sim 6.1-6.2$ and $\sim 7.3$ April.
+ A closer inspection of the data revealed that these bins had small source on time (~25% of the whole bin)., A closer inspection of the data revealed that these bins had small source on time $\sim 25$ of the whole bin).
+ Therefore. it is likely that it is a fake drop. caused by the not. good reconstruction of the [lux in presence of too few events. although sullicient for a high 7's.," Therefore, it is likely that it is a fake drop, caused by the not good reconstruction of the flux in presence of too few events, although sufficient for a high $TS$."
+ We presented. the variability. analysis of the 5 rav observation with fermi/fLJXE. of the blazar 3C 454.3 pointed during the outburst of carly April 2010., We presented the variability analysis of the $\gamma-$ ray observation with /LAT of the blazar 3C 454.3 pointed during the outburst of early April 2010.
+ Contrary to the usual scanning mode. this time performed a 200 ks pointed observation (Wallace ct al.," Contrary to the usual scanning mode, this time performed a $200$ ks pointed observation (Wallace et al."
+ 2010). which euarantees a higher sensitivity and a longer source ontinic.," 2010), which guarantees a higher sensitivity and a longer source ontime."
+ These particular settings. together with the exceptionally. high Lux emitted. by the blazar. enabled us to probe the shortest time scales that we found to be in the range of ~25 hours. confirming the indications found by Bonnoli et al. (," These particular settings, together with the exceptionally high flux emitted by the blazar, enabled us to probe the shortest time scales that we found to be in the range of $\sim 2-5$ hours, confirming the indications found by Bonnoli et al. ("
+2010). Tavecchio οἱ al. (,"2010), Tavecchio et al. ("
+2010) and in agreement with the current theories.,2010) and in agreement with the current theories.
+ Although the significance of the events is not very strong ( 30). the fact that it is present in the light curves with cillerent time bins is an indication that the findings are well grounded.," Although the significance of the events is not very strong $\sim 3\sigma$ ), the fact that it is present in the light curves with different time bins is an indication that the findings are well grounded."
+ On the other side. we have not. found significant evidence of sub-hour time scales. but the present data set is not vet conclusive.," On the other side, we have not found significant evidence of sub-hour time scales, but the present data set is not yet conclusive."
+ We noted that the [Lux increased at the end of the pointed observations with 72.7 hours., We noted that the flux increased at the end of the pointed observations with $\tau\sim 2.7$ hours.
+ Fherefore. we could expect (or reasonably hope) that hours or less time scales may. be effectively. probed with ΕΕ under conditions of even greater fluxes. likely as occurred in December 2009 (Bonnoli et al.," Therefore, we could expect (or reasonably hope) that hours or less time scales may be effectively probed with /LAT under conditions of even greater fluxes, likely as occurred in December 2009 (Bonnoli et al."
+ 2010)., 2010).
+ This research has mace use of cata obtained. from. the Ligh Enerev Astrophysics Science Archive Research Center (IIEASARC). provided by NASA's Goddard Space Flight Center.," This research has made use of data obtained from the High Energy Astrophysics Science Archive Research Center (HEASARC), provided by NASA's Goddard Space Flight Center."
+ We are grateful to the /LAT Collaboration for having performed this pointed observation and mace the data publicly available immediately., We are grateful to the /LAT Collaboration for having performed this pointed observation and made the data publicly available immediately.
+ This work has been partially supported. by PRIN-ALUR 2007 and ASL Cant 1405540600., This work has been partially supported by PRIN-MiUR 2007 and ASI Grant I/088/06/0.
+where it is assumed that the disk is axisymmetric and has zero vertical thickness. and X») denotes the surface mass density distribution.,"where it is assumed that the disk is axisymmetric and has zero vertical thickness, and $\Sigma(r)$ denotes the surface mass density distribution."
+ Doing the angular integral gives In the limit >0. the energy change is which agrees with the corresponding expression in Binney&Tremaine(1987) using the impulse approximation. as expected. while in the limit5x for the prograde case and for a retrograde encounter To evaluate these. or eq. (34)).," Doing the angular integral gives In the limit $\alpha \rightarrow 0$, the energy change is which agrees with the corresponding expression in \citet{BT87} using the impulse approximation, as expected, while in the limit $\alpha \rightarrow \infty$ for the prograde case and for a retrograde encounter To evaluate these, or eq. \ref{Ener_sl}) ),"
+ it is necessary to specify the radial dependence of ο). Le.. the shape of the rotation curve of the victim.," it is necessary to specify the radial dependence of $\alpha(r)$, i.e., the shape of the rotation curve of the victim."
+ We have also evaluated the next order correction term given by equation (10)). vielding: the upper and lower signs terms refer to prograde and retrograde encounters. respectively. and Av. = 0because the encounter is assumed to be coplanar.," We have also evaluated the next order correction term given by equation \ref{correction}) ), yielding: where the upper and lower signs terms refer to prograde and retrograde encounters, respectively, and $\Delta$ $_{z} = 0$ because the encounter is assumed to be coplanar."
+ We note that the /7 term does not couple to the phase o of an individual star. because it describes the overall reaction of the victim to the perturber: 1.8. the gain of angular momentum of the entire disk.," We note that the $H$ term does not couple to the phase $\phi_0$ of an individual star, because it describes the overall reaction of the victim to the perturber; i.e. the gain of angular momentum of the entire disk."
+ On top of this global contribution. each star also receives a resonant sensitive and ay dependent contribution.," On top of this global contribution, each star also receives a resonant sensitive and $\phi_0$ dependent contribution."
+ A similar effect occurs in linear order., A similar effect occurs in linear order.
+ There. the overall contribution leads to a motion of the center of mass of the victim.," There, the overall contribution leads to a motion of the center of mass of the victim."
+ But since our caleulation is performed in the center of mass reference frame. this o independent contribution does not show up explicitly in the linear order results.," But since our calculation is performed in the center of mass reference frame, this $\phi_0$ independent contribution does not show up explicitly in the linear order results."
+ There all terms depend on the phase oo., There all terms depend on the phase $\phi_0$.
+" We further note that the argument of the trigonometric functions involving o is different in first and second order: o in linear order and 20, in second order.", We further note that the argument of the trigonometric functions involving $\phi_0$ is different in first and second order: $\phi_0$ in linear order and $2\phi_0$ in second order.
+ This implies. that the coupling of both orders leads to an asymmetric perturbation in . and y. whereas the linear term alone always produces symmetric results as can be seen from the velocity increments in the .c and y directions.," This implies, that the coupling of both orders leads to an asymmetric perturbation in $x$ and $y$, whereas the linear term alone always produces symmetric results as can be seen from the velocity increments in the $x$ and $y$ directions."
+ Next. consider an encounter from a parabolic trajectory. which ts the case analyzed by Press&Teukolsky(1977) in their study of tidal interactions between stars.," Next, consider an encounter from a parabolic trajectory, which is the case analyzed by \citet{PT77} in their study of tidal interactions between stars."
+ This situation is relevant for collisions between galaxies in the field or in loose groups. where the orbits are highly elongated.," This situation is relevant for collisions between galaxies in the field or in loose groups, where the orbits are highly elongated."
+ For this reason. TT72 focused on this geometry in particular. as in the examples shown in Figure l.," For this reason, TT72 focused on this geometry in particular, as in the examples shown in Figure 1."
+ The relative velocity between the victim and perturber in this case is their mutual escape velocity. treating the interaction as that between two point masses. and Is given by: At the minimum separation between the victim and the perturber the relative velocity attains its maximum value Vy: and the relative velocity can be written as We orient the disk as 1n the earlier derivation. so the orbits within the victim disk are given again by eqs. (2))," The relative velocity between the victim and perturber in this case is their mutual escape velocity, treating the interaction as that between two point masses, and is given by: At the minimum separation between the victim and the perturber the relative velocity attains its maximum value $_{0}$: and the relative velocity can be written as We orient the disk as in the earlier derivation, so the orbits within the victim disk are given again by eqs. \ref{vicorb}) )"
+ and (12))., and \ref{phaset}) ).
+ To specify the orbit path. employ polar coordinates and write," To specify the orbit path, employ polar coordinates and write"
+one snapshot to another.,one snapshot to another.
+ The map at 12900 us TiO free and there are mostly CN lines: in this case. the surface intensity contrast is less strong than in the ΠΟ bands.," The map at 12900 is TiO free and there are mostly CN lines: in this case, the surface intensity contrast is less strong than in the TiO bands."
+ Fig., Fig.
+ 8 shows the comparison to the data taken in 2004., \ref{young_comparison2} shows the comparison to the data taken in 2004.
+ There. the filters used span wavelength regions corresponding to TiO absorption bands with different strengths centered at and (transitions A—XA(y) and ΕΥΠ-- X2A(e). see Fig. 1)).," There, the filters used span wavelength regions corresponding to TiO absorption bands with different strengths centered at and (transitions ${\rm A}^3\Phi - {\rm X}^3\Delta (\gamma)$ and ${\rm E}^3\Pi - {\rm X}^3\Delta (\epsilon)$ , see Fig. \ref{filters}) )."
+ Again. the same snapshot fitted the whole dataset from the same epoch.," Again, the same snapshot fitted the whole dataset from the same epoch."
+ The departure from circular symmetry is more evident than in the H band. with the first and the second lobes already showing large visibility fluctuations.," The departure from circular symmetry is more evident than in the H band, with the first and the second lobes already showing large visibility fluctuations."
+ The RHD simulation shows an excellent agreement with the data both in the visibility curves and closure phases., The RHD simulation shows an excellent agreement with the data both in the visibility curves and closure phases.
+ However. within the same observation epoch we had to scale the size of the simulated star to a different apparent diameter at each observed wavelength.," However, within the same observation epoch we had to scale the size of the simulated star to a different apparent diameter at each observed wavelength."
+ For example. in Fig.," For example, in Fig."
+ 8 the apparent diameter varies from 47.3 to 52.6 mas., \ref{young_comparison2} the apparent diameter varies from 47.3 to 52.6 mas.
+ Thus we infer that our RHD simulation fails to reproduce the TiO molecular band strengths probed by the three filters., Thus we infer that our RHD simulation fails to reproduce the TiO molecular band strengths probed by the three filters.
+ As already pointed out in Paper L. our RHD simulations are constrained by execution time and they use a grey approximation for the radiative transfer that 1s. well justified in the stellar interior and is a crude approximation in the optically thin layers: as a consequence the thermal gradient is too shallow and weakens the contrast between strong anc weak lines (2)..," As already pointed out in Paper I, our RHD simulations are constrained by execution time and they use a grey approximation for the radiative transfer that is well justified in the stellar interior and is a crude approximation in the optically thin layers; as a consequence the thermal gradient is too shallow and weakens the contrast between strong and weak lines \citep{2006sf2a.conf..455C}."
+ The intensity maps look too sharp with respect to the observations., The intensity maps look too sharp with respect to the observations.
+ The implementation of non-grey opacities with five wavelength bins employed to describe the wavelength dependence of radiation. fields (see??.fordetails) should change the mean temperature structure and the temperature fluctuations.," The implementation of non-grey opacities with five wavelength bins employed to describe the wavelength dependence of radiation fields \citep[see][for
+ details]{1994A&A...284..105L, 1982A&A...107....1N} should change the mean temperature structure and the temperature fluctuations."
+ The mean thermal gradient in the outer layers. where TiO absorption has a large effect. should increase.," The mean thermal gradient in the outer layers, where TiO absorption has a large effect, should increase."
+ Moreover. in a next step. the inclusion of the radiatior pressure in. the simulations should lead to a different density/pressure structure with a less steep decline of density with radius.," Moreover, in a next step, the inclusion of the radiation pressure in the simulations should lead to a different density/pressure structure with a less steep decline of density with radius."
+ We expect that the intensity maps probing TiO bands with different strengths will eventually show larger diameter variations due to the molecular absorption as a result of these refinements., We expect that the intensity maps probing TiO bands with different strengths will eventually show larger diameter variations due to the molecular absorption as a result of these refinements.
+ In ?.. the authors managed to model the data in the 7000 ffilter with two best-fitting parametric models. consisting of a circular disk with superimposed bright features (Fig. 9..," In \cite{2000MNRAS.315..635Y}, the authors managed to model the data in the 7000 filter with two best-fitting parametric models, consisting of a circular disk with superimposed bright features (Fig. \ref{images_2},"
+ central panel) or dark features (left panel)., central panel) or dark features (left panel).
+ We compared these parametric models to our best-fitting synthetic image of Fig., We compared these parametric models to our best-fitting synthetic image of Fig.
+ 7 (top row)., \ref{young_comparison1} (top row).
+ The convolved image. displayed in Fig.," The convolved image, displayed in Fig."
+ 9 (right panel). shows a better qualitative agreement with the bright features parametric model.," \ref{images_2} (right panel), shows a better qualitative agreement with the bright features parametric model."
+ In fact. 3D simulations show that the surface contrast is enhanced by the presence of significant molecular absorbers like TiO which contribute in layers where waves and shocks start to dominate.," In fact, 3D simulations show that the surface contrast is enhanced by the presence of significant molecular absorbers like TiO which contribute in layers where waves and shocks start to dominate."
+ The location of bright spots Is then a consequence of the underlying activity., The location of bright spots is then a consequence of the underlying activity.
+ We have used radiation hydrodynamics simulations of red supergiant stars to explain interferometric observations of Betelgeuse from the optical to the infrared region., We have used radiation hydrodynamics simulations of red supergiant stars to explain interferometric observations of Betelgeuse from the optical to the infrared region.
+ The picture of the surface of Betelgeuse resulting from this work ts the following: (1) a granulation pattern is undoubtedly present on the surface and the convection-related structures have a strong signature in the visibility curve and closure phases at high spatial frequencies in the H band and on the first and second lobes in the optical region. (, The picture of the surface of Betelgeuse resulting from this work is the following: (i) a granulation pattern is undoubtedly present on the surface and the convection-related structures have a strong signature in the visibility curve and closure phases at high spatial frequencies in the H band and on the first and second lobes in the optical region. (
+1) In the H band. Betelgeuse is characterized by a granulation pattern that is composed of convection-related structures of different sizes.,"ii) In the H band, Betelgeuse is characterized by a granulation pattern that is composed of convection-related structures of different sizes."
+ There are small to medium scale. granules (5-15 mas) and a large convective cell (520 mas)., There are small to medium scale granules (5–15 mas) and a large convective cell $\approx$ 30 mas).
+ This supports previous detections carried out with our RHD simulation in the K band (Paper I). using parametric models and the same dataset (?).. ? with VLT/NACO observations and ο with VLTI/AMBER observations.," This supports previous detections carried out with our RHD simulation in the K band (Paper I), using parametric models and the same dataset \citep{2009A&A...508..923H}, \cite{2009A&A...504..115K} with VLT/NACO observations and \cite{2009A&A...503..183O} with VLTI/AMBER observations."
+ Moreover. we have demonstrated that H»O molecules contribute more than CO and CN to the position of the visibility curve’s first null (and thus to the measured stellar radius) and to the small scale surface structures. (," Moreover, we have demonstrated that $_2$ O molecules contribute more than CO and CN to the position of the visibility curve's first null (and thus to the measured stellar radius) and to the small scale surface structures. ("
+111) In the optical. Betelgeuse's surface appears more complex with areas up to 50 times brighter than the dark ones.,"iii) In the optical, Betelgeuse's surface appears more complex with areas up to 50 times brighter than the dark ones."
+ This picture is the consequence of the underlaying activity characterized by interactions between shock waves and non-radial pulsations in layers where there are strong TiO molecular bands., This picture is the consequence of the underlaying activity characterized by interactions between shock waves and non-radial pulsations in layers where there are strong TiO molecular bands.
+ These observations provide a wealth of information about both the stars and our RHD models., These observations provide a wealth of information about both the stars and our RHD models.
+ The comparison with the observations in the TiO bands allowed us to suggest which approximations must be replaced with more realistic treatments in the simulations., The comparison with the observations in the TiO bands allowed us to suggest which approximations must be replaced with more realistic treatments in the simulations.
+ New models with wavelength resolution (1.e.. non-grey opacities) are in progress and they will be tested against these observations.," New models with wavelength resolution (i.e., non-grey opacities) are in progress and they will be tested against these observations."
+ From the observational. point of view. further multi-epoch observations. both in the optical andin the infrared. are needed to assess the time variability of convection.," From the observational point of view, further multi-epoch observations, both in the optical andin the infrared, are needed to assess the time variability of convection."
+refvariseetion)).,).
+ Therefore. the velocity ranges listed in are the widest. velocity. spread: over. which emission has been detected. considering all the observations available.," Therefore, the velocity ranges listed in \\ref{resotable} are the widest velocity spread over which emission has been detected, considering all the observations available."
+ EVhis is the same convention as used for the nisers in the Calactic centre (2)..., ÊThis is the same convention as used for the masers in the Galactic centre \citep{caswell10mmb1}.
+ The mean. spread in velocitv for ἃ source. is |. the median is +. and. of the sources have. emission spread over. velocity ranges less than 10. kkmss.|.," The mean spread in velocity for a source is $^{-1}$ , the median is $^{-1}$, and of the sources have emission spread over velocity ranges less than 10 $^{-1}$."
+ Only three sources of the 119. (<3%)) have ranges greater than 1: 6.795-0.257 kkmss.1). 7.166|0.131 ) and 10.472|0.027. kkmiss.1).," Only three sources of the 119 $<$ ) have ranges greater than $^{-1}$: 6.795-0.257 $^{-1}$ ), 7.166+0.131 $^{-1}$ ) and 10.472+0.027 $^{-1}$ )."
+ This rarity in velocity. spreads greater than ss! is comparable to both 2 (« A)) and 2? (14 ))., This rarity in velocity spreads greater than $^{-1}$ is comparable to both \citet{caswell09a} $<$ ) and \citet{caswell10mmb1} $\sim$ ).
+ Consistent with the results for the central part of the Galaxy (?).. all the masers in the longitude range 6 to 207 lie well inside the velocity coverage of the MMD (which was chosen to fully sample the velocity range of the CO emission of ?)).," Consistent with the results for the central part of the Galaxy \citep{caswell10mmb1}, all the masers in the longitude range $^{\circ}$ to $^{\circ}$ lie well inside the velocity coverage of the MMB (which was chosen to fully sample the velocity range of the CO emission of \citealt{dame01}) )."
+ There are three sources in the region with velocities of peak emission exceeding (6.368-0.052.. 1.601-0.139 and 7.632-0.109) and there are six sources with negative velocities of peak emission (6.189-0.358. 6.881|0.003. 8.832-0.028. 10.629-0.333.. 10.724-0.334 and 15.665-0.499).," There are three sources in the region with velocities of peak emission exceeding $^{-1}$ (6.368-0.052, 7.601-0.139 and 7.632-0.109) and there are six sources with negative velocities of peak emission (6.189-0.358, 6.881+0.093, 8.832-0.028, 10.629-0.333, 10.724-0.334 and 15.665-0.499)."
+ Only one of the six sources has a significantly. negative velocity («-10 ss. 1). 6.189-0.355.," Only one of the six sources has a significantly negative velocity $<$ -10 $^{-1}$ ), 6.189-0.358."
+ 6668-MlIz methanol masers have the potential to trace he kinematies of Galactic structure: their. velocities are closely linked. to the svstemic velocities (22): and they are only detected: towards regions of high-mass star formation (??77).. regions which are intrinsically associated with the spiral armis.," 6668-MHz methanol masers have the potential to trace the kinematics of Galactic structure: their velocities are closely linked to the systemic velocities \citep{szymczak07, pandian09}; and they are only detected towards regions of high-mass star formation \citep{pestalozzi02b,minier03,xu08}, , regions which are intrinsically associated with the spiral arms."
+ Figure 4. shows longitudes and velocities of the 6668-MlIz methanol sources detected with the MMD survey »etween longitudes 345° and 207 (this paper and 2))., Figure \ref{3-kpcExtCir} shows longitudes and velocities of the 6668-MHz methanol sources detected with the MMB survey between longitudes $^{\circ}$ and $^{\circ}$ (this paper and \citealt{caswell10mmb1}) ).
+ The mascers are shown in relation to the 3-kpe arms as defined by 7 and the spiral zum loci based on the logarithmic spirals of νι with the updates of 2 and the rotation curve of ? with LAU standard LSR parameters.," The masers are shown in relation to the 3-kpc arms as defined by \citet{dame08} and the spiral arm loci based on the logarithmic spirals of \citet{georgelin76}, with the updates of \citet{taylor93} and the rotation curve of \citet{brand93} with IAU standard LSR parameters."
+ Figure 4 includes an inset of the assumed spiral arm pattern. oriented so as to readily. recognise how individual portions of the arms transform to the£e domain.," Figure \ref{3-kpcExtCir} includes an inset of the assumed spiral arm pattern, oriented so as to readily recognise how individual portions of the arms transform to the domain."
+ By plotting in theLe domain we can investigate Galactic structure without the need to correctly assign individual sources to near and far kinematic distances., By plotting in the domain we can investigate Galactic structure without the need to correctly assign individual sources to near and far kinematic distances.
+ Ehe figure shows that there appear to be masers lving between the arms. but. before closer. inspection. of individual anomalies. there are two issues that one must be aware of: the thickness of the arm loci and the choice of rotation curve.," The figure shows that there appear to be masers lying between the arms, but before closer inspection of individual anomalies, there are two issues that one must be aware of: the thickness of the arm loci and the choice of rotation curve."
+ Firstly. the arms plotted in£e diagrams are purposely thin to reduce obscuration. but as such they do not account for the real spatial radial thickness and velocity dispersion within the arms.," Firstly, the arms plotted in diagrams are purposely thin to reduce obscuration, but as such they do not account for the real spatial radial thickness and velocity dispersion within the arms."
+ Secondlsv. the locations (physical and in the£e domain) of the arms depends on the Galactic rotation curve adopted.," Secondly, the locations (physical and in the domain) of the arms depends on the Galactic rotation curve adopted."
+ If the spiral pattern. is. broadly correct. any change in adopted rotation curve will only alter the position of the arms slightly. leaving them as continuous structures. not only in the Galactic disk. but also in theLe domain.," If the spiral pattern is broadly correct, any change in adopted rotation curve will only alter the position of the arms slightly, leaving them as continuous structures, not only in the Galactic disk, but also in the domain."
+ For the Galactic structure investigation presented rere we use the 2? curve. since it closely approximates he assumptions made in deriving the spatial pattern of he spiral arms.," For the Galactic structure investigation presented here we use the \citet{brand93} curve, since it closely approximates the assumptions made in deriving the spatial pattern of the spiral arms."
+" Furthermore we resist adopting the most recently suggested rotation curve and revised LS1t suggested w ὃν, as it is preliminary ancl requires revision (2).."," Furthermore we resist adopting the most recently suggested rotation curve and revised LSR suggested by \citet{reid09}, as it is preliminary and requires revision \citep{mcmillan10}."
+ However. we note that near the Galactic Centre. where we did choose he Ποιά et al.," However, we note that near the Galactic Centre, where we did choose the Reid et al."
+ parameters. (2)... this choice had. negligible impact on the interpretation of that region.," parameters \citep{caswell10mmb1}, this choice had negligible impact on the interpretation of that region."
+ Lt would alfect he current discussion in only one instance (noted. below). oit would have more impact at longitudes further from the Galactic centre.," It would affect the current discussion in only one instance (noted below), but would have more impact at longitudes further from the Galactic centre."
+ ‘Table 2. lists 45 sources associated. kinematically and spatially with the 3-kpe arm features (denoted. by. filled cireles in rel3-kpcelext Cir))., Table \ref{3-kpclist} lists 45 sources associated kinematically and spatially with the 3-kpc arm features (denoted by filled circles in \\ref{3-kpcExtCir}) ).
+ “Phis includes the 42 originally identified in ? and three additional sources. from ?— and the current paper.," This includes the 42 originally identified in \citet{green09b} and three additional sources, from \citet{caswell10mmb1} and the current paper."
+ X further three sources. (two. new to the survey) between longitudes 15° and 20° (17.862|0.074. 19.496|0.115 ancl 19.755-0.128) have velocities in excess of and are potential candidates of an extended 3-kpe structure.," A further three sources (two new to the survey) between longitudes $^{\circ}$ and $^{\circ}$ (17.862+0.074, 19.496+0.115 and 19.755-0.128) have velocities in excess of $^{-1}$ and are potential candidates of an extended 3-kpc structure."
+ They do not align. with spiral armi loci (subject to the caveats already mentioned). but they could be accounted. for by the far side of the 3-kpe arms if this is extended as a continuous ring structure with a radius of ~83.5kkpe (e.g.2)..," They do not align with spiral arm loci (subject to the caveats already mentioned), but they could be accounted for by the far side of the 3-kpc arms if this is extended as a continuous ring structure with a radius of $\sim$ kpc \citep[e.g.][]{sevenster99}."
+ Xn alternative explanation requiring a shift in velocity of the Crux-Seutum arm (by 30 +) seems less likely., An alternative explanation requiring a shift in velocity of the Crux-Scutum arm (by $\sim$ $^{-1}$ ) seems less likely.
+ Also within the 15° and 20 longitude range are eight sources with velocities between ss and kkmss+ (15.607-0.255. 16.585-0.051.16.831|0.079. 17.029-0.071. 18.262-0.244. 18.661|0.034.18.6670.025 and 18.999-0.239) and seven sources with velocities between  and (16.112-0.303.. 18.735-0.227. 18.834-0.300. 18.874|0.053. 19.600-0.234. 19.614|0.011 and 19.701-0.267) which do not align with spiral arm loci.," Also within the $^{\circ}$ and $^{\circ}$ longitude range are eight sources with velocities between $^{-1}$ and $^{-1}$ (15.607-0.255, 16.585-0.051,16.831+0.079, 17.029-0.071, 18.262-0.244, 18.661+0.034,18.667+0.025 and 18.999-0.239) and seven sources with velocities between $^{-1}$ and $^{-1}$ (16.112-0.303, 18.735-0.227, 18.834-0.300, 18.874+0.053, 19.609-0.234, 19.614+0.011 and 19.701-0.267) which do not align with spiral arm loci."
+" Fhese sources may also be accounted for by the continuation bevond 15"" of the 3-kpc arms structure as a ring.", These sources may also be accounted for by the continuation beyond $^{\circ}$ of the 3-kpc arms structure as a ring.
+ We note that if an alternative rotation model were adopted. the shift of the spiral arm loci might account for some of these sources. but would. leave a clillerent group. of orphans.," We note that if an alternative rotation model were adopted, the shift of the spiral arm loci might account for some of these sources, but would leave a different group of orphans."
+ For example. using the LAU solar parameter values and [la rotation curve of 2? would. shift. the loci of the Crux-Seutum and Carina-Sagittarius arms by about (}15kkmss+.," For example, using the IAU solar parameter values and flat rotation curve of \citet{reid09} would shift the loci of the Crux-Scutum and Carina-Sagittarius arms by about $^{-1}$."
+ This allows 16.585-0.051 +). 16.831|0.079 ss. 3). 18.999-0.239 1) and 19.614] *) to be associated with the arms. but. 0.011.clisassociates kkmiss.16.403-0.181 (peak velocity. of 1). 16.662-0.331 (peak velocity. of ty. 18.733-0.224 (peak velocity. of ty ancl19.881- (peak velocity of +).," This allows 16.585-0.051 $^{-1}$ ), 16.831+0.079 $^{-1}$ ), 18.999-0.239 $^{-1}$ ) and 19.614+0.011 $^{-1}$ ) to be associated with the arms, but disassociates 16.403-0.181 (peak velocity of $^{-1}$ ), 16.662-0.331 (peak velocity of $^{-1}$ ), 18.733-0.224 (peak velocity of $^{-1}$ ) and19.884-0.534 (peak velocity of $^{-1}$ )."
+ Llenee. regardless of parameter ancl model. choice. thereare 15 sources unaccounted for by the spiral arms which are candidates for an extended 3-kpc arm structure.," Hence, regardless of parameter and model choice, thereare 18 sources unaccounted for by the spiral arms which are candidates for an extended 3-kpc arm structure."
+ The actual structure of the 3-kpe arms has. been variously interpreted as an expanding ring (requiring an explosive precursor). as material orbiting a bar (cither on a circular orbit or an elliptical orbit. with its major axis," The actual structure of the 3-kpc arms has been variously interpreted as an expanding ring (requiring an explosive precursor), as material orbiting a bar (either on a circular orbit or an elliptical orbit with its major axis"
+importance to see how the lines evolve and if quiescence Is reached.,importance to see how the lines evolve and if quiescence is reached.
+ Though the emission component had nearly disappeared these last years 2004).. the emissions continued to decline until 2004.," Though the emission component had nearly disappeared these last years , the emissions continued to decline until 2004."
+ Now. 1108 has finally reached its quiescent state since the data taken between 2005 and 2009 are very similar (Figs.," Now, 108 has finally reached its quiescent state since the data taken between 2005 and 2009 are very similar (Figs."
+ | and 2))., \ref{108prof} and \ref{108hb}) ).
+ Based on past behaviour2001). it is expected that the emission will strengthen again m the future. as is seen in 1191612.," Based on past behaviour, it is expected that the emission will strengthen again in the future, as is seen in 191612."
+ We note however that. for 1191612. the quiescent state occurs during one-third of the cyele. which would correspond to ~18 yyrs for 1108.," We note however that, for 191612, the quiescent state occurs during one-third of the cycle, which would correspond to $\sim$ yrs for 108."
+ Assuming that the variations are due to an oblique magnetic rotator configuration. this quiescence timescale would however depend on the exact geometry of the system. which may be different between the two objects despite their obvious similarities.," Assuming that the variations are due to an oblique magnetic rotator configuration, this quiescence timescale would however depend on the exact geometry of the system, which may be different between the two objects despite their obvious similarities."
+ from literature. 1t is only known that rremained in absorption during about 10 years. while (44471 did so during several decades2001).," from literature, it is only known that remained in absorption during about 10 years, while $\lambda$ 4471 did so during several decades."
+. However. the exact evolution was not recorded and it cannot be concluded from the older data how long the lines remained constant. in à quiescent state.," However, the exact evolution was not recorded and it cannot be concluded from the older data how long the lines remained constant, in a quiescent state."
+ Previously. a long-term monitoring (3. yrs. with monthly observations) provided a large set of low-resolution spectra.," Previously, a long-term monitoring (3 yrs, with monthly observations) provided a large set of low-resolution spectra."
+ These data revealed the variability of 3 lines: Ηα..Hf.. and .A 446862008a).," These data revealed the variability of 3 lines: , and $\lambda$ 4686."
+. However. this variability was of low amplitude: the EWs measured forHa.. the most variable line. display a dispersion which. though it was 10 times larger than that observed for the neighbouring narrow Diffuse Interstellar Band (DIB). correspond to a change of only in EW (cit also corresponds to a variation in the peak's amplitude of only between the two extreme profiles). whereas the other prototypical Of?p stars have llines varying from absorption to emission. with EWs and peak’s amplitudes varying by >>100¢.," However, this variability was of low amplitude: the EWs measured for, the most variable line, display a dispersion which, though it was 10 times larger than that observed for the neighbouring narrow Diffuse Interstellar Band (DIB), correspond to a change of only in EW (it also corresponds to a variation in the peak's amplitude of only between the two extreme profiles), whereas the other prototypical Of?p stars have lines varying from absorption to emission, with EWs and peak's amplitudes varying by $>>$."
+. Moreover. the aand llines seemed relatively constant. a behaviour different from that of 1108 and 1191612 - though subtle changes may have remained hidden in such low-resolution data.," Moreover, the and lines seemed relatively constant, a behaviour different from that of 108 and 191612 - though subtle changes may have remained hidden in such low-resolution data."
+ A period search further revealed a period of 7.031+0.003d when studying closely the line profile variations ofHa., A period search further revealed a period of $\pm$ 0.003d when studying closely the line profile variations of.
+".. However. the data sampling. aimed at studying monthly-to-yearly variations, was not adequate for detecting such a short period. which thus required confirmation with a more intense temporal sampling."," However, the data sampling, aimed at studying monthly-to-yearly variations, was not adequate for detecting such a short period, which thus required confirmation with a more intense temporal sampling."
+ A new. short-term monitoring was therefore undertaken. yielding 20 high-resolution spectra over 2 weeks.," A new, short-term monitoring was therefore undertaken, yielding 20 high-resolution spectra over 2 weeks."
+ For consistency. the radial velocities and equivalent widths (EWs) were estimated as in(20082): they are given in Table 1..," For consistency, the radial velocities and equivalent widths (EWs) were estimated as in; they are given in Table \ref{tab:148937}."
+ While the aand DIBs display constant EWs. the Balmer. υἱ 34686. and 2155876 lines show a larger dispersion.," While the and DIBs display constant EWs, the Balmer, $\lambda$ 4686, and $\lambda$ 5876 lines show a larger dispersion."
+ When plotted against time (Fig. 3)).," When plotted against time (Fig. \ref{148937rvew}) ),"
+4 obvious modulations are detected., obvious modulations are detected.
+ aand .A446586 lines display maximum emission whenHf..Hy.. and A4 55876 present à minimum absorption.," and $\lambda$ 4686 lines display maximum emission when, and $\lambda$ 5876 present a minimum absorption."
+ The variability detected when looking at EWs ts confirmed by the Temporal Variance spectrum1793]ful96. which compares. in each wavelength bin. the variations between spectra with the scatter expected from random noise.," The variability detected when looking at EWs is confirmed by the Temporal Variance spectrum, which compares, in each wavelength bin, the variations between spectra with the scatter expected from random noise."
+ It is important to note that the variability has not the same amplitude for all lines of a given element: for the Balmer lines. the TVS is very significant forΗα.. significant forHB.. and barely significant forHy.," It is important to note that the variability has not the same amplitude for all lines of a given element: for the Balmer lines, the TVS is very significant for, significant for, and barely significant for."
+. This might explain why the variability of. e.g.. (44471 is not detected - its amplitude. expected to be lower than for .155876. is certainly hidden by the noise.," This might explain why the variability of, e.g., $\lambda$ 4471 is not detected - its amplitude, expected to be lower than for $\lambda$ 5876, is certainly hidden by the noise."
+ Note that no significant variation of the A 44650 lines is detected., Note that no significant variation of the $\lambda$ 4650 lines is detected.
+ Periodograms were calculated for the 20-exposures dataset using the techniques of in the wavelength interval where lines vary significantly2001)., Periodograms were calculated for the 20-exposures dataset using the techniques of in the wavelength interval where lines vary significantly.
+. The results are shown in Fig. 5: , The results are shown in Fig. \ref{148937four}: :
+the highest peak is located at a frequency of 0.1315d7!. 0.135347!. 0.1373d7!. and 0.139747! for 4440686.Hp.. 455876. andΗα.. respectively.," the highest peak is located at a frequency of $^{-1}$ , $^{-1}$, $^{-1}$, and $^{-1}$ for $\lambda$ 4686, $\lambda$ 5876, and, respectively."
+ These frequencies translate into a period of 7.16-7.6040.40d., These frequencies translate into a period of $\pm$ 0.40d.
+ The relatively large error. compared to the old dataset simply stems from the different observational timespans (2 weeks vs yyrs) while the period differences from one line to the other have two origins: the noise. which affects more the fainter lines such as (55876. and the different amplitudes of the variations (the signal is indeed more difficult to catch in case of low amplitude changes. e.g. for 455876).," The relatively large error, compared to the old dataset simply stems from the different observational timespans (2 weeks vs yrs) while the period differences from one line to the other have two origins: the noise, which affects more the fainter lines such as $\lambda$ 5876, and the different amplitudes of the variations (the signal is indeed more difficult to catch in case of low amplitude changes, e.g. for $\lambda$ 5876)."
+ The Iline. which is both strong and the most variable line. yields the most reliable period. te. 7.16£0.40d from the sole new dataset.," The line, which is both strong and the most variable line, yields the most reliable period, i.e. $\pm$ 0.40d from the sole new dataset."
+ Taking the errors intoaccount!.. this agrees with the results reported by(2008a).," Taking the errors into, this agrees with the results reported by."
+. The 7d period was thus not an artefact from our inadequate time sampling., The 7d period was thus not an artefact from our inadequate time sampling.
+ 1148937 therefore appears quite similar to the two other Galactic Of?p stars. though with a shorter period. a smaller amplitude of the variations and with the sole exception of the apparent constancy in the llines.," 148937 therefore appears quite similar to the two other Galactic Of?p stars, though with a shorter period, a smaller amplitude of the variations and with the sole exception of the apparent constancy in the lines."
+ A spectropolarimetric monitoring should be undertaken to see if this could be related toe.g.a low inclination of the star's rotation axis., A spectropolarimetric monitoring should be undertaken to see if this could be related toe.g.a low inclination of the star's rotation axis.
+(1995).,.
+. The information criteria (Table 23) choose AZ» and AZ; as the best two models., The information criteria (Table \ref{M_j:IC}) ) choose $M_2$ and $M_3$ as the best two models.
+ However. they disagree as AIC suggests 1». whereas BIC proposes A5.," However, they disagree as AIC suggests $M_2$ , whereas BIC proposes $M_3$."
+ The maximum likelihood functions under these two models are 26235.67 for Ado and 26237.19 for Δι., The maximum likelihood functions under these two models are $-26235.67$ for $M_2$ and $-26237.19$ for $M_3$.
+" Although Avo maximises the likelihood compared to A;. the fact that it contains one more parameter than Ad, does not allow Bayes factor estimations to clearly favour Alo."," Although $M_2$ maximises the likelihood compared to $M_3$, the fact that it contains one more parameter than $M_3$ does not allow Bayes factor estimations to clearly favour $M_2$ ."
+ Therefore. A4 cannot be excluded.," Therefore, $M_3$ cannot be excluded."
+ Nevertheless. we select model AM» for further analysis.," Nevertheless, we select model $M_2$ for further analysis."
+ Table 3 provides all the information we need for the parameter estimation., Table \ref{M_1:estimation} provides all the information we need for the parameter estimation.
+ We will not comment separately on each parameter becausethere are too many., We will not comment separately on each parameter becausethere are too many.
+ Instead we will try to focus on the, Instead we will try to focus on the
+A population of star-forming ealaxics at high. recshilts are characterized. by their strong. Lyman-a (Lya) line emission.,A population of star-forming galaxies at high redshifts are characterized by their strong $\alpha$ $\alpha$ ) line emission.
+ Such Lya emitters (LAXIS) have been found. at various redshifts by narrow-band surveys using S-1O m class telescopes (Lluetal.1998.1999.2002:Ixocaira2003:2004.2005:Iveetal. 2008).," Such $\alpha$ emitters (LAEs) have been found at various redshifts by narrow-band surveys using 8-10 m class telescopes \citep{Hu98, Hu99, Hu02,Kodaira03, Shimasaku03, Shimasaku06, 
+Ha2004, Ou04, Ou05, Ouchi08, Taniguchi05, Matsuda04, Matsuda05, Iye06}."
+. lt is generally thought that the strong Lye emission physically originates from star-forming regions regions) in à voung starburst galaxy., It is generally thought that the strong $\alpha$ emission physically originates from star-forming regions regions) in a young starburst galaxy.
+ Some LAEs have very [large equivalent widths (LEM >...) exceeding 400A. which is dillieult to explain with ordinary stellar population svnthesis models (e.g... Charlot.&Fall.(1993):Schaerer.. (2003))).," Some LAEs have very large equivalent widths ${\rm EW_{Ly\alpha}}$ ) exceeding $400 {\rm \AA}$, which is difficult to explain with ordinary stellar population synthesis models (e.g., \citet{CF93, Schaerer03}) )."
+ Alternative physical models include cooling radiation from a primordial collapsing gas (Llaiman.Spaans.&Quataert.2000:Fardalctal. 2001).. from a galactic wind-driven shell (Taniguchi.&Shiova.2000).. and from supernova remnants (Mori.Umemura.&Ferrara.2004:Mori.Umeniura. 2006).," Alternative physical models include cooling radiation from a primordial collapsing gas \citep{Haiman00, Fardal01}, , from a galactic wind-driven shell \citep{TS00}, and from supernova remnants \citep{MUF04, MU2006}."
+.. Recent large LAL surveys. provided. an array of statistical properties of LALs such as the Lye luminosity function. two-point angular correlation function. and. the evolution of them.," Recent large LAE surveys provided an array of statistical properties of LAEs such as the $\alpha$ luminosity function, two-point angular correlation function and the evolution of them."
+ The observations generally suggest that LAEsS are not simply a subset of star-forming ealaxies., The observations generally suggest that LAEs are not simply a subset of star-forming galaxies.
+ Indeed. theoretical models: proposed. κο far do not. Lully explain the observed. properties.," Indeed, theoretical models proposed so far do not fully explain the observed properties."
+ In. particular. reproducing very large equivalent: widths of some bright LIES appears to be challenging.," In particular, reproducing very large equivalent widths of some bright LAEs appears to be challenging."
+ Lya photons are easily absorbed by dust and thus it is naively expected that LAL is à very voung anc dust-free. galaxy., $\alpha$ photons are easily absorbed by dust and thus it is naively expected that LAE is a very young and dust-free galaxy.
+ While some observations ancl theoretical studies actually support the notion. (Cawiseret.al.2006.)07:Mori.&Umoemura.2006:Shimizuetal. 2007).. more recent. multi-wavelength observations of LAEs in optical. infrared. and. sub-millimeter suggest. that there are LES that are indeed old and dustv (Finkelsteinetal.2007:Laiμαeinetal.2009c:Tamura2009:Ono 2010).," While some observations and theoretical studies actually support the notion \citep{Ga2006, Ga2007, MU2006, S2007}, more recent multi-wavelength observations of LAEs in optical, infrared and sub-millimeter suggest that there are LAEs that are indeed old and dusty \citep{Fin2007, Lai2008, Matsuda2007, Uchimoto2008, Fin2009c, SMGLAE, Ono2010}."
+.Interestingly. such a population increases with decreasing redshift (Nilssonetal. 2009).," Interestingly, such a population increases with decreasing redshift \citep{Nilsson2009}."
+. There is even an evidence ju some sub-millimetergalaxies show strong Lya emission, There is even an evidence that some sub-millimetergalaxies show strong $\alpha$ emission
+ , 
+ihe minimum Lorentz [actor of the injected electrons.,the minimum Lorentz factor of the injected electrons.
+ As the flare progresses. the minimum Lorentz [actor of the evolving electrons will eventually drop to the optimum values for the enission al [frequencies below (he turn-over frequency.," As the flare progresses, the minimum Lorentz factor of the evolving electrons will eventually drop to the optimum values for the emission at frequencies below the turn-over frequency."
+" However. if the initial 5,,5, is hieh enough. (he optimum value might only be reached after the crossing time."," However, if the initial $\gamma_{{{}} min}$ is high enough, the optimum value might only be reached after the crossing time."
+" In this case. the ERC emission at the frequency (hat corresponds to this optimum value will continue to grow even after time /,.C when the shock [ront exits the excitation region and the acceleration of electrons stops."," In this case, the ERC emission at the frequency that corresponds to this optimum value will continue to grow even after time $t_{{{}} ac}$ when the shock front exits the excitation region and the acceleration of electrons stops."
+ This phenomenon should not affect the SSC flares in the same fashion since seed photons from a broad range of Irequencies contribute equally to the SSC emission al a given [requency of observation., This phenomenon should not affect the SSC flares in the same fashion since seed photons from a broad range of frequencies contribute equally to the SSC emission at a given frequency of observation.
+ The frequencies at. which electrons with Lorentz [actor sin enit svuchrotron radiation are generally lower (han the svuchrotron sell-absorption frequency lor realistic parameters., The frequencies at which electrons with Lorentz factor $\gamma_{min}$ emit synchrotron radiation are generally lower than the synchrotron self-absorption frequency for realistic parameters.
+ Therefore. svnchirotron [flares should not be expected to exhibit this effect. either.," Therefore, synchrotron flares should not be expected to exhibit this effect, either."
+" The ERC light curves lor the viewing angle 8&,,;=90° are presented in Fig.", The ERC light curves for the viewing angle $\theta_{obs}=90^\circ$ are presented in Fig.
+ 7 lor “nin—100 and Fig., \ref{melc.90} for $\gamma_{min}=100$ and Fig.
+" 8. lor 5,,5,=LO.", \ref{melc.90a} for $\gamma_{min}=10$.
+" At a viewing angle of 90° the light curves at higher frequencies. defined by [4<1. peak al ~/,,/2 as a result of 1) rapid decay of electrons that dominate the observed emission al higher frequencies and 2) the circular geometry of the source along the line of sight."," At a viewing angle of $90^{\circ}$ the light curves at higher frequencies, defined by $t^E_{\nu}\lesssim{}t_{ac}$, peak at $\sim{}t_{{{}}
+ac}/2$ as a result of 1) rapid decay of electrons that dominate the observed emission at higher frequencies and 2) the circular geometry of the source along the line of sight."
+" At lower frequencies defined by (/>1, the maximum is closer to /,,. since this is when the emission [fills the entire volume of the source."," At lower frequencies defined by $t^E_{\nu}\gtrsim{}t_{ac}$, the maximum is closer to $t_{ac}$ since this is when the emission fills the entire volume of the source."
+ The mechanism that causes extra delay of the ERC emission. which was described above. affects ERC light curves at any viewing angle. including θε.=90°.," The mechanism that causes extra delay of the ERC emission, which was described above, affects ERC light curves at any viewing angle, including $\theta_{obs}=90^\circ$."
+" IIowever. when the viewing angle 90"" the ellect is not as obvious since the delays due to the geometrical shape of the source and energy stratification are equally important."," However, when the viewing angle $\sim90^{\circ}$ the effect is not as obvious since the delays due to the geometrical shape of the source and energy stratification are equally important."
+ In the calculations discussed here (he parameters have been selected such (hat. inverse Compton energv losses dominate over svnchirotron losses. ρω2230ug.," In the calculations discussed here the parameters have been selected such that inverse Compton energy losses dominate over synchrotron losses, $u_{rad}\approx30u_B$."
+ This. in particular. means that the decay time of svnehrotron emission is shorter (han that caleulatecl [rom svuchrotron losses alone by a factor 7ρω αμ.," This, in particular, means that the decay time of synchrotron emission is shorter than that calculated from synchrotron losses alone by a factor $\sim{}u_{rad}/u_B$ ."
+ In Paper I we neglected external emission: the break frequency. of the svnchrotron spectrum was ab ~10/7Hz while the svnchrotron sell-absorption frequency ~LO! Hz.," In Paper I we neglected external emission; the break frequency of the synchrotron spectrum was at $\sim10^{12}\,\mbox{Hz}$ while the synchrotron self-absorption frequency $\sim10^{10}\,\mbox{Hz}$ ."
+ The dominance of ERC losses in the present calculations shifts the break frequency of the svichrotron spectrum to ~10?Hz. which is less than the sell-absorption lrequency.," The dominance of ERC losses in the present calculations shifts the break frequency of the synchrotron spectrum to $\sim10^9\,\mbox{Hz}$, which is less than the self-absorption frequency."
+ Therefore. the svnchrotron light curves at all Irequencies of interest originatem [rom a volume of the source that is limited by frequency stratification.," Therefore, the synchrotron light curves at all frequencies of interest originate from a volume of the source that is limited by frequency stratification."
+" These lieht curves are expected (o peak al /,./2 when 0,4,~907.", These light curves are expected to peak at $\sim{}t_{ac}/2$ when $\theta_{obs}\sim90^{\circ}$.
+" In contrast. when 6,=0° the 5P""nchrotron [lares at. high lrequencies (at which /,<ac ο) ave characterized by a quick rise. flat top. and equally rapid decay after the crossing time."," In contrast, when $\theta_{obs}=0^{\circ}$ the synchrotron flares at high frequencies (at which $t_\nu < t_{\rm ac}$ ) are characterized by a quick rise, flat top, and equally rapid decay after the crossing time."
+ Comparing this behavior with that of the ERC light curves. one can assert that there must be a clelay of about half the erossing(ime betweenthe svuchrotvon fIares ancl the ERC," Comparing this behavior with that of the ERC light curves, one can assert that there must be a delay of about half the crossingtime betweenthe synchrotron flares and the ERC"
+rotation axis and the dipole axis (c.g... with Ax907 or 2707). Gi) the enuission is produced due to the shell-dviven instability. aud Gii) the rotation axis inclination ὁ differs ∖from €!907.,"rotation axis and the dipole axis (e.g., with $\lambda\simeq 90^{\circ}$ or $270^{\circ}$ ), (ii) the emission is produced due to the shell-driven instability, and (iii) the rotation axis inclination $i$ differs from $90^{\circ}$."
+ As popa result.ovo] the (ccounterpart3343» of: tle5 pulvas D is invisible because the emission source is occulted by the star at that time.," As a result, the counterpart of the pulse B is invisible because the emission source is occulted by the star at that time."
+ Figure 8 shows the simulatecdk dynamic spectrum and light curve for the following simulation parameters:, Figure \ref{FigMod} shows the simulated dynamic spectrum and light curve for the following simulation parameters:
+The prototype T dwarf was discovered in 1995 as a companion to the nearby M dwarf star GI229 (Nakajima et al.,The prototype T dwarf was discovered in 1995 as a companion to the nearby M dwarf star Gl229 (Nakajima et al.
+ 1995). and today we know of 155 T dwarfs (www.dwarfarchives.org as of 10/2009).," 1995), and today we know of 155 T dwarfs (www.dwarfarchives.org as of 10/2009)."
+ The majority of these are a result of the near-infrared 2 Micron All Sky Survey (?.2MASS) and the deep optical Sloan Digital Sky Survey (?.SDSS)..," The majority of these are a result of the near-infrared 2 Micron All Sky Survey \citep[2MASS]{2006AJ....131.1163S}
+ and the deep optical Sloan Digital Sky Survey \citep[SDSS]{2000AJ....120.1579Y}."
+ Once discovered. significant efforts were undertaken to determine their distances (22?) to map out the lower end of the Hertzsprung-Russell (H-R) diagram and to constrain models.," Once discovered, significant efforts were undertaken to determine their distances \citep{2004AJ....127.2948V,
+ 2003AJ....126..975T, 2002AJ....124.1170D} to map out the lower end of the Hertzsprung-Russell (H-R) diagram and to constrain models."
+ However. in these surveys only a handful of objects have spectral types T7 or later and more examples of intrinsically fainter and cooler objects are crucial to understand both such objects and. in turn. exoplanets of similar masses.," However, in these surveys only a handful of objects have spectral types T7 or later and more examples of intrinsically fainter and cooler objects are crucial to understand both such objects and, in turn, exoplanets of similar masses."
+ In 2005 the UKIRT Infrared Deep Sky Survey afterUKIDSS) began with the Wide Field Camera (?.CAM) large-field infrared camera., In 2005 the UKIRT Infrared Deep Sky Survey \citep[hereafter UKIDSS]{2007MNRAS.379.1599L} began with the Wide Field Camera \citep[WFCAM]{2007AA...467..777C} large-field infrared camera.
+ This survey is going magiitudes fainter than the 2MASS survey and has already revealed a number of extremely faint and cool T dwarfs (2???).. A," This survey is going magnitudes fainter than the 2MASS survey and has already revealed a number of extremely faint and cool T dwarfs \citep{2009MNRAS.397..258L, 2008MNRAS.391..320B, 2008MNRAS.390..304P,
+ 2007MNRAS.379.1423L}."
+s part of the followup for the UKIDSS surveys we have started a program to determine parallaxes for the coolest objects being discovered., As part of the followup for the UKIDSS surveys we have started a program to determine parallaxes for the coolest objects being discovered.
+ The first and most obvious application of measured distances is to derive a luminosity which is used to populate the H-R diagram and constrain models., The first and most obvious application of measured distances is to derive a luminosity which is used to populate the H-R diagram and constrain models.
+ Distances are also used to identify unresolved companions via over-luminosities. and to provide space motions which provide indications of an object's origin and age.," Distances are also used to identify unresolved companions via over-luminosities, and to provide space motions which provide indications of an object's origin and age."
+ There is a large degree of uncertainty in current model analyses of the late-type T dwarfs. due to known madequacies in the atmospheric models (e.g. Leggett et al.," There is a large degree of uncertainty in current model analyses of the late-type T dwarfs, due to known inadequacies in the atmospheric models (e.g. Leggett et al."
+ 2009)., 2009).
+ However. evolutionary models of cooling brown dwarfs are well understood. and show that all brown dwarfs older than about 200 Myr have a radius within about of Jupiter's (e.g. ?)..," However, evolutionary models of cooling brown dwarfs are well understood, and show that all brown dwarfs older than about 200 Myr have a radius within about of Jupiter's \citep[e.g.][]{2001RvMP...73..719B}."
+ Hence there is a tight relationship between luminosity and effective temperature (7.4) for brown dwarfs (via the Stefan-Boltzman law)., Hence there is a tight relationship between luminosity and effective temperature $T_{\rm eff}$ ) for brown dwarfs (via the Stefan-Boltzman law).
+ Parallax determinations combined with flux-calibrated observed spectra provide luminosities. and hence the best measurements of Τομ for brown dwarfs currently available.," Parallax determinations combined with flux-calibrated observed spectra provide luminosities, and hence the best measurements of $T_{\rm eff}$ for brown dwarfs currently available."
+ Once Tay is known. gravity and metallicity can be constrained by comparing synthetic spectra to the observations (e.g.?)..," Once $T_{\rm eff}$ is known, gravity and metallicity can be constrained by comparing synthetic spectra to the observations \citep[e.g.][]{2007ApJ...656.1136S}."
+ ? use luminosity to determine Zr for the T8.5 Wolf 940B. and show that an analysis using synthetic spectra alone would have overestimated the temperature by ~ 100 K. for this 600 K object.," \citet{2009MNRAS.395.1237B} use luminosity to determine $T_{\rm eff}$ for the T8.5 Wolf 940B, and show that an analysis using synthetic spectra alone would have overestimated the temperature by $\sim$ 100 K, for this 600 K object."
+ The latest spectral type currently defined is T9 and three such objects are currently known: ULAS (?.hereafterULAS0034).. CFBDS J005910.90-011401.3 (?) and ULAS J133553.45 +113005.2 (?)..," The latest spectral type currently defined is T9 and three such objects are currently known: ULAS \citep[hereafter ULAS
+0034]{2007MNRAS.381.1400W}, CFBDS J005910.90-011401.3 \citep{2008AA...482..961D} and ULAS J133553.45 +113005.2 \citep{2008MNRAS.391..320B}."
+" They have effective temperatures in the range 500 to 700 K. and masses of 5 to 50 M, corresponding to ages of 0.1 to 10 Gyr (e.g. Leggett et al."," They have effective temperatures in the range 500 to 700 K, and masses of 5 to 50 $_J$ corresponding to ages of 0.1 to 10 Gyr (e.g. Leggett et al."
+ 2009)., 2009).
+ Here we present the first results of the UKIRT parallax program for ULAS 0034., Here we present the first results of the UKIRT parallax program for ULAS 0034.
+ In Sect., In Sect.
+ 2 we describe the observations and. procedures for the program. in Sect.," 2 we describe the observations and procedures for the program, in Sect."
+ 3 we provide the results for ULAS 0034 and in Sect., 3 we provide the results for ULAS 0034 and in Sect.
+ 4 we find the most consistent set of physical parameters for ULAS 0034 by combining available observations and models., 4 we find the most consistent set of physical parameters for ULAS 0034 by combining available observations and models.
+ observations were all made on the UKIRT 3.8 m telescope using the WFCAM imager., observations were all made on the UKIRT 3.8 m telescope using the WFCAM imager.
+ This ts the instrument being used to carry out the UKIDSS surveys. and the calibrated images were taken directly from the pipeline for that All observations are carried out in queue override mode. allowing us to be very flexible in the scheduling. maximising the parallax factor and observing close to meridian passage.," This is the instrument being used to carry out the UKIDSS surveys, and the calibrated images were taken directly from the pipeline for that All observations are carried out in queue override mode, allowing us to be very flexible in the scheduling, maximising the parallax factor and observing close to meridian passage."
+ The observational sequence we adopt is very similar to the J-band imaging sequence in the Large Area Survey (LAS) component of UKIDSS (e.g. ?).. , The observational sequence we adopt is very similar to the $J$ -band imaging sequence in the Large Area Survey (LAS) component of UKIDSS \citep[e.g.][]{2007MNRAS.375..213W}. .
+A longer exposure time was used to provide an increased signal-to-noise for the targets.," A longer exposure time was used to provide an increased signal-to-noise for the targets,"
+velocity via ἐξ3/4.4.,velocity via $\beta = 3/4 \beta_{sh}$.
+ We therefore construct a relationship between emission time and fluid velocity similar to that between emission time and shock velocity: We assume that the jet does not spread sideways throughout the relativistic phase of its evolution., We therefore construct a relationship between emission time and fluid velocity similar to that between emission time and shock velocity: We assume that the jet does not spread sideways throughout the relativistic phase of its evolution.
+ However. at some point the blast wave become spherical.," However, at some point the blast wave become spherical."
+" For if we keep the opening angle fixed. but do take E; to dictate the ST solution instead of Ej, we would underestimate the final flux by integrating over an integration domain that is too small."," For if we keep the opening angle fixed, but do take $E_j$ to dictate the ST solution instead of $E_{iso}$ we would underestimate the final flux by integrating over an integration domain that is too small."
+" We will assume that the jet starts spreading sideways when it has reached the nonrelativistic phase and we take this moment to be given by At this point. the jet starts spreading sideways with the speed of sound c,. leading to In the nonrelativistic regime df? and df, are identical."," We will assume that the jet starts spreading sideways when it has reached the nonrelativistic phase and we take this moment to be given by At this point, the jet starts spreading sideways with the speed of sound $c_s$, leading to In the nonrelativistic regime $d t_{obs}$ and $d t_e$ are identical."
+" The ST solution for the speed of sound for adiabatic index 5/3 18 given by c,2r/V20t. leading to until spherical symmetry is reached."," The ST solution for the speed of sound for adiabatic index $5/3$ is given by $c_s = r/\sqrt{20} t$, leading to until spherical symmetry is reached."
+ By contrast. Rhoads 90 take O~z(0;-cfet. (where ¢’ time in the comoving frame. and € a solid angle) as the starting point.," By contrast, Rhoads '99 take $\Omega \approx \pi ( \theta_j + c_s t' / c t_e )^2$ (where $t'$ time in the comoving frame, and $\Omega$ a solid angle) as the starting point."
+ From the local fluid conditions the local emissivity can be calculated., From the local fluid conditions the local emissivity can be calculated.
+ In the case of slow cooling. we define The definition for fast cooling is analogous (see also ?)).," In the case of slow cooling, we define The definition for fast cooling is analogous (see also \citealt{Sari1998}) )."
+" The peak emissivity is given by The synchrotronPy break frequencyJj 7”,AW depends on its corresponding critical electron Lorentz factor ~/,.AW leading to For the cooling break frequency an identical relation between frequency and electron Lorentz factor holds."," The peak emissivity is given by The synchrotron break frequency $\nu'_m$ depends on its corresponding critical electron Lorentz factor $\gamma'_m$, leading to For the cooling break frequency an identical relation between frequency and electron Lorentz factor holds."
+ The eritical Lorentz factor is now given by which follows from the electron kinetic equation when synchrotron losses dominate over adiabatic expansion and the cooling time is approximated by the lab frame time since the explosion., The critical Lorentz factor is now given by which follows from the electron kinetic equation when synchrotron losses dominate over adiabatic expansion and the cooling time is approximated by the lab frame time since the explosion.
+ Synchrotron self-absorption is included in the model using the assumption that emission and absorption occur in a homogeneous shell., Synchrotron self-absorption is included in the model using the assumption that emission and absorption occur in a homogeneous shell.
+ The solution to the linear equation of radiative transfer then dictates that we need to replace eq., The solution to the linear equation of radiative transfer then dictates that we need to replace eq.
+" A2. by The optical depth 7=—a,AR!.", \ref{flux_model2_equation} by The optical depth $\tau \approx - \alpha_\nu \Delta R$.
+". Emissivity and absorption translate between frames using e,=SOpre, and respectively."," Emissivity and absorption translate between frames using $\epsilon'_{\nu'} = \gamma^2 (1 - \beta
+\mu)^2 \epsilon_\nu$ and $\alpha'_{\nu'} = \alpha_\nu / \gamma (1 -
+\beta \mu)$ respectively."
+" In our simplified model we calculate the absorption coefficient a7, under the assumption that electron cooling does not influence it.", In our simplified model we calculate the absorption coefficient $\alpha'_{\nu'}$ under the assumption that electron cooling does not influence it.
+" This assumption is justified when the self-absorption break frequency 7, lies well below the cooling break frequency v,. which is the case for all applications of the model in this paper."," This assumption is justified when the self-absorption break frequency $\nu_a$ lies well below the cooling break frequency $\nu_c$, which is the case for all applications of the model in this paper."
+ Also approximating the synchrotron spectral shape by just two sharply connected power laws we then find for the self-absorption coefficient:, Also approximating the synchrotron spectral shape by just two sharply connected power laws we then find for the self-absorption coefficient:
+One of the remaining problems vet to be understood within star formation is the collapse of a cloud into a star. and the relatively slow star formation rate observed.,"One of the remaining problems yet to be understood within star formation is the collapse of a cloud into a star, and the relatively slow star formation rate observed."
+ This problem. suggests there is some form. of support preventing the clouds from collapse (at least. initially)., This problem suggests there is some form of support preventing the clouds from collapse (at least initially).
+ For instance. if typical (AL 107.+ AL.) molecular clouds had no support. and so were in frec-fall collapse. with all the mass going into stars. cloud. lifetimes would be unrealistically short. and the star formation rate would be up to three orders of magnitude greater than observed. (Alouschovias1976).," For instance, if typical (M $\sim$ $^{3-4}$ $_{\odot}$ ) molecular clouds had no support, and so were in free-fall collapse, with all the mass going into stars, cloud lifetimes would be unrealistically short, and the star formation rate would be up to three orders of magnitude greater than observed \citep{mous76}."
+. Phermal pressure is weak compared to the gravitational stresses in the cloud. so the support may. come from the magnetic field that permeates the eas or the pressure of turbulent. edcdies indeed. it is likely that these two mechanisms are coupled.," Thermal pressure is weak compared to the gravitational stresses in the cloud, so the support may come from the magnetic field that permeates the gas or the pressure of turbulent eddies – indeed it is likely that these two mechanisms are coupled."
+ Spinning elongated. dust. erains in such star forming regions can become aligned to the local magnetic field. such that the grains’ semi-major axes are perpendicular (on average) to the magnetic field lines.," Spinning elongated dust grains in such star forming regions can become aligned to the local magnetic field, such that the grains' semi-major axes are perpendicular (on average) to the magnetic field lines."
+ The net thermal emission from such grains is polarised (perpendicular to the direction of the magnetic field lines). and can be used to trace the morphology of the magnetic field as seen projected onto the plane of the sky.," The net thermal emission from such grains is polarised (perpendicular to the direction of the magnetic field lines), and can be used to trace the morphology of the magnetic field as seen projected onto the plane of the sky."
+ Further information can be gained on the magnetic field via Zeeman splitting (Sarma.Troland&Romney 2001).. which vields the line of sight magnetic Ποιά strength.," Further information can be gained on the magnetic field via Zeeman splitting \citep{sarma}, which yields the line of sight magnetic field strength."
+ lon/neutral linewidths (Ποιοetal.2000).. can also be used to gain information on the orientation of the magnetic field.," Ion/neutral linewidths \citep{houde1}, can also be used to gain information on the orientation of the magnetic field."
+ In order to understand the importance of the magnetic Ποιά on Large scales. consistent with tracing the magnetic Ποια. in the envelope of these cores. we present SCUBA S5Üm imaging polarimetry of 20 star forming cores.," In order to understand the importance of the magnetic field on large scales, consistent with tracing the magnetic field in the envelope of these cores, we present SCUBA $\mu$ m imaging polarimetry of 20 star forming cores."
+ We also introduce a new. improved (more reoesentative) method of data reduction for SCUBA polarimetry. which presents the errors in polarisation in a truer light.," We also introduce a new, improved (more representative) method of data reduction for SCUBA polarimetry, which presents the errors in polarisation in a truer light."
+ We discuss the implied magnetic field morphology across the cores. as well as the estimated. field srength for some of the regions using the moclificcd Chandrasekhar&Fermi(1953.henceforth.CE) method.," We discuss the implied magnetic field morphology across the cores, as well as the estimated field strength for some of the regions using the modified \citet[][
+henceforth CF]{cf} method."
+ We use this estimate to incicate the importance of the magnetic field. on these scales on the support of these clouds. and compare the mean magnetic field. direction to the outUow directions. to see if they correspond in the plane of the sky.," We use this estimate to indicate the importance of the magnetic field on these scales on the support of these clouds, and compare the mean magnetic field direction to the outflow directions, to see if they correspond in the plane of the sky."
+ The sample of high-mass star-forming regions were selected from the Arcetri water maser atlas (Comorettoetal.1990: based. on the following criteria: i)," The sample of high-mass star-forming regions were selected from the Arcetri water maser atlas \citep{comoretto,brand} based on the following criteria: i)"
+Padmanabhan etal. (2007). 3. Mission Concept pphotometry times. whileit alsocrucial to (PSF).many The former criterion required tobe ableknowledgeto long-termpoint-spread trend. eliminate insi rument-related photometry. because they need programs and exquisite PSFvery control/knowledge suchin /Z//PPARCOS.. manyGALA repeat ΔΙΑ ancl the canceled £AALE citepGpproject (Johnston aud the proposed ALEX 2003) AM EX50030ndO SS citepAMENX3s005.manissions(Johnslonelal. D 2006).. aastrometric missions. andtheir & Gaume 2002:Gaume& Olling 2002)]].,"The moderate dispersion can be achieved by a low-power dispersive element prism in the light-path (such as was planned by or, or ESO's WIFI imager) The above strategy implies that will visit the average star 168 times in 5 years, while 61 observations are taken during those 168 6.6-hour We simulated the expected number of transits in the following manner: 1) we perform a full-length mission simulation that measures the number of visits at arbitrary points on the sky, 2) these data are then used to determine a sky-averaged probability for detecting transit-like events with a duration determined by the period of the planet (transits last longer in longer-period orbits), 3) here we assume a G2V primary and a planet with Jupiter's radius, 4) we use the modified TT2002 $PDF$, multiplied this by the probability that the planet is seen edge-on, and the probability of observing 5 transits as determined in (2) above, 5) we use a simple star-count model of the solar neighborhood to predict how many dwarf stars are with our magnitude limit, and 6) generate the total number of observable planetary transits."
+ wepropose the following LEAVITT significantlyto precision astrophysics with the d," The combined effects of not observing continuously, and the period distribution of the ESGPs indicates that is $\ga
+50$ complete for periods shorter than 18 days."
+iscovery ofCepheid variables inthe Magel-, This compares very well with 2.4 days for (for which we performed the same simulations).
+" ]anic Alternatively. LEAVITT could standlor: ""LEgacy."," In fact, outperforms by a factor of 100 (460) at a period of 10 (20) days."
+ Astrophysics andTrans, This is mostly due to s very slow precession rate of 75 days.
+iting AIDEN-class mission., will produce roughly 20 times more PTs than.
+ space The basicproperty (transit needstohave," On the other hand, can find planets with a radius of the Earth at $V=12$, while can detect Neptune-size objects at that magnitude."
+}).,).
+ This provides supporting evidence that both populations trace different evolutionary stages of the same inassive galaxy population. as would be expected in the ULIRG-quasar transition scenario proposed by Sanders ct ((1988)).," This provides supporting evidence that both populations trace different evolutionary stages of the same massive galaxy population, as would be expected in the ULIRG-quasar transition scenario proposed by Sanders et \citeyear{san88}) )."
+ J123707|6211 is a prototypical example of an SM caving an carly merecr stage. found in the peak epoch of ealaxy formation.," J123707+6214 is a prototypical example of an SMG during an early merger stage, found in the peak epoch of galaxy formation."
+ Higher resolution. dvuamical mapping of eenission in this intriguing system (and others) ds desirable to narrow down ooo through ανασα] ass nieasurenieuts over several resolution elemeuts. as possible with the full EVLA in the future.," Higher resolution, dynamical mapping of emission in this intriguing system (and others) is desirable to narrow down $\alpha_{\rm CO}$ through dynamical mass measurements over several resolution elements, as possible with the full EVLA in the future."
+ A more complete census of oobservatious of SM will provide the necessary context to interpret the results of such investigations., A more complete census of observations of SMGs will provide the necessary context to interpret the results of such investigations.
+ Such studies provide the most direct micaus to constrain the eas fraction. total mass and evolutionary state of SAICs. which is necessary to better uuderstaud the evolutionary path of massive ealaxics through their most active phases. aud to constrain the molecular gas mass deusity of the universe.," Such studies provide the most direct means to constrain the gas fraction, total mass and evolutionary state of SMGs, which is necessary to better understand the evolutionary path of massive galaxies through their most active phases, and to constrain the molecular gas mass density of the universe."
+ We thank the referee for a critical reading of the manuscript aud for a helpful report., We thank the referee for a critical reading of the manuscript and for a helpful report.
+ DR acknowledges support from NASA through Hubble Fellowship eraut IST-HE-512535.01. awarded by STScl. operated by AURA for NASA. uncer contract 55-26555.," DR acknowledges support from NASA through Hubble Fellowship grant HST-HF-51235.01 awarded by STScI, operated by AURA for NASA, under contract 5-26555."
+ The EVLA is a facility of NRAO. operated by AUT. uuder a cooperative aereeimenut with the NSF.," The EVLA is a facility of NRAO, operated by AUI, under a cooperative agreement with the NSF."
+the UDF. the cosmie variance for dark matter decreases with redshift while for all surveys the cosmic variance for galaxies increases with redshift.,"the UDF, the cosmic variance for dark matter decreases with redshift while for all surveys the cosmic variance for galaxies increases with redshift."
+ As we have shown. cosmic variance depends strongly on the geometry of the survey.," As we have shown, cosmic variance depends strongly on the geometry of the survey."
+ In section 4.1 we investigated the geometric effect that arises because of the ratio between line-of-sight distance and the transverse distance in the redshift bins., In section \ref{sec:redshift} we investigated the geometric effect that arises because of the ratio between line-of-sight distance and the transverse distance in the redshift bins.
+ Now we investigate the effects that arise due to different angular geometry on the sky., Now we investigate the effects that arise due to different angular geometry on the sky.
+" For this we assume a survey with a fixed total area of 1.0deg""."," For this we assume a survey with a fixed total area of $1.0{~\rm
+ deg}^2$."
+" We vary the ratio between the two observation angles e,/o» froma,/as~ Oto nj/0321.", We vary the ratio between the two observation angles $\alpha_1/\alpha_2$ from $\alpha_1/\alpha_2\approx0$ to $\alpha_1/\alpha_2=1$.
+ This is done for several mean redshifts with a bin size of Az= 0.5., This is done for several mean redshifts with a bin size of $\Delta z=0.5$ .
+ Figure + shows the geometry effects on the cosmic variance., Figure \ref{fig:ratio} shows the geometry effects on the cosmic variance.
+" It plots 7,5 as a function of the ratio between the two observation angles.", It plots $\sigma_{dm}$ as a function of the ratio between the two observation angles.
+ The result agrees with our explanation in section 4.1: for a low ratio the cosmic variance Is very low., The result agrees with our explanation in section \ref{sec:redshift}: for a low ratio the cosmic variance is very low.
+ If one increases the ratio the cosmic variance quickly increases and reaches its maximum at a)ασ=I., If one increases the ratio the cosmic variance quickly increases and reaches its maximum at $\alpha_1/\alpha_2=1$.
+ Here the mean distance between the observed galaxies 1s the smallest. so that many galaxies are likely to be correlated. resulting in large cosmic variance.," Here the mean distance between the observed galaxies is the smallest, so that many galaxies are likely to be correlated, resulting in large cosmic variance."
+ This effect can also be seen by comparing the cosmic variance for GEMS and EGS at a fixed redshift bin., This effect can also be seen by comparing the cosmic variance for GEMS and EGS at a fixed redshift bin.
+ We find that although the area of the GEMS field is larger than that of the EGS. GEMS has a higher cosmic variance.," We find that although the area of the GEMS field is larger than that of the EGS, GEMS has a higher cosmic variance."
+ However. we note that unless the ratio between the observation angles is smaller than ~20%. this effect is small.," However, we note that unless the ratio between the observation angles is smaller than $\sim20\%$, this effect is small."
+ In agreement with Newman&Davis(2002).. we find that the size of the survey area is much more important than the axis ratio.," In agreement with \citet{newman2002}, we find that the size of the survey area is much more important than the axis ratio."
+ Inferences about galaxy number densities and. related quantities such as luminosity and stellar mass functions are subject to uncertainties due to cosmic variance., Inferences about galaxy number densities and related quantities such as luminosity and stellar mass functions are subject to uncertainties due to cosmic variance.
+ It can be a significant source of uncertainty. especially for surveys which cover only small areas and for massive high-redshift galaxies.," It can be a significant source of uncertainty, especially for surveys which cover only small areas and for massive high-redshift galaxies."
+ We have derived a simple recipe to compute cosmic variance for five surveys as a function of mean redshift z. redshift bin size Az and the stellar mass of the galaxy population 11," We have derived a simple recipe to compute cosmic variance for five surveys as a function of mean redshift $\bar{z}$, redshift bin size $\Delta z$ and the stellar mass of the galaxy population $m_*$."
+" For this recipe we first calculated the dark matter cosmic variance as a function of mean redshift and for a reference bin size of Az,=0.2."," For this recipe we first calculated the dark matter cosmic variance as a function of mean redshift and for a reference bin size of $\Delta
+z_{\rm red}=0.2$."
+ This was done by integrating the dark matter correlation function obtained from CDM theory for a pencil beam geometry in a redshift bin., This was done by integrating the dark matter correlation function obtained from CDM theory for a pencil beam geometry in a redshift bin.
+ We provide fitting functions for five surveys (UDF. GOODS. GEMS. AEGIS. COSMOS).," We provide fitting functions for five surveys (UDF, GOODS, GEMS, AEGIS, COSMOS)."
+ We then showed that the dependence of cosmic variance on bin size Az is nearly independent of mean redshift and survey geometry., We then showed that the dependence of cosmic variance on bin size $\Delta z$ is nearly independent of mean redshift and survey geometry.
+ This means that cosmic variance for a different bin size can be estimated from the value computed with the reference bin size using a simple conversion factor., This means that cosmic variance for a different bin size can be estimated from the value computed with the reference bin size using a simple conversion factor.
+ We have used the galaxy bias predictions by Mosteretal.(2009) (for six stellar mass bins and six thresholds) and the dark matter cosmic variance to compute the cosmic variance for galaxies., We have used the galaxy bias predictions by \citet{moster2009} (for six stellar mass bins and six thresholds) and the dark matter cosmic variance to compute the cosmic variance for galaxies.
+ We also presented a formula to compute the overall cosmic variance for a volume-weighted average of multiple separated fields., We also presented a formula to compute the overall cosmic variance for a volume-weighted average of multiple separated fields.
+ Applying this recipe to GOODS. GEMS and COSMOS at z22 and Az =0.5. the relative cosmic variance for galaxies with mw.>10''M.. Gn.~ 10M.) is ~38% (19%) for GOODS. while it is ~27% (13%) for GEMS and ~12% (6%) for COSMOS.," Applying this recipe to GOODS, GEMS and COSMOS at $\bar{z}=2$ and $\Delta z=0.5$ , the relative cosmic variance for galaxies with $m_*>10^{11}\msun$ $m_*\sim10^{10}\msun$ ) is $\sim 38\%$ $19\%$ ) for GOODS, while it is $\sim 27\%$ $13\%$ ) for GEMS and $\sim 12\%$ $6\%$ ) for COSMOS."
+ At z=3.5and Az=0.5 we found a relative cosmic variance for galaxies with i.>10''M. , At $\bar{z}=3.5$and $\Delta z=0.5$ we found a relative cosmic variance for galaxies with $m_*>10^{11}\msun$ 
+The complete invisibility of the non-lossils 1s not quite as successful for LOO kpe</?«200 kpe as we are still slightly overproducing the number of Ly:>10? L.. satellites compared to observations.,The complete invisibility of the non-fossils is not quite as successful for $100$ $<R<200$ kpc as we are still slightly overproducing the number of $L_V > 10^5$ $_\odot$ satellites compared to observations.
+" However, we are better able to reproduce the sudden steepening in the observed luminosity function in Figure 16 than with any of our other suppression mechanisms (Figures 10... 11.. and especially 12))."," However, we are better able to reproduce the sudden steepening in the observed luminosity function in Figure \ref{LF.nonf}
+ than with any of our other suppression mechanisms (Figures \ref{LF.ml500}, \ref{LF.high10}, and especially \ref{LF.sup}) )."
+ This feature may be unique to the Milky Way so we are not unduly concerned with matching it., This feature may be unique to the Milky Way so we are not unduly concerned with matching it.
+" In addition, if the are dark, our argument for the existence of primordial fossils becomes straightforward."," In addition, if the non-fossils are dark, our argument for the existence of primordial fossils becomes straightforward."
+" There are only 30 polluted fossils in this distance bin, only 75'% of the —LO observed galaxies, with any dwarf with Lyc10! L. difficult. if not impossible, to detect with current surveys."," There are only $\sim 30$ polluted fossils in this distance bin, only $75\%$ of the $\sim 40$ observed galaxies, with any dwarf with $L_V<10^4$ $_\odot$ difficult, if not impossible, to detect with current surveys."
+ Beyond the MW.3 virial radius (/?~200 kpc). turning the non-fossils dark easily places the primordial luminosity function into agreement with observations.," Beyond the MW.3 virial radius $R\sim200$ kpc), turning the non-fossils dark easily places the primordial luminosity function into agreement with observations."
+ This allows lor the formation of post-reionization populations of stars in the polluted fossils and non-fossils., This allows for the formation of post-reionization populations of stars in the polluted fossils and non-fossils.
+ We remind the reader that the >=0 halos at these radii would likely be on first approach to the Milky Way system and more likely to accrete and retain gas at later times., We remind the reader that the $z=0$ halos at these radii would likely be on first approach to the Milky Way system and more likely to accrete and retain gas at later times.
+ The diffuse primordial population in these distant non-[ossils is an observational test of star l[ormation in pre-reionizauion dwarts and the existence of pre-reionization fossils., The diffuse primordial population in these distant non-fossils is an observational test of star formation in pre-reionization dwarfs and the existence of pre-reionization fossils.
+" Although the existence of pre-reionization fossils seems likely, observations do not unequivocally demonstrate their existence due to the large uncerlainues in estimatüng the number of yet undiscovered ultra-[aint dwarfs (?).."," Although the existence of pre-reionization fossils seems likely, observations do not unequivocally demonstrate their existence due to the large uncertainties in estimating the number of yet undiscovered ultra-faint dwarfs \citep{Tollerudetal:08}."
+ In this section. we summarize three observational tests for the existence of fossils of the first galaxies that we propose based on the results in Paper I and the present work.," In this section, we summarize three observational tests for the existence of fossils of the first galaxies that we propose based on the results in Paper I and the present work."
+" The first test of our model is especially interesung as it can be performed using HST observations and does not require waiting for future all sky surveys deeper than SDSS, like PanStar or LSST. to be online."," The first test of our model is especially interesting as it can be performed using HST observations and does not require waiting for future all sky surveys deeper than SDSS, like PanStar or LSST, to be online."
+energies at the theoretical values expected for a thin-thermal plasma emission (Meweetal.1985).,energies at the theoretical values expected for a thin-thermal plasma emission \cite{Mewe1985}.
+. In this fitting process. in addition to Wea lines from highly ionized atoms. we also included some WK? lines. but in this case fixed the line energy and the ratio of the normalisation between Kea and. Wee ines at the values predicted by thin-thermal plasma models (Aleweetal.1985).," In this fitting process, in addition to $\alpha$ lines from highly ionized atoms, we also included some $\beta$ lines, but in this case fixed the line energy and the ratio of the normalisation between $\alpha$ and $\beta$ lines at the values predicted by thin-thermal plasma models \cite{Mewe1985}."
+. Similarly the fluorescent. lines from neutral ion (Ixo. I2) were fitted at fixed line energies.," Similarly the fluorescent lines from neutral iron $\alpha$, $\beta$ ) were fitted at fixed line energies."
+ As or the line width. we only allowed that of the 6.7-keV. line. which has by far the best statistics. to be free.," As for the line width, we only allowed that of the 6.7-keV line, which has by far the best statistics, to be free."
+ Phe other line widths were fixed at the values expected for a thin-thermal asma (Alewectal.1985).. taking the energy. resolution of the detector into account:. we [fixed the width of the 3.13-keV. line. which is presumably a WKe-line [rom 1elium-like argon. to be 9 eV because this line is actually a combination of resonance. forbidden. intercombination. and satellite lines.," The other line widths were fixed at the values expected for a thin-thermal plasma \cite{Mewe1985}, taking the energy resolution of the detector into account;, we fixed the width of the 3.13-keV line, which is presumably a $\alpha$ -line from helium-like argon, to be 9 eV because this line is actually a combination of resonance, forbidden, intercombination, and satellite lines."
+ With the best-fit to the MOS spectra determined. we then fitted the pn spectrum with the same mocel.," With the best-fit to the MOS spectra determined, we then fitted the pn spectrum with the same model."
+ Lhe MOS and pn results were found to be consistent with each other except for slight dillerences in the normalisation and energy scale (the latter plausibly attributable to a detector gain error)., The MOS and pn results were found to be consistent with each other except for slight differences in the normalisation and energy scale (the latter plausibly attributable to a detector gain error).
+ The pn and MOS. spectra were aligned when we applied a gain correction of —0.1 per cent to the former., The pn and MOS spectra were aligned when we applied a gain correction of $\sim$ 0.1 per cent to the former.
+ Since the calibration uncertainty is ~0.3 per cent. (Ixirsch Y02). this dilference is within the probable uncertainty.," Since the calibration uncertainty is $\sim$ 0.3 per cent (Kirsch 2002), this difference is within the probable uncertainty."
+ In 1e analysis described below the pn gain was fixed at 1.001., In the analysis described below the pn gain was fixed at 1.001.
+ Table 1 summarises the best-fitting results for re separate MOS. and. pn fits including probable line --dentifications. while Fig.," Table \ref{tbl:fit-eachline} summarises the best-fitting results for the separate MOS and pn fits including probable line identifications, while Fig."
+ 2. clisplavs the pn spectrum with 1 best-fitting model., \ref{fig:pn-fit} displays the pn spectrum with the best-fitting model.
+ The quoted uncertainties in Table and heneclorth are at the 90 per cent confidence level for one interesting parameter. unless otherwise mentioned.," The quoted uncertainties in Table \ref{tbl:fit-eachline} and henceforth are at the 90 per cent confidence level for one interesting parameter, unless otherwise mentioned."
+ In yarticular. the most. distinct four lines are identified: with Ίνα lines from Lle-like sulfur. argon (Ar). caleium (Ca). and iron.," In particular, the most distinct four lines are identified with $\alpha$ lines from He-like sulfur, argon (Ar), calcium (Ca), and iron."
+ ‘Table 1. also gives the results of a simultaneous Lit to all three detectors (but with all the line energies fixed as indicated)., Table \ref{tbl:fit-eachline} also gives the results of a simultaneous fit to all three detectors (but with all the line energies fixed as indicated).
+ The elobal normalisation [actor between the pn and MOS spectra was found to be Q.S7 (which may he accounted for at least in part by photons in the spectrum-accumulating region in the pn falling on a chip gap)., The global normalisation factor between the pn and MOS spectra was found to be 0.87 (which may be accounted for at least in part by photons in the spectrum-accumulating region in the pn falling on a chip gap).
+ We apply an appropriate correction for “the lost pn photons” in he analysis which follows., We apply an appropriate correction for “the lost pn photons” in the analysis which follows.
+ Next we estimated the ionization temperature of the xdasma from the ratio of He-like ancl L-like Ix-lines for each atomic species. using the theoretical results of Alewe 1985).," Next we estimated the ionization temperature of the plasma from the ratio of He-like and H-like K-lines for each atomic species, using the theoretical results of Mewe \shortcite{Mewe1985}."
+. Fig., Fig.
+ 3. summiarises the derived temperatures., \ref{fig:line-kt} summarises the derived temperatures.
+ The asma temperature is found to vary significantly from atom. o atom: for example. the ionization temperatures of sulfur and iron are 1 keV and 4 keV. respectively.," The plasma temperature is found to vary significantly from atom to atom; for example, the ionization temperatures of sulfur and iron are 1 keV and 4 keV, respectively."
+ Fhis implies tha he spectrum consists of multiple temperature components., This implies that the spectrum consists of multiple temperature components.
+ In fact. when we tried to apply a single-tempoerature therma model to the spectrum. it is clearly rejected.," In fact, when we tried to apply a single-temperature thermal model to the spectrum, it is clearly rejected."
+ We note that a he energv of the silicon (Si) line the continuum is strongly allected by absorption. hence the estimate of the ionization emperature of silicon may. be subject to some adcditiona systematic uncertainty.," We note that at the energy of the silicon (Si) line the continuum is strongly affected by absorption, hence the estimate of the ionization temperature of silicon may be subject to some additional systematic uncertainty."
+ On the basis of the results presented above. we have applied a tvo-tempcrature thin-thermal plasma model AL) both subject to the same absorption component to the LEPLC spectra observed. fromEast.," On the basis of the results presented above, we have applied a two-temperature thin-thermal plasma model ) both subject to the same absorption component to the EPIC spectra observed from."
+ We initially adopt. the solar abundance ratio (Anders&Crevesse1989) and. for the absorption. assume the cross sections tabulated by Morrison AeCammon (1983).," We initially adopt the solar abundance ratio \cite{Anders1989} and, for the absorption, assume the cross sections tabulated by Morrison McCammon (1983)."
+ The ratio of the normalisation between pn ancl MOS. detector. was fixed at the value obtained previously., The ratio of the normalisation between pn and MOS detector was fixed at the value obtained previously.
+ Fig., Fig.
+ 4 (upper panel) shows the pn. and AIOSI and 2 spectra with the best-fitting model. whereas," \ref{fig:simul-fit} (upper panel) shows the pn, and MOS1 and 2 spectra with the best-fitting model, whereas"
+"assuming a pure proton spectrum of the form E7-οEe, where * is the cosmic ray spectral index and E. the cut- energy.","assuming a pure proton spectrum of the form $E^{\gamma}\cdot e^{-{E\over E_c}}$, where $\gamma$ is the cosmic ray spectral index and $E_c$ the cut-off energy."
+" Figure 3 testifies that the estimated cosmic ray energy spectrum is consistent to harder than the isotropic flux (at a 4.6 o level) with a cutoff, and with the mostsignificant excess in the multi-TeV range (Abdoetal. 2008)."," Figure 3 testifies that the estimated cosmic ray energy spectrum is consistent to harder than the isotropic flux (at a 4.6 $\sigma$ level) with a cutoff, and with the mostsignificant excess in the multi-TeV range \citep{abdo2}."
+". 'The localized excess regions of multi-TeV cosmic rays cover the same portion of the sky where the tail-in excess was observed at cosmic ray energies below TeV, and have similar seasonal modulations."," The localized excess regions of multi-TeV cosmic rays cover the same portion of the sky where the tail-in excess was observed at cosmic ray energies below TeV, and have similar seasonal modulations."
+ It is likely that these are two manifestations of the same phenomenology and that the heliotail has an important role., It is likely that these are two manifestations of the same phenomenology and that the heliotail has an important role.
+" IceCube has recently reported first view of the multi-'TeV medium scale anisotropy in athe southern hemisphere (i.e. only modulations that are smaller than 60?) (Desiatietal. 2010),, which might add novel information to this observation."," IceCube has recently reported a first view of the multi-TeV medium scale anisotropy in the southern hemisphere (i.e. only modulations that are smaller than $^{\circ}$ ) \citep{beyond2010}, which might add novel information to this observation."
+" While no explanation has ever being attempted to explain the broad sub-TeV tail-in excess, a number of interpretations have been provided to address the existence of the most significant localized excess of multi-TeV cosmic rays."," While no explanation has ever being attempted to explain the broad sub-TeV tail-in excess, a number of interpretations have been provided to address the existence of the most significant localized excess of multi-TeV cosmic rays."
+ Some proposed models rely on astrophysical origin., Some proposed models rely on astrophysical origin.
+ In Salvati&Sacco it is noted that the two localized excess regions (2008)observed by MILAGRO surround the present day apparent location of Geminga pulsar., In \citet{salvati} it is noted that the two localized excess regions observed by MILAGRO surround the present day apparent location of Geminga pulsar.
+" The supernova that gave birth to the pulsar exploded about 340,000 years ago, when its distance to the Sun was estimated to be about 90 pc."," The supernova that gave birth to the pulsar exploded about 340,000 years ago, when its distance to the Sun was estimated to be about 90 pc."
+" Even if the proper motion of the pulsar induced by the explosion moved it further away (the present distance of Geminga pulsar is estimated to be about 155-35 pc), 10 TeV cosmic rays produced by the supernova have propagated about 65 pc away, if we assume Bohm from the source to here: approximately consistent with the distance of Geminga at the time of explosion."," Even if the proper motion of the pulsar induced by the explosion moved it further away (the present distance of Geminga pulsar is estimated to be about $\pm$ 35 pc), 10 TeV cosmic rays produced by the supernova have propagated about 65 pc away, if we assume Bohm from the source to here: approximately consistent with the distance of Geminga at the time of explosion."
+ From this distance a total cosmic ray energy of about 1.5.104? erg must be emitted by the supernova to produce the observed fractional excess., From this distance a total cosmic ray energy of about $1.5\cdot 10^{49}$ erg must be emitted by the supernova to produce the observed fractional excess.
+ This value is consistent with the commonly required efficiency ( 196)) with which a supernova energy output must be converted into cosmic rays if they are to maintain the galactic cosmic ray density., This value is consistent with the commonly required efficiency $\sim$ ) with which a supernova energy output must be converted into cosmic rays if they are to maintain the galactic cosmic ray density.
+" 'The major problem with this explanation is that Bohm diffusion through such large distances cannot possibly explain the localized nature of the observed excesses, but it would rather produce at most a broad faint dipolar anisotropy in arrival direction."," The major problem with this explanation is that Bohm diffusion through such large distances cannot possibly explain the localized nature of the observed excesses, but it would rather produce at most a broad faint dipolar anisotropy in arrival direction."
+" In addition, if Bohm scaling for the scattering in the immediate vicinity of the supernova shock is plausible, it seems unlikely that this regime persists during the propagation of cosmic rays through the interstellar medium (Drury&Aha-ronian2008)."," In addition, if Bohm scaling for the scattering in the immediate vicinity of the supernova shock is plausible, it seems unlikely that this regime persists during the propagation of cosmic rays through the interstellar medium \citep{drury}."
+. On the other hand the structure of the interstellar magnetic field would very hardly maintain multi-TeV cosmic rays focussed within a 10? beam., On the other hand the structure of the interstellar magnetic field would very hardly maintain multi-TeV cosmic rays focussed within a $^{\circ}$ beam.
+" The supposed opposite scenario of a free-streaming of cosmic rays along a sort of magnetic ""nozzle"" (Drury&Aha-ronian 2008), that would explain the localized nature of the observation, would also be extremely unlikely also because the propagation would have been so fast that we would not have the observation in the first place anymore."," The supposed opposite scenario of a free-streaming of cosmic rays along a sort of magnetic ""nozzle"" \citep{drury}, that would explain the localized nature of the observation, would also be extremely unlikely also because the propagation would have been so fast that we would not have the observation in the first place anymore."
+ It is possible to argue a scenario with a combination of slow Bohm diffusion regime (close to the supernova due to turbulence induced by the explosion) and fast free-streaming along magnetic field through the interstellar medium (which could partially explain the localized natures of the observed excesses)., It is possible to argue a scenario with a combination of slow Bohm diffusion regime (close to the supernova due to turbulence induced by the explosion) and fast free-streaming along magnetic field through the interstellar medium (which could partially explain the localized natures of the observed excesses).
+" But this interpretation would have problems, as the time of propagation from the source should be made long, which contradicts the idea of localization of the intensive scattering only near the source."," But this interpretation would have problems, as the time of propagation from the source should be made long, which contradicts the idea of localization of the intensive scattering only near the source."
+" Perhaps some sort of leaky magnetic field bottles are formed near the source, which make the propagation slow compared to the Bohm diffusion time, thus mitigating the problem."," Perhaps some sort of leaky magnetic field bottles are formed near the source, which make the propagation slow compared to the Bohm diffusion time, thus mitigating the problem."
+ However this possibility would require fine tuning and we have not seen this idea discussed in the literature., However this possibility would require fine tuning and we have not seen this idea discussed in the literature.
+" The coincidence of the most significant localized excess observed by MILAGRO with the heliotail, supports the idea that the heliosphere could somehow have a role."," The coincidence of the most significant localized excess observed by MILAGRO with the heliotail, supports the idea that the heliosphere could somehow have a role."
+ The possibility that we are seeing the effects of neutron production in the gravitationally focussed tail of the interstellar material was considered by Drury&Aharonian (2008)., The possibility that we are seeing the effects of neutron production in the gravitationally focussed tail of the interstellar material was considered by \citet{drury}.
+". As the Solar system surrounded by Solar Wind moves through the interstellar medium, the complex interaction between the two media create the heliotail."," As the Solar system surrounded by Solar Wind moves through the interstellar medium, the complex interaction between the two media create the heliotail."
+ Cosmic rays propagating through the direction of the tail interact with the matter and magnetic fields to produce neutrons and hence a localized excess of cosmic ray in that direction., Cosmic rays propagating through the direction of the tail interact with the matter and magnetic fields to produce neutrons and hence a localized excess of cosmic ray in that direction.
+" But while the target size has about the right size compared to the decay length of multi-TeV neutrons ( 0.1 pc), the increase of the gravitating matter density is too low to account for the observed excess."," But while the target size has about the right size compared to the decay length of multi-TeV neutrons $\sim$ 0.1 pc), the increase of the gravitating matter density is too low to account for the observed excess."
+" While it is possible to argue that the large angular scale anisotropy of cosmic rays arrival direction might be generated by a combination of astrophysical phenomena, such as the distribution of nearby recent supernova explosions (Erlykin&Wolfendale 2006),, propagation effects (Battaner, Castellano Masip 2009, Malkov et al."," While it is possible to argue that the large angular scale anisotropy of cosmic rays arrival direction might be generated by a combination of astrophysical phenomena, such as the distribution of nearby recent supernova explosions \citep{erlykin}, , propagation effects (Battaner, Castellano Masip 2009, Malkov et al."
+We estimated the NCCS ACN vield using four almost indepeudent methods (the fourth method is partially depeudenut on the estimation of the third one). each one of which is biased ciffereuth:.,"We estimated the NCCS AGN yield using four almost independent methods (the fourth method is partially dependent on the estimation of the third one), each one of which is biased differently."
+" Towever. all the estimation methods produced remarkably simular results, allowing to calculate the average NCCS AGN vield ~600+100 cauclidates."," However, all the estimation methods produced remarkably similar results, allowing to calculate the average NCCS AGN yield $\sim$ $\pm$ candidates."
+ We have shown that there is valuable new scicuce hat can be derived from the NCCS data set. apart roni inuproviug the global knowledge about the sky at Heh declinatious.," We have shown that there is valuable new science that can be derived from the NCCS data set, apart from improving the global knowledge about the sky at high declinations."
+ It is possible. therefore. to consider possible extensions of this rather Dnuited project. as defined above.," It is possible, therefore, to consider possible extensions of this rather limited project, as defined above."
+ We can relatively casily extend the surveved region five more degrees of declination. increasing the surveved region- to more than —700 doe.," We can relatively easily extend the surveyed region for five more degrees of declination, increasing the surveyed region to more than 700 $^2$."
+JD This will increase the overlap with SDSS-surveyed regions. allowing an nproved cross-calibration between the two surveys.," This will increase the overlap with SDSS-surveyed regions, allowing an improved cross-calibration between the two surveys."
+ Alternatively. we could repeat the observations of the already. surveved region in oue of the R aud I bauds to check the variability of the detected sources.," Alternatively, we could repeat the observations of the already surveyed region in one of the R and I bands to check the variability of the detected sources."
+ We could meree the NCCS data with 2MASS extcuding the waveleneth base to IR., We could merge the NCCS data with 2MASS extending the wavelength base to IR.
+" Though 2MASS is relatively shallow. it could provide additional information about NCCSob jects brighter than ~16"" 17""."," Though 2MASS is relatively shallow, it could provide additional information about NCCS objects brighter than $\sim$ $^m$ $^m$."
+"Spectroscopicl follow-up of the promising NCCS sources could be performed at the Wise Observatory for objects brighter than ~16"" or at a larger telescope for the faint oues.", Spectroscopic follow-up of the promising NCCS sources could be performed at the Wise Observatory for objects brighter than $\sim$ $^m$ or at a larger telescope for the faint ones.
+ Based ou observations made with the NASA Calaxy Evolution Explorer., Based on observations made with the NASA Galaxy Evolution Explorer.
+ CALEX is operated for NASA by the California Institute of Technology under NASA coutract NAS5-08031., GALEX is operated for NASA by the California Institute of Technology under NASA contract NAS5-98034.
+ This study males use of data from the (http:/Awww.scdss.ore/collaboration/credits.htiul)., This study makes use of data from the (http://www.sdss.org/collaboration/credits.html).
+and actually interpret observed events as constraints on the actual compact object binary population. as is our goal.,"and actually interpret observed events as constraints on the actual compact object binary population, as is our goal."
+ To improve on past models for the physical population of inspiraline compact binary svslenis. we use galaxy catalogs to model the actual distribution of galaxies in the local universe and we use stellar synthesis calculations (specilicallythoseofBelezvuski.&Dulik2002) (o model the mass distribution of binaries within each galaxy.," To improve on past models for the physical population of inspiraling compact binary systems, we use galaxy catalogs to model the actual distribution of galaxies in the local universe and we use stellar synthesis calculations \cite[specifically those of ][]{belc02} to model the mass distribution of binaries within each galaxy."
+ From the constructed. population models. we determine the compact binary coalescence rate and distribution with binary svstem mass (hat we expect the LIGO detector svstem to observe. taking full account of each galaxys distance and declination. the LIGO detector svstems noise spectrum. ancl its position ancl orientation on Earth.," From the constructed population models, we determine the compact binary coalescence rate and distribution with binary system mass that we expect the LIGO detector system to observe, taking full account of each galaxy's distance and declination, the LIGO detector system's noise spectrum, and its position and orientation on Earth."
+ Our principal goal in relating a physical population model to the distribution that we expect modern GW detectors to observe is to enable observations by those detectors to constrain the population model (see Dulik&Belezvuski2003. and Bulik.Belezvnski.&Rucak2004. [or recent studies with similar goals)., Our principal goal in relating a physical population model to the distribution that we expect modern GW detectors to observe is to enable observations by those detectors to constrain the population model (see \citealt{bulik} and \citealt{bulik2} for recent studies with similar goals).
+ Through the calculations described here. comparison of future observed rates or rate upper limits constrain stellar svnthesis models aud the overall binary. compact object population.," Through the calculations described here, comparison of future observed rates or rate upper limits constrain stellar synthesis models and the overall binary compact object population."
+ While our principal interest is in preparing for this kind of interpretation of fortheoming; observations. as a byv-product of our investigations we have improved detection rate predictions as well.," While our principal interest is in preparing for this kind of interpretation of forthcoming observations, as a by-product of our investigations we have improved detection rate predictions as well."
+ In relsec:backeround we present an overview of the various approaches used so [ar lor the extrapolation of Galactic detection rates to extragalactic distances and introduce our novel ealaxv-by-galaxy. approach. whereby we calculate the detectability of DC inspiral for each ealaxv in our catalog.," In \\ref{sec:background} we present an overview of the various approaches used so far for the extrapolation of Galactic detection rates to extragalactic distances and introduce our novel galaxy-by-galaxy approach, whereby we calculate the detectability of BCO inspiral for each galaxy in our catalog."
+ In relseczmethod we describe how we calculate. from the detailed extra-galactie population model described in re[secibackground.. the observed distribution of BCOs.," In \\ref{sec:method} we describe how we calculate, from the detailed extra-galactic population model described in \\ref{sec:background}, the observed distribution of BCOs."
+ In relsec:resulis we discuss our results. including the LIGO detector svstem's efliciency [or detecting binaries from dilferent galaxies in (he nearby universe. (he expected observed coalescing binary mass distribution. new detection rate predictions. and the implications of the geography of the nearby univese for detection of binary compact object svstems.," In \\ref{sec:results}
+ we discuss our results, including the LIGO detector system's efficiency for detecting binaries from different galaxies in the nearby universe, the expected observed coalescing binary mass distribution, new detection rate predictions, and the implications of the geography of the nearby univese for detection of binary compact object systems."
+ We end in with a summary of our main conclusions., We end in \\ref{sec:conclusions} with a summary of our main conclusions.
+and so corresponding to LEcldington approximation /=34A (c.g.Alihalas&Weibel-Alibalas1999).,and so corresponding to Eddington approximation $J=3K$ \citep[e.g.\][]{mihal99}.
+.. Thus given a distribution of ionization. we can solve for the cilfuse Ποιά using this expression. in which the cilfusive transport corrects the OLS intensity.," Thus given a distribution of ionization, we can solve for the diffuse field using this expression, in which the diffusive transport corrects the OTS intensity."
+ Note. that since ο~1 in the centre of the region.the free path 1/58 may be Iarge.," Note, that since $x\sim 1$ in the centre of the region,the free path $1/\kappa$ may be large."
+ lützerveld(2005). solves an outward-only svsteni where for slab symmetry d— 1. £=1. for evlindrical svmmetee d=2. €=2s and for spherical symmetry d=3. €=dx.," \cite{ritze05} solves an outward-only system where for slab symmetry $d=1$ , $\xi=1$, for cylindrical symmetry $d=2$, $\xi = 2\pi$ and for spherical symmetry $d=3$, $\xi = 4\pi$."
+ Note that these equations are given here or the total ionization intensity integrated over the surface at the specified radius., Note that these equations are given here for the total ionization intensity integrated over the surface at the specified radius.
+ Ritzervelcl states that they apply to he density of photons per unit volume. but this is clearly not the case. as if all the recombination cocllicicnts are zero. then Lay=cons!.," Ritzerveld states that they apply to the density of photons per unit volume, but this is clearly not the case, as if all the recombination coefficients are zero, then $L_{\rm dir} = L_{\rm dif} = {\it const}$."
+ This system approximates he transport in the limit of nearly complete ionization within the region., This system approximates the transport in the limit of nearly complete ionization within the region.
+" The factor e=eoe, allows for the referential absorption of the dilfuse photons.", The factor $c=a_0/a_1$ allows for the preferential absorption of the diffuse photons.
+ The results presented by. Hitzerveld.(2005) can. be eeneralized to the case of arbitrary ratios of absorption cross section., The results presented by \cite{ritze05} can be generalized to the case of arbitrary ratios of absorption cross section.
+ Adding equations (16)) and (17)). we find and hence A general (implicit) expression for the direct and diffuse fields can be derived. in terms of the total field. independent of the density distribution or geometry.," Adding equations \ref{e:ritzerdir}) ) and \ref{e:ritzerdif}) ), we find and hence A general (implicit) expression for the direct and diffuse fields can be derived in terms of the total field, independent of the density distribution or geometry."
+ Dividing equation (18)) by equation (16)). we find and hence Given Ly1} from equation (19)) above. equation (21)) may then be solved for Lai (explicitly in the case e=1. to recover the analytic solutions which Ritzervelcl presents) and then Lay=ListLas.," Dividing equation \ref{e:ritzersum}) ) by equation \ref{e:ritzerdir}) ), we find and hence Given $L_{\rm tot}(r)$ from equation \ref{e:itot}) ) above, equation \ref{e:gendirect}) ) may then be solved for $L_{\rm dir}$ (explicitly in the case $c=1$, to recover the analytic solutions which Ritzerveld presents), and then $L_{\rm dif} = L_{\rm tot}-L_{\rm
+dir}$."
+ We use an iterative technique to find solutions. using similar techniques to those applied. clsewhere (Hammer&Seaton1963:Rubin LOGS).," We use an iterative technique to find solutions, using similar techniques to those applied elsewhere \citep{humme63,rubin68}."
+.. Llere. we integrate the cilfuse field using a large number of coaxial zones.," Here, we integrate the diffuse field using a large number of coaxial zones."
+ We first derive an approximate ionization structure using the case B approximation. and. use this to estimate the outward-going diluse field. components.," We first derive an approximate ionization structure using the case B approximation, and use this to estimate the outward-going diffuse field components."
+ We integrate the diffuse field. transport inwards through this structure., We integrate the diffuse field transport inwards through this structure.
+ As part of this sweep. we calculate an improved. estimate for the ionization structure taking account of the new inward radiation field. beams.," As part of this sweep, we calculate an improved estimate for the ionization structure taking account of the new inward radiation field beams."
+ We then integrate the clirect radiation and outward-going cdilluse field outward: through the resulting structure. using the inware diffuse field at its closest point to the star as the initial condition for the outward beams. updating the ionization structure again.," We then integrate the direct radiation and outward-going diffuse field outward through the resulting structure, using the inward diffuse field at its closest point to the star as the initial condition for the outward beams, updating the ionization structure again."
+ We iterate these inward and outward sweeps until the ionization structure has. converged this convergence is rapid in practice., We iterate these inward and outward sweeps until the ionization structure has converged – this convergence is rapid in practice.
+ The results. for the ionization structure are very sensitive to the details of the numerical scheme for the case where the density varies as rE in particular the Strómmegren radius can vary significantly.," The results for the ionization structure are very sensitive to the details of the numerical scheme for the case where the density varies as $r^{-2}$, in particular the Strömmgren radius can vary significantly."
+ As noted by Francoetal.(1990).. the ionization front can in this case escape to infinity: having the ionization front at 20 times the inner radius of the density. distribution requires a rather precisely tuned: value of the incident. radiation field.," As noted by \cite{franc90}, the ionization front can in this case escape to infinity: having the ionization front at 20 times the inner radius of the density distribution requires a rather precisely tuned value of the incident radiation field."
+" ""Ehis feeds through to the numerical sensitivity.", This feeds through to the numerical sensitivity.
+ As a result of this. considerable care has to be taken in the radiation field integration to ensure the rate at. which photons are removed. from the radiation field is consisten with the ionization rate within the zone (Abeletal.1999:2006).," As a result of this, considerable care has to be taken in the radiation field integration to ensure the rate at which photons are removed from the radiation field is consistent with the ionization rate within the zone \citep{abele99,willi02,melle06,whale06}."
+. To do this. we will adapt the approach clescribec by Williams(2002) to the case of multiple incident beans of radiation aud. dilfuse sources.," To do this, we will adapt the approach described by \cite{willi02} to the case of multiple incident beams of radiation and diffuse sources."
+ We specify the radiation field in coaxial cvlindrica zones. which intersect spherical shells within which the physical variables are assumed to be uniform.," We specify the radiation field in coaxial cylindrical zones, which intersect spherical shells within which the physical variables are assumed to be uniform."
+ We take the Ath component of the density. [iek to occupy a spherica shell of outer radius 7;=fA. and t1e j-th group of clilluse photons to be those with an impact parameter between 251 and Ay. where Ry=ja.," We take the $i$ -th component of the density field to occupy a spherical shell of outer radius $r_i = i\Delta$, and the $j$ -th group of diffuse photons to be those with an impact parameter between $R_{j-1}$ and $R_j$, where $R_j = j\Delta$."
+ Integrating equation (3)) over a spatial zone. we have where Substituting from the transport equations. (1)) and (2)). then gives where the problem is uniform in space. so the angular integral over the zone volume is trivial.," Integrating equation \ref{e:ion}) ) over a spatial zone, we have where Substituting from the transport equations, \ref{e:dir}) ) and \ref{e:dif}) ), then gives where the problem is uniform in space, so the angular integral over the zone volume is trivial."
+ Combining the spatial integral i7dr with the beam solid angle integral dQ’. however. vields an integralwhich may be taken as over the zone volume for beams which are incident from a single direction.," Combining the spatial integral $r^2 dr$ with the beam solid angle integral $d\Omega'$ however, yields an integralwhich may be taken as over the zone volume for beams which are incident from a single direction."
+ This may be cdiscretized as the cüllerence in cülfuse Hux within each radiation bin across the zone.," This may be discretized as the difference in diffuse flux within each radiation bin across the zone,"
+spectra of many quasars. and (4) assuming no beamine. the characteristic huninosities and abundances of bright quasars near z~2.5 are consistent with (he characteristic masses and abundances of (heir remnant DIIs at 2=0 (see. e.g. Yu Tremaine 2002): strong beaming would likely invalidate (his successful agreement.,"spectra of many quasars, and (4) assuming no beaming, the characteristic luminosities and abundances of bright quasars near $z\sim 2.5$ are consistent with the characteristic masses and abundances of their remnant BHs at $z=0$ (see, e.g. Yu Tremaine 2002); strong beaming would likely invalidate this successful agreement."
+ Combining this with the finding above that the quasar is not magnified through eravitational lensing by a signilicant [actor implies that SDSS 10304-0524 needs to indeed contain a very massive black hole., Combining this with the finding above that the quasar is not magnified through gravitational lensing by a significant factor implies that SDSS 1030+0524 needs to indeed contain a very massive black hole.
+ Assuming the quasar radiates al the Edcdington Iuminositv and magnilied through. gravitational lensing bv a factor of 5. the minimum mass for its resident DIL is 1x105NM. and it then has to have formed at redshift 2>9.," Assuming the quasar radiates at the Eddington luminosity and magnified through gravitational lensing by a factor of $5$, the minimum mass for its resident BH is $4\times10^8~{\rm
+M_\odot}$, and it then has to have formed at redshift $z>9$."
+ We have analvzed the flux distribution of the emission of (he quasar wil the highest known redshift. SDSS 10304-0524 at 2=6.28. discovered by the Sloan Digital Sky Survey.," We have analyzed the flux distribution of the emission of the quasar with the highest known redshift, SDSS 1030+0524 at $z=6.28$, discovered by the Sloan Digital Sky Survey."
+ From its specirunm. we infer (he presence of a large (4.5Mpc) ionized region around this QSO.," From its spectrum, we infer the presence of a large $\sim
+4.5$ Mpc) ionized region around this QSO."
+ The large size of (his ionized region makes it impossible for Chis source to be intrinsically faint. or to be verv voung.," The large size of this ionized region makes it impossible for this source to be intrinsically faint, or to be very young."
+ We find that SDSS 10302-0524 could not have been magnified through gravitational lensine by more than a [actor of ~5., We find that SDSS 1030+0524 could not have been magnified through gravitational lensing by more than a factor of $\sim5$.
+ The line/continuumn ratio of 5D55 10304-0524 is observed to be (vice that of the median 2>2.25 quasar. indicating (hat (his quasar is also unlikely to be significantly. beamed.," The line/continuum ratio of SDSS 1030+0524 is observed to be twice that of the median $z>2.25$ quasar, indicating that this quasar is also unlikely to be significantly beamed."
+ Combining (hese two lactsM., Combining these two facts.
+".. If the quasar is not lensed. ancl is shining at (or below) the Exldington luninosity of its resident SMDIIL. then the inferred niass is at least Mj,=2xLO?\L.. and its minimum age is 2x10* ves."," If the quasar is not lensed, and is shining at (or below) the Eddington luminosity of its resident SMBH, then the inferred mass is at least $M_{\rm bh}=2\times10^9~{\rm M_\odot}$, and its minimum age is $2\times 10^{7}~$ yrs."
+" These numbers can only be modified by gravitational lensing bv relatively small factors: if the source is magnified by the maximum allowed [actor of 5. the BIL mass is ~4x105M... and in this case. its age has to be longer than 10 yrs. placing its formation redshift at 2,>9."," These numbers can only be modified by gravitational lensing by relatively small factors: if the source is magnified by the maximum allowed factor of $5$, the BH mass is $\sim 4\times10^8~{\rm
+M_\odot}$, and in this case, its age has to be longer than $10^8~$ yrs, placing its formation redshift at $z_f>9$."
+ As the formation of such massive black holes in the universe al such high redshifts is already presenting a theoretical challenge. it is important to have limits on the maenilicalion of their fInxes by gravitational lensing.," As the formation of such massive black holes in the universe at such high redshifts is already presenting a theoretical challenge, it is important to have limits on the magnification of their fluxes by gravitational lensing."
+ Although SDSS 10302-0524 is eurrentlv the only high redshift quasar to which our method is applicable. (he constraints we have derived for this source can be repeated and applied to future quasars that will be discovered at 26.3. prior to the reionization epoch.," Although SDSS 1030+0524 is currently the only high redshift quasar to which our method is applicable, the constraints we have derived for this source can be repeated and applied to future quasars that will be discovered at $z\ga 6.3$, prior to the reionization epoch."
+ We thank Michael Strauss and Dan Vanden Berk for providing the spectrum of SDSS 10304-0524. and the mean SDSS quasar spectrum. in electronic form. together with helpful narratives.," We thank Michael Strauss and Dan Vanden Berk for providing the spectrum of SDSS 1030+0524, and the mean SDSS quasar spectrum, in electronic form, together with helpful narratives."
+ ZIT acknowledges support by NASA through the Hubble Fellowship grant, ZH acknowledges support by NASA through the Hubble Fellowship grant
+ ZIT acknowledges support by NASA through the Hubble Fellowship grant., ZH acknowledges support by NASA through the Hubble Fellowship grant
+About a decade ago estimates of the percentage of planet-hosting solar-like stars stood al ~ICS3%. although it was already suspected| that the percentage might be much higher for stars with higher metallicities: ~25%—30% for stars with twice the solar metallicity (Fe/ll> >0.3,"About a decade ago estimates of the percentage of planet-hosting solar-like stars stood at $\sim 3 \%$, although it was already suspected that the percentage might be much higher for stars with higher metallicities: $\sim 25\%-30\%$ for stars with twice the solar metallicity \citep[$\rm{Fe/H}> >0.3;."
+":??). ?. extrapolated from the detected parameter space of planetary svslenms {ο below detection sensilivily aud concluded that the real fraction of planet hosting Sun-like stars should be at least [or My,sini>0.92.4, and P?«€13 vears and at least [or M,sin;>OLA) and P«60 vis.", \citet{Lineweaver2003} extrapolated from the detected parameter space of planetary systems to below detection sensitivity and concluded that the real fraction of planet hosting Sun-like stars should be at least for $M_p \sin i > 0.3M_{\rm J}$ and $P < 13$ years and at least for $M_p \sin i > 0.1M_{\rm J}$ and $P < 60$ yrs.
+ Thev also suggested that since this area of the planet mass-period plane is only. a fraction of that occupied by our own Solar System. these fraction may be even larger once the entire parameter space is sampled.," They also suggested that since this area of the planet mass-period plane is only a fraction of that occupied by our own Solar System, these fraction may be even larger once the entire parameter space is sampled."
+ This was (he situation when Soker Subag (2005) updated the PN formation channel statistics (See Column 5 of Table 1)., This was the situation when Soker Subag (2005) updated the PN formation channel statistics (See Column 5 of Table 1).
+ Today we have strong evidence that (he planetary Iraction increases with bot metallicity and mass of the host star (?).. although our knowledge of the planet-hosting star [raction as a [funetion of planet mass ancl orbital separation is grosslv incomplete.," Today we have strong evidence that the planetary fraction increases with both metallicity and mass of the host star \citep{Johnson2010}, although our knowledge of the planet-hosting star fraction as a function of planet mass and orbital separation is grossly incomplete."
+ A new finding in this respect is that of 2? who determined that ~262% of stars having a main sequence mass of 1.5<M/M. <2.0 host massive planets atseparations (but still less than 3 AU)., A new finding in this respect is that of \citet{Bowler2010} who determined that $\sim 26^{+9}_{-8}\%$ of stars having a main sequence mass of $1.5 \la M/$ $ \la 2.0$ host massive planets at (but still less than 3 AU).
+ The new finding is not only of a larger Iraction of planet-hosting stars. but also puts the planets at larger orbital separations. where (hey can more easily be available to shape winds from AGB stars: if the planets are too close to their host. thev interact with the star during the RGB phase. aud either prevent the star [rom reaching the AGB at all (ancl form a PN). or they just get destroved before the star reaches the AGB (???)..," The new finding is not only of a larger fraction of planet-hosting stars, but also puts the planets at larger orbital separations, where they can more easily be available to shape winds from AGB stars: if the planets are too close to their host, they interact with the star during the RGB phase, and either prevent the star from reaching the AGB at all (and form a PN), or they just get destroyed before the star reaches the AGB \citep{Nelemans1998,Villaver2007b,Villaver2009}."
+ There is also an inclication [rom (he analvsis of a handful of thick disk stars that the planetary fraction for low metallicity. may be higher than for vounger stars at similar metallicities (?)..," There is also an indication from the analysis of a handful of thick disk stars that the planetary fraction for low metallicity, may be higher than for younger stars at similar metallicities \citep{Sheehan2010}."
+ Another relevant finding is that of brown dwar! ancl planetary mass companions around extreme horizontal branch (ENB) stars. also termed sdO and sdB subcdwarfs.," Another relevant finding is that of brown dwarf and planetary mass companions around extreme horizontal branch (EHB) stars, also termed sdO and sdB subdwarfs."
+ Although such [incdings regard stars that have gone through the RGB.nef the AGB evolution. we can draw information of a general interest.," Although such findings regard stars that have gone through the RGB, the AGB evolution, we can draw information of a general interest."
+ EIID stars are on the horizontal branch (IB) and have small envelope mass. because their RGB progenitor lost most of its envelope (?)..," EHB stars are on the horizontal branch (HB) and have small envelope mass, because their RGB progenitor lost most of its envelope \citep{D'Cruz1996}."
+ Current consensus is (hat inmost cases the mass loss was caused or enhanced by (he interaction with a slellar-mass Companion (?).., Current consensus is that inmost cases the mass loss was caused or enhanced by the interaction with a stellar-mass companion \citep{Han2002}. .
+Here we report WEPC2 observatious of the Quaoar-Meywot. Ixuiper elt binary.,Here we report WFPC2 observations of the Quaoar-Weywot Kuiper belt binary.
+ From tlese observalions we [iuc that Weswot is ou au elliptical orbit with eccentricity of 0.114- 0.01. pe3od of 12.[38d-0.005 days. aud a semi-major axis of 1.15-0.08 x 10! kin.," From these observations we find that Weywot is on an elliptical orbit with eccentricity of 0.14 ± 0.04, period of 12.438±0.005 days, and a semi-major axis of 1.45±0.08 × 10^4 km."
+ Tje orbi treyreals a surpsiugly high Quaoar-Weywot system mass of 1.64E0.3 x 10?! ke., The orbit reveals a surpsingly high Quaoar-Weywot system mass of 1.6±0.3 × 10^21 kg.
+ Using tle stl[ace properties of the Uraniau aud Neptuuian satellites as a proxy for Quaoal’s surface. we reanalvze the size estimate trom Brown aud Trujillo (2001).," Using the surface properties of the Uranian and Neptunian satellites as a proxy for Quaoar's surface, we reanalyze the size estimate from Brown and Trujillo (2004)."
+ We Litd. [rou a uealt of available published size estimates. a diameter lor Quaoar of 8004-TO lan.," We find, from a mean of available published size estimates, a diameter for Quaoar of 890±70 km."
+" We find Quaoar's deusity to be p= 1.24ειοcm"". possibly he highest density in tle ltiper belt. lxu"," We find Quaoar's density to be ρ = 4.2±1.3, possibly the highest density in the Kuiper belt. \end{abstract}\tikzmark{mainBodyEnd121}
+
+%%% BEGIN DOCUMENT
+\tikzmark{mainBodyStart122}\begin{document}"
+iper belt Objects: Iucdividual Quaoar. general ]xuiper belt objecs are icy members of a planetesiual population of our Solar system beyond Neptune.," Kuiper belt Objects: Individual Quaoar, general Kuiper belt objects are icy members of a planetesimal population of our Solar system beyond Neptune."
+" To date. these bodies have appeared to [αἱ mainly into two density Classes: the si1all IXBOs have low densiies, pSLecanaloE. sugeeMODsting that these objects are largely compose of ices (Margotetal.2001:Stausberry2006:Grundy2007.2008)."," To date, these bodies have appeared to fall mainly into two density classes: the small KBOs have low densities, $\rho \lesssim 1 \dense$, suggesting that these objects are largely composed of ices \citep{Margot2004,Stansberry2006,Grundy2007,Grundy2008}."
+. The largest objects ineasured Lave densities Pn>Dee[n]|ο (Buieetal.1997:Brown2006) with the hielesl observed deusity. tliat of Hatunea. being p2.6&cm! (Babinowitzetal.2007) Sugeooesting a hieher rock coitent. bu still atjoderate fraction of ice i1 the body.," The largest objects measured have densities $\rho \gtrsim 2 \dense$ \citep{Buie1997,Brown2006} with the highest observed density, that of Haumea, being $\rho~2.6 \dense$ \citep{Rabinowitz2006,Lacerda2007} suggesting a higher rock content, but still a moderate fraction of ice in the body."
+ ΤΙese two density classes suggest separate formatio1 patli-ways resulting iu high ceusities for the largest objects. and low deusijes for smaller objects.," These two density classes suggest separate formation path-ways resulting in high densities for the largest objects, and low densities for smaller objects."
+ The discovery of Quaoar's satelite. Wevwot (Brown&Suer2007) provices all opportuui to determine the mass and density of a Ixuiper belt object smaller than the large. high. deusi objects. but larger than the small oyjects which have low densities.," The discovery of Quaoar's satellite, Weywot \citep{Brown2007IAUC} provides an opportunity to determine the mass and density of a Kuiper belt object smaller than the large, high density objects, but larger than the small objects which have low densities."
+ To that etd. the goal of tus work was to determiue Weywot's orbit. aud Quaoars taass auc cleusity.," To that end, the goal of this work was to determine Weywot's orbit, and Quaoar's mass and density."
+ Here wer'eport observatic olf Quaoar using the Wide-Field Planetary Camera 2 (WEPC2) aboard the Hubbe Space Telesco[, Here we report observations of Quaoar using the Wide-Field Planetary Camera 2 (WFPC2) aboard the Hubble Space Telescope.
+ Iu Section 2 we report the observations aud procedure for rieasuring Weywots orbit., In Section \ref{sec:Observations} we report the observations and procedure for measuring Weywot's orbit.
+ Iu Section , In Section \ref{sec:size} 
+o one at a higher frequency.,to one at a higher frequency.
+ Therefore. the value of the spectral index will also vary with time. being positive when he emission is decreasing and negative when it is increasing.," Therefore, the value of the spectral index will also vary with time, being positive when the emission is decreasing and negative when it is increasing."
+ The measurement of a positive spectral index in 2003 and a negative one in 2009 strengthens then the interpretation of he observed variability as due to an unstable stellar wind., The measurement of a positive spectral index in 2003 and a negative one in 2009 strengthens then the interpretation of the observed variability as due to an unstable stellar wind.
+ As shown in Panagia&Felli(1975) if the radio Hux is due ο a stellar wind. we can calculate its mass loss rate as and the size of the emitting region as where cosmic abundances. full ionisation. and a 107 Ix emperature of the electron gas have been assumed.," As shown in \citet{panafelli}, if the radio flux is due to a stellar wind, we can calculate its mass loss rate as and the size of the emitting region as where cosmic abundances, full ionisation, and a $^4$ K temperature of the electron gas have been assumed."
+" I we ink the racio Hux to the fast wind component. we can take 100 km Lows its terminal velocity and. calculate from the racio [Lux density in 2001 (when t16 spectral index was likely »ositive) M—L610"" M.  and Reven,~Lot cm (~ 07.01). assuming a distance of 5.8 kpe."," If we link the radio flux to the fast wind component, we can take 100 km $^{-1}$ as its terminal velocity and calculate from the radio flux density in 2001 (when the spectral index was likely positive) $\dot{\mathrm{M}}\sim 1.6\times10^{-6}$ $_\odot$ $^{-1}$ and $_{8.4 \,\mathrm{GHz}}\sim 10^{15}$ cm $\sim0''$ .01), assuming a distance of 5.8 kpc."
+ The mass loss rate found is quite high for a postACD star. where values smaller than T M. C loasare expected.," The mass loss rate found is quite high for a post-AGB star, where values smaller than $^{-7}$ $_\odot$ $^{-1}$ are expected."
+- evertheless. large rates (10.7 7M. D) have been estialec in the prototypical pre-PN CRL 618. which has undergone at least two distinct οjXsodes Of mass loss in the orm of a slow wind. in the last 2500 vr (Sánchez-Contrerasetal 2004).," Nevertheless, large rates $^{-5}$ $^{-4}$ $_\odot$ $^{-1}$ ) have been estimated in the prototypical pre-PN CRL 618, which has undergone at least two distinct episodes of mass loss in the form of a slow wind, in the last 2500 yr \citep{sanchez}."
+.. Also. enhanced mass loss in LRAS 17423-1755 las been ound by Llugeinsetal(2004).. who have estimated hat its molecular envelope (70.6 M.) was ejected in less han 1500 vr at a mass loss rate o>lot M. we+.," Also, enhanced mass loss in IRAS 17423-1755 has been found by \citet{huggins}, who have estimated that its molecular envelope $\sim$ 0.6 $_\odot$ ) was ejected in less than 1500 yr at a mass loss rate $>10^{-4}$ $_\odot$ $^{-1}$."
+" Lhe estimation of the linear size of the emitting region. allows us to improve the calculation of the emission measure and convert it into an electron density of 10"" em", The estimation of the linear size of the emitting region allows us to improve the calculation of the emission measure and convert it into an electron density of $\times10^6$ $^{-3}$.
+ If weassume that. like in CRL GIs. the radio lux arises. [rom⋅ a slow stellar wind. (15 km 1 7). its mass loss rate becomes 2.40107 M. to with Rea approximately unchanged.," If weassume that, like in CRL 618, the radio flux arises from a slow stellar wind (15 km $^{-1}$ ), its mass loss rate becomes $2.4\times10^{-7}$ $_\odot$ $^{-1}$ , with $_{8.4}$ approximately unchanged."
+apparent at intermediate inclinations but still exists towards an edge-on view.,apparent at intermediate inclinations but still exists towards an edge-on view.
+ The additional scattering region also has an impact on the rotation Aw between the soft and hard X-ray polarisation angle., The additional scattering region also has an impact on the rotation $\Delta \psi$ between the soft and hard X-ray polarisation angle.
+ The equatorial wedge covers a fraction of the primary radiation going towards the inner surfaces of the torus and it changes the irradiation geometry for both the torus and the polar cones., The equatorial wedge covers a fraction of the primary radiation going towards the inner surfaces of the torus and it changes the irradiation geometry for both the torus and the polar cones.
+" These effects diminish the resulting Aw, at least when the optical depth in the polar outflows is high."," These effects diminish the resulting $\Delta \psi$, at least when the optical depth in the polar outflows is high."
+ We have investigated different half-opening angles of the pure electron scattering wedge as well as a reduction of its optical depth., We have investigated different half-opening angles of the pure electron scattering wedge as well as a reduction of its optical depth.
+" The obtained values of Aw change for such cases, but the qualitative behaviour stays the same."," The obtained values of $\Delta \psi$ change for such cases, but the qualitative behaviour stays the same."
+ We have conducted accurate modelling of the expected X-ray polarisation induced by complex reprocessing in the active nucleus of NGC 1068., We have conducted accurate modelling of the expected X-ray polarisation induced by complex reprocessing in the active nucleus of NGC 1068.
+ This work is motivated by the apparent misalignment of the ionisation cones with respect to the torus axis as discussed by Rabanetal.(2009)., This work is motivated by the apparent misalignment of the ionisation cones with respect to the torus axis as discussed by \citet{raban2009}.
+ The orientation of the ionisation cones was inferred from slit spectroscopy of the narrow line region (NLR) using the STIS in combination with kinetic modelling (Dasetal. 2006)., The orientation of the ionisation cones was inferred from slit spectroscopy of the narrow line region (NLR) using the in combination with kinetic modelling \citep{das2006}.
+ But Rabanetal.(2009) also point out that this orientation is still a matter of debate., But \citet{raban2009} also point out that this orientation is still a matter of debate.
+" From the observation of infrared emission lines, Poncelet, Sol Perrin (2008) deduce a different position angle of the NLR that is more in agreement with theHST imaging results of [OIII] emission reported by Evansetal.(1991)."," From the observation of infrared emission lines, Poncelet, Sol Perrin (2008) deduce a different position angle of the NLR that is more in agreement with the imaging results of [OIII] emission reported by \citet{evans1991}."
+. But then the UV-imaging probably suffers from significant foreground absorption., But then the UV-imaging probably suffers from significant foreground absorption.
+" In addition to that, we want to point out that all these measurements are taken on large spatial scales reaching out to 100 parsec from the central engine."," In addition to that, we want to point out that all these measurements are taken on large spatial scales reaching out to 100 parsec from the central engine."
+ It is not straightforward to derive from these observations what the actual geometry of the ionisation cones across a few parsec is., It is not straightforward to derive from these observations what the actual geometry of the ionisation cones across a few parsec is.
+ One could imagine that the outflows become deflected or twisted on intermediate scales and thus their orientation would vary with distance., One could imagine that the outflows become deflected or twisted on intermediate scales and thus their orientation would vary with distance.
+" As Rabanetal.(2009) conclude, further investigation is needed to really constrain the geometry of the innermost outflows."," As \citet{raban2009} conclude, further investigation is needed to really constrain the geometry of the innermost outflows."
+" In this context, X-ray polarimetry is going to give unambiguous constraints."," In this context, X-ray polarimetry is going to give unambiguous constraints."
+" The scattered X-ray emission must come from very close to the central engine, as is indicated from X-ray spectroscopy results revealing the presence of high-ionisation emission lines of iron (Marshalletal.1993;Matt 2004)."," The scattered X-ray emission must come from very close to the central engine, as is indicated from X-ray spectroscopy results revealing the presence of high-ionisation emission lines of iron \citep{marshall1993,matt2004}."
+". Furthermore, any foreground absorption effects seen in the UV should much less interfere in the X-ray range so that X-ray polarimetry enables a direct view on the innermost parts of the outflow."," Furthermore, any foreground absorption effects seen in the UV should much less interfere in the X-ray range so that X-ray polarimetry enables a direct view on the innermost parts of the outflow."
+" Our modelling results have shown that the polarisation position angle in the soft X-ray range is clearly dominated by the polar scattering, which enables a precise measurement of the orientation of the scattering material."," Our modelling results have shown that the polarisation position angle in the soft X-ray range is clearly dominated by the polar scattering, which enables a precise measurement of the orientation of the scattering material."
+ Another aim of this investigation is to verify if the rotation Aw of the X-ray polarisation angle between soft and hard X-ray energies can be used to measure the misalignment between the polar scattering regions and the obscuring torus., Another aim of this investigation is to verify if the rotation $\Delta \psi$ of the X-ray polarisation angle between soft and hard X-ray energies can be used to measure the misalignment between the polar scattering regions and the obscuring torus.
+ We have found that coherent modelling of the different scattering components leads to values of Av that at heavily obscured viewing directions allows one to distinguish the torus from the ionisation cones., We have found that coherent modelling of the different scattering components leads to values of $\Delta \psi$ that at heavily obscured viewing directions allows one to distinguish the torus from the ionisation cones.
+ The relation between Aw and the actual misalignment also depends on i., The relation between $\Delta \psi$ and the actual misalignment also depends on $i$.
+" The highest values of Aw, which are most favourable for future X-ray polarimetry observations, are found at moderate viewing angles."," The highest values of $\Delta \psi$, which are most favourable for future X-ray polarimetry observations, are found at moderate viewing angles."
+" At strictly edge-on viewing directions, Aw becomes very small, especially for higher optical depths of the ionisation cones."," At strictly edge-on viewing directions, $\Delta \psi$ becomes very small, especially for higher optical depths of the ionisation cones."
+" But, as we discuss in the following, our modelling approach is rather conservative, and in reality we should expect more favourable conditions for the measurement of Aw."," But, as we discuss in the following, our modelling approach is rather conservative, and in reality we should expect more favourable conditions for the measurement of $\Delta \psi$."
+ The results we present in this work give lower limits on Δ for NGC 1068 assuming the least favourable conditions for a rotation of the polarisation angle between low and high photon energies., The results we present in this work give lower limits on $\Delta \psi$ for NGC 1068 assuming the least favourable conditions for a rotation of the polarisation angle between low and high photon energies.
+ The net Stokes flux and therefore the resulting polarisation angle is determined by the competition between the disc/torus system and the polar scattering cones., The net Stokes flux and therefore the resulting polarisation angle is determined by the competition between the disc/torus system and the polar scattering cones.
+ Especially at higher inclination the polar scattering has a strong impact as the scattering-induced polarisation rises towards orthogonal scattering angles., Especially at higher inclination the polar scattering has a strong impact as the scattering-induced polarisation rises towards orthogonal scattering angles.
+ The efficiency of polar scattering depends on Τοοπο and is low for optically thin scattering cones., The efficiency of polar scattering depends on $\tau_{\rm cone}$ and is low for optically thin scattering cones.
+ It rises with increasing Toone and then goes through amaximum until multiple scattering effects, It rises with increasing $\tau_{\rm cone}$ and then goes through amaximum until multiple scattering effects
+pick-out a roughly diagonal line in Table 2.. and Óuplies that the radius of remmauts at the cud of the quasi-Sedov-Tavlor stage8 has less variance with ny0 than would otherwise: be the case.,"pick-out a roughly diagonal line in Table \ref{tab:rad_qst}, and implies that the radius of remnants at the end of the quasi-Sedov-Taylor stage has less variance with $n_{0}$ than would otherwise be the case."
+" However. since: // ""nis also clepeudeut on the nuuber density of cold clouds we do not expect a particularly tight relationship."," However, since $f'$ is also dependent on the number density of cold clouds we do not expect a particularly tight relationship."
+ Simple. approximations. that fit. the evolution. of. . ↑∐↸∖↥⋅⋜⋯∶↴∙⊾↸∖∪↕↴∖↴↿∏⋉∖↥⋅∐∪↖↽⋜↧↥⋅↸∖⋯∐⋜⋯↑↴∖↴↖↖⇁↕∐↸⊳∐↸⊳∪∐≼⊔∐⊳⊓↖⇁↸∖↕⋅↖↽ : ⋅ mass load. and which are complementary to simular," Simple approximations that fit the evolution of the range of supernova remnants which conductively mass load, and which are complementary to similar"
+We have exiiiued the morphological properties of PNe οι the USNO-PN sample using the radio images of these objects. found iu the literature.,We have examined the morphological properties of PNe from the USNO-PN sample using the radio images of these objects found in the literature.
+ We use conrparisons with radio morphological classes eiven by Aaquist&Kwok(1996) which are based on theimodcl (PES)., We use comparisons with radio morphological classes given by \cite{1996ApJ...462..813A} which are based on the (PES).
+ The PES model describes divergency in observed. (Cenipirical) morphologics through different sky projections of relatively simple spherical shell nodels with both radial and latitude density eracdients and with different ionization depths, The PES model describes divergency in observed (empirical) morphologies through different sky projections of relatively simple spherical shell models with both radial and latitude density gradients and with different ionization depths.
+ USNO-PN sample objects can be classified as follows: circular (NGC 7293 or Helix. Nebula. see. Zijlstraetal. 1989)): open elliptical (NCC) 6853. or Diuubbell Nebula. see Dienell 1983)): elliptical. (NGC 6720 or Biug Nebula. see Ceoreeetal. 1971): svammetric (A2! or S271. see Salterctal.1981 2): and S-type I)," USNO-PN sample objects can be classified as follows: circular (NGC 7293 or Helix Nebula, see \cite{1989AAS...79..329Z}) ); open elliptical (NGC 6853 or Dumbbell Nebula, see \cite{1983IAUS..103...69B}) ); elliptical (NGC 6720 or Ring Nebula, see \cite{1974AA....35..219Gnp}) ); symmetric (A21 or S274, see \cite{1984AA...137..291Snp}) ); and S-type )."
+ Ilowever. a urther search of the literature reveals that the intrinsic uorphologies for all five of these objects is most likely. to )e bipolar. i.c. that the structures of the emüittiug regions (or shells) range from a thick disk (Helix Nebula) to a riaxial ellipsoid (À21. and Ring Nebula) to a ike shape (À21. and Duiibbell Nebula) (see O'Dell1998.eureal.1999.Tia&KwokοDellet2007 and Aleaburneal. 20003).," However, a further search of the literature reveals that the intrinsic morphologies for all five of these objects is most likely to be bipolar, i.e. that the structures of the emitting regions (or shells) range from a thick disk (Helix Nebula) to a triaxial ellipsoid (A24, and Ring Nebula) to a barrel-like shape (A21, and Dumbbell Nebula) (see \cite{1998AJ....116.1346Onp,
+1999ApJ...517..782H,1999AAS..138..275Hnp,2007AJ....134.1679O} and \cite{2005RMxAA..41..109M}) )."
+" The sixth object for which we found reliable. radio observatious (AT) appears to be a ""classical PN with a well defined spherical shel Nilourisctal. 1996).", The sixth object for which we found reliable radio observations (A7) appears to be a “classical” PN with a well defined spherical shell \cite{1996AA...310..603X}) ).
+" Wo note that Ixastueretal.(1996) ax Zuckerman&Catley(1988) detected. a shocl-excited Tl, enussion from the Πο. Ringe anc Duubbell uebulae. Προς an ionization bound case for these objects."," We note that \cite{1996ApJ...462..777K} and \cite{1988ApJ...324..501Z}
+ detected a shock-excited $_2$ emission from the Helix, Ring and Dumbbell nebulae, implying an ionization bounded case for these objects."
+ Additionally. no spectroscopic evidence for the existence of a high-velocity stellar wind las been fouud for the previously meutiouc objects. including AT (Cerruti-Sola&Perinotto1985.Patriarcli&Peorinetto L99L)).," Additionally, no spectroscopic evidence for the existence of a high-velocity stellar wind has been found for the previously mentioned objects, including A7 \cite{1985ApJ...291..237C,1991AAS...91..325P}) )."
+ Central stars of these objects are nost likely well pass the lyvdrogcu-shel ring phase aud approaching the white dwarf cooling sequence., Central stars of these objects are most likely well pass the hydrogen-shell burning phase and approaching the white dwarf cooling sequence.
+ A wide range of shell structures is predicted in the ivdrodyvuuuical models based oun the ISW. model aud it is very likely that a large number of PNe classified * apparent iuorphologv as spherical or elliptical. nay du fact posses a bipolar structure.," A wide range of shell structures is predicted in the hydrodynamical models based on the ISW model and it is very likely that a large number of PNe classified by apparent morphology as spherical or elliptical, may in fact posses a bipolar structure."
+" Even though uorphological classificatious based on optical muaeius ""wor elliptical objects. Zuckerman&Aller(1986). predict hat approximately of all PN in the Galaxy are actually bipolar."," Even though morphological classifications based on optical imaging favor elliptical objects, \cite{1986ApJ...301..772Z} predict that approximately of all PN in the Galaxy are actually bipolar."
+ Nevertheless. we conclude that the USNO-PN sample is not representative of the majority of PNe because of its “narrowness” in morphology (d.c. aliiost all appear to be bipolar).," Nevertheless, we conclude that the USNO-PN sample is not representative of the majority of PNe because of its “narrowness” in morphology (i.e. almost all appear to be bipolar)."
+ The main results of this paper may be sunmniarized as follows:, The main results of this paper may be summarized as follows:
+Figure & displays the observed cumulative white chwart παλπιοτν function iu M67 from Richer 1998 (heavy solid line) compared with svuthetic luuinosity fuuctious or clusters with ages of 3. 1 and 5 Cis.,"Figure 8 displays the observed cumulative white dwarf luminosity function in M67 from Richer 1998 (heavy solid line) compared with synthetic luminosity functions for clusters with ages of 3, 4 and 5 Gyrs."
+ The svuthetic Tuctions have been shifted so as to represent a cluster with an apparent Wo distance modulus of 9.59 (Aloutegomery 1993) and they have been normalized to cota he same nunber of white dwarfs that are observed iu he cluster (58)., The synthetic functions have been shifted so as to represent a cluster with an apparent $V$ distance modulus of 9.59 (Montgomery 1993) and they have been normalized to contain the same number of white dwarfs that are observed in the cluster (58).
+ All the svuthetic functions lave Salpeter IMEs but the actual choice of the IME. within rather broad limits. has rather little effect on the final results as could be deduced from Figure 6.," All the synthetic functions have Salpeter IMFs but the actual choice of the IMF, within rather broad limits, has rather little effect on the final results as could be deduced from Figure 6."
+ Frou the faint eud of the huuinositv fuuction seen m MOT. it is clear that the cooling age of the cluster is larger than 3 Corr. less than 5 Cyr and that the £C yr svuthetie luminosity function is an excellent match to the Iunuinositv function of the faiutest observed cluster white dwarfs.," From the faint end of the luminosity function seen in M67, it is clear that the cooling age of the cluster is larger than 3 Gyr, less than 5 Gyr and that the 4 Gyr synthetic luminosity function is an excellent match to the luminosity function of the faintest observed cluster white dwarfs."
+ This result iudicates that a properly constructed svuthetic white dwart Iwuuinosity function compared with data should be a robust aud reliable age indicator. aud hat it will be an important tool in establishing ages for old clusters iu the Galaxy.," This result indicates that a properly constructed synthetic white dwarf luminosity function compared with data should be a robust and reliable age indicator, and that it will be an important tool in establishing ages for old clusters in the Galaxy."
+ We note in passing that the observed white dwarf muinositv function im M67 has a well populated tail of stars to high hunimositv. many more stars than are xedieted. bv the models.," We note in passing that the observed white dwarf luminosity function in M67 has a well populated tail of stars to high luminosity, many more stars than are predicted by the models."
+ Varyiug the IMF. even by a rather large amount. could not make the fit ofthe svuthetic uction to the observations sieuificantly better as the WeCULSOL niss range among the bright M67 white dwafs is quite πα.," Varying the IMF, even by a rather large amount, could not make the fit of the synthetic function to the observations significantly better as the precursor mass range among the bright M67 white dwarfs is quite small."
+ The origina of this tail is not currently understood but aight be related to the high binary raction in the cluster (see Richer 1998 for further discussion) or to some deficiency in the cooling models whichoverestimates the true rate of cooling ofvoung white cvarfs., The origin of this tail is not currently understood but might be related to the high binary fraction in the cluster (see Richer 1998 for further discussion) or to some deficiency in the cooling models which the true rate of cooling of young white dwarfs.
+ Π the binary scenario is correct. the excess ununboer of bright white dwarfs could be produced by making a relatively huge number of bela core white charts via truncated stellar evolution.," If the binary scenario is correct, the excess number of bright white dwarfs could be produced by making a relatively large number of helium core white dwarfs via truncated stellar evolution."
+ Such objects fade less rapidly than C-O white dwarts as they have a ereater heat capacity per unit nass., Such objects fade less rapidly than C-O white dwarfs as they have a greater heat capacity per unit mass.
+ The mücroleusiug ορποιές in the direction of the LAIC κοσ to be indicating that 60+20% of the dark matter in the Galactic halo is tied up iu objects (Alcock 1997a. 1997b: Renault0.5035. 1997).," The microlensing experiments in the direction of the LMC seem to be indicating that $60 \pm 20$ of the dark matter in the Galactic halo is tied up in $0.5^{+0.3}_{-0.2} M_{\odot}$ objects (Alcock 1997a, 1997b; Renault 1997)."
+ This naturally. sugeests old white dwarfs although. other possibilities exist (e.g. neutron stars. primordial black holes).," This naturally suggests old white dwarfs although other possibilities exist (e.g. neutron stars, primordial black holes)."
+ The possibility that white cdwairfs are nmuportaut contributors to the Galactic dark iatter bas been considered for some time now (Larson 1986: Silk 1991: Cary 1991). but with the mucrolensing results this scenario its taken on increased viability.," The possibility that white dwarfs are important contributors to the Galactic dark matter has been considered for some time now (Larson 1986; Silk 1991; Carr 1994), but with the microlensing results this scenario has taken on increased viability."
+ Chabricr 1999. οματια 1996. Gibsou aud Mould 1997. and Adams aud Lanehlin 1996 have all pointed out hat if indeed this scenario is correct. the IME of the white dwart precursors could not have had a Salpeter form but welt have been more Gaussian im shape aud peaked uear 27AL...," Chabrier 1999, Chabrier 1996, Gibson and Mould 1997, and Adams and Laughlin 1996 have all pointed out that if indeed this scenario is correct, the IMF of the white dwarf precursors could not have had a Salpeter form but might have been more Gaussian in shape and peaked near $2.7M_{\odot}$."
+" For this reason we have ealeulated halo luuinosity Muctions for both Salpeter IMEs aud those of the form. Qu)=exp(C787)""nHuny72Jey with; wereT=a—2.7. 4=*2.2 aud »=5/5 (Chabrier 1996)."," For this reason we have calculated halo luminosity functions for both Salpeter IMFs and those of the form $\Phi(m) = \exp^{-({\overline m/m)}^{\beta_1}}m^{-\beta_2}$ with $\overline m = 2.7$, $\beta_1 = 2.2$ and $\beta_2 = 5.75$ (Chabrier 1996)."
+" Under this scenario. old white dwarfs will be pleutiful in the Galactic halo but difficult to detect because of their intrinsic faintuess,"," Under this scenario, old white dwarfs will be plentiful in the Galactic halo but difficult to detect because of their intrinsic faintness."
+ The local nuuber of such objects can be determined simply from the local dark matter density (0.0070AZ. νο) (Alcock 19972: Chalvier aud Mévra 1997: Could. Flynn and Bahcall 1996) and the mean white dwarf iiass (Mi p?) through," The local number of such objects can be determined simply from the local dark matter density $0.0079 M_{\odot}$ $^3$ ) (Alcock 1997a; Chabrier and Mérra 1997; Gould, Flynn and Bahcall 1996) and the mean white dwarf mass $\left\langle M_{WD}\right\rangle$ ) through"
+shocks are widespread in these galaxies.,shocks are widespread in these galaxies.
+ We have excluded some points from the central regions of each system where our fitting routine was unable to satisfactorily fit 2 Gaussians to the complex line profiles or where both emission line components are broad enough to exceed the shock cutoff., We have excluded some points from the central regions of each system where our fitting routine was unable to satisfactorily fit 2 Gaussians to the complex line profiles or where both emission line components are broad enough to exceed the shock cutoff.
+ In both systems the ggas traces the HII regions which follow the rotational motion of the parent galaxies., In both systems the gas traces the HII regions which follow the rotational motion of the parent galaxies.
+ In IC 1623 the is dominated by the rotation of the eastern ggasgalaxy and is primarilycomparable to the overall shape of the rotation observed in the CO gas observed by Yunetal.(1994)., In IC 1623 the gas is primarily dominated by the rotation of the eastern galaxy and is comparable to the overall shape of the rotation observed in the CO gas observed by \citet{Yun94}.
+". There is a gap in consistent with the dust lane seen in the HST images, as ggasthe shock excitation is correlated with the dust lane."," There is a gap in gas consistent with the dust lane seen in the HST images, as the shock excitation is correlated with the dust lane."
+ In this region our emission-line analysis only traces the tidally-induced shock excited gas between the two systems because the underlying star formation of the western galaxy is completely extinguished by dust., In this region our emission-line analysis only traces the tidally-induced shock excited gas between the two systems because the underlying star formation of the western galaxy is completely extinguished by dust.
+" The kinematics of the shocked gas are quite distinct, exhibiting very little correlation with the rotation seen in the ggas."," The kinematics of the shocked gas are quite distinct, exhibiting very little correlation with the rotation seen in the gas."
+ The shocked gas in the northern portion of IC 1623a is blue shifted up to 200 The kinematic distinction between the HII region gas and the shocked is more subtle in NGC 3256., The shocked gas in the northern portion of IC 1623a is blue shifted up to 200 The kinematic distinction between the HII region gas and the shocked gas is more subtle in NGC 3256.
+" Where both a shocked componentgas and an HII region component are both measured, however, the shocked gas is mostly blue-shifted."," Where both a shocked component and an HII region component are both measured, however, the shocked gas is mostly blue-shifted."
+ Over the larger part of the visibly star forming regions of NGC 3256 there is a mild shocked component blue shifted a few 10s ofs-!., Over the larger part of the visibly star forming regions of NGC 3256 there is a mild shocked component blue shifted a few 10s of.
+. The dust lane extinguishes the HII region gas as in IC 1623 and in spaxels where both a aand a ccomponent are measured along the dust lane the shocked gas is generally redshifted a few 10s of!., The dust lane extinguishes the HII region gas as in IC 1623 and in spaxels where both a and a component are measured along the dust lane the shocked gas is generally redshifted a few 10s of.
+". In addition, NGC 3256 hosts a galactic wind seen quite clearly in blue shifted Na D (Heckmanetal.1990,2000)."," In addition, NGC 3256 hosts a galactic wind seen quite clearly in blue shifted Na D \citep{Heckman90,Heckman00}."
+. The spectra affected are within a few arcseconds of the nucleus and have a very broad blueshifted tail extending several hundred s-!and a broader narrow component such that both components have strong line ratios and are considered aand thus excluded from the velocity plots., The spectra affected are within a few arcseconds of the nucleus and have a very broad blueshifted tail extending several hundred and a broader narrow component such that both components have strong line ratios and are considered and thus excluded from the velocity plots.
+" Finally, there is a ~ 8""in diameter of regionccomponent strongly blue high-oshifted-up to nearly 200 -lying northeast of the nucleus in a region with little dust and several closely-packed star clusters, possibly a second localized outflow."," Finally, there is a region $\sim8$ in diameter of component strongly blue shifted-up to nearly 200 -lying northeast of the nucleus in a region with little dust and several closely-packed star clusters, possibly a second localized outflow."
+ Our analysis clearly separates the HII region gas from the shocked gas where both are detected., Our analysis clearly separates the HII region gas from the shocked gas where both are detected.
+" In both systems along the dust lanes and in the periphery of the galaxies where star formation either decreases or is extinguished the emission line spectra are dominated by a single, shocked component."," In both systems along the dust lanes and in the periphery of the galaxies where star formation either decreases or is extinguished the emission line spectra are dominated by a single, shocked component."
+" However, our results indicate that even in the regions where emission is dominated by star formation there is still an underlying broad component associated with shocks, induced either by outflows generated by the circumnuclear star formation or by the tidal motions of the gas."," However, our results indicate that even in the regions where emission is dominated by star formation there is still an underlying broad component associated with shocks, induced either by outflows generated by the circumnuclear star formation or by the tidal motions of the gas."
+ To characterize this shock excitation we employ new slow shock models which reproduce our measured emission-line ratios., To characterize this shock excitation we employ new slow shock models which reproduce our measured emission-line ratios.
+ These models are introduced in, These models are introduced in
+noted by ?: the ISO data clearly. show such a correlation for all in our seriesspectrum.,noted by \citet{1987ApJ...318..356H}; the ISO data clearly show such a correlation for all series in our spectrum.
+ This may be due tothe Earth atmosphere affecting the estimate of the weak high level Pfuud series Hues iu the «ρούτα of IEuunaun Simon., This may be due to the Earth atmosphere affecting the estimate of the weak high level Pfund series lines in the spectrum of Hamann Simon.
+ The line widths oftheir stronger lines (c.g. Bra) agree well with those of our spectrmm., The line widths of their stronger lines (e.g. $\alpha$ ) agree well with those of our spectrum.
+ They also agree well with the line widths derived from high-resolution spectra published bv ? aud ?.., They also agree well with the line widths derived from high-resolution spectra published by \citet{1985ApJ...290..325L} and \citet{1983A&A...127..279C}.
+" The FWIIM of the weakest lines Guore than 500 kin D) significantly exceeds that of the photospherie t ysini of 230— km 5,.", The FWHM of the weakest lines (more than 500 km $^{-1}$ ) significantly exceeds that of the photospheric $\sin$ i of 230 km $^{-1}$.
+ While electron scatteriuetC» iav affect the width of the strongest lines (especially at short wavelengths where free-free opacities are small compared to the electron scattering opacity) the IR lines do not show evidence for prominent electron scattering wines., While electron scattering may affect the width of the strongest lines (especially at short wavelengths where free-free opacities are small compared to the electron scattering opacity) the IR lines do not show evidence for prominent electron scattering wings.
+ This suggests that the line width is mainly due to kinematical broadening.c» which would nuplv that the disc is rotatingC» more rapidly than the uuderlving star (ignoring errors in vsini).," This suggests that the line width is mainly due to kinematical broadening, which would imply that the disc is rotating more rapidly than the underlying star (ignoring errors in $\sin$ i)."
+ This cau only occur if some niechanis adds aneular imonientuun to the material injected into the disc., This can only occur if some mechanism adds angular momentum to the material injected into the disc.
+ We note that ? found liue widths of weak cuiission lines in the optical spectra of Be stars that also siguificautlv exceeded the width expected on the basis of the observed sui Trausfer of angular moment lav lead. to spin-downu of the star (?).., We note that \citet{1987A&A...173..299H} found line widths of weak emission lines in the optical spectra of Be stars that also significantly exceeded the width expected on the basis of the observed $\sin$ i. Transfer of angular momentum may lead to spin-down of the star \citep{1998A&A...333L..83P}.
+" Usine the formali given by ?.. and assuniug a full opeuing augle of 2 degrees and py of 3410. 13 & Ὁ, the outflow velocity in the dise near the star can be of the order of 1 kus + without significant spin-down of the star during its main sequence Πίο fiue."," Using the formalism given by \citet{1998A&A...333L..83P}, , and assuming a full opening angle of 2 degrees and $\rho_0$ of $\times$ $^{-11}$ g $^{-3}$, the outflow velocity in the disc near the star can be of the order of 1 km $^{-1}$ without significant spin-down of the star during its main sequence life time."
+ Such à. value is in agreement with the observed line shape., Such a value is in agreement with the observed line shape.
+ The picture which emerges from our analysis of the infrared spectrum of5 Cas is that of a circumstellar region of verv high deusitv. perhaps exceeding 3 ἨἩ oο 3 which is rotating rapidly and whose rotational velocity decreases with distance.," The picture which emerges from our analysis of the infrared spectrum of $\gamma$ Cas is that of a circumstellar region of very high density, perhaps exceeding 3 $^{-11}$ g $^{-3}$, which is rotating rapidly and whose rotational velocity decreases with distance."
+ The rotational velocity iu the disce near the star exceeds that of the photosphere., The rotational velocity in the disc near the star exceeds that of the photosphere.
+ Thisregion is heated by radiatiou from the central star. aud the," Thisregion is heated by radiation from the central star, and the"
+at racius 2. where 6 is the fractional solid angle of the outllow (we shall see later that 6 is of order unity).,"at radius $R$, where $b$ is the fractional solid angle of the outflow (we shall see later that $b$ is of order unity)."
+ As the outflow shock is eflicienthy cooled. the censity is strongly increased (hy a factor ~(e£o)?) there.," As the outflow shock is efficiently cooled, the density is strongly increased (by a factor $\sim (v/\sigma)^2$ ) there."
+ ‘Phus the outflow density in contact with the host LSAL is ‘Phe ISAT density is roughly the isothermal value so combining we find Thus the bubble is Bavleigh.Tavlor stable. and can propag»ropagate outwards and suppressDPI accretion provided tha Since b~ghl. we see that central accretion and SABI growth is likely to be suppressed for AL~AL;(mom) (," Thus the outflow density in contact with the host ISM is The ISM density is roughly the isothermal value so combining we find Thus the bubble is Rayleigh–Taylor stable, and can propagate outwards and suppress accretion provided that Since $b \sim \dot m \sim 1$, we see that central accretion and SMBH growth is likely to be suppressed for $M \sim M_{\sigma}({\rm
+ mom})$. ("
+Note that aceretion of gas within the black hole sphere of influence can continue for a time after the outllow bubble begins to propagate to significant radii.),Note that accretion of gas within the black hole sphere of influence can continue for a time after the outflow bubble begins to propagate to significant radii.)
+" This suggests that SMDII masses in AGN are likely to be only a factor of a few below the AZo limit (2)). as they can grow only when the externally imposed. accretion rate is not far above the Eddington value and so potentially spend. longer at. such Masses,"," This suggests that SMBH masses in AGN are likely to be only a factor of a few below the $M - \sigma$ limit \ref{msigmom}) ), as they can grow only when the externally imposed accretion rate is not far above the Eddington value and so potentially spend longer at such masses."
+ For an energydriven outflow. the preshock outtlow velocity follows from With this change. and. compression only by a factor 4 in the outllow shock. (since cooling is negligible). the SAIBL growth suppression condition corresponding to (11)) becomes which is clearlv impossible to. satisfy for. realistic parameters.," For an energy–driven outflow, the preshock outflow velocity follows from as With this change, and compression only by a factor 4 in the outflow shock (since cooling is negligible), the SMBH growth suppression condition corresponding to \ref{m}) ) becomes which is clearly impossible to satisfy for realistic parameters."
+" ""Thus an energvdriven bubble is never ltavleigh.Tavlor stable and cannot halt SMDLII growth.", Thus an energy–driven bubble is never Rayleigh–Taylor stable and cannot halt SMBH growth.
+ We conclude that SMDBII growth is not halted at the energydriven sequence (3)). and given an adequate mass supply can continue all the way to the momentumclriven limit (2)).," We conclude that SMBH growth is not halted at the energy--driven sequence \ref{msigen}) ), and given an adequate mass supply can continue all the way to the momentum–driven limit \ref{msigmom}) )."
+ We have seen that SMDLL fcocback occurs through massive outllows carrving the Edcington momentum., We have seen that SMBH feedback occurs through massive outflows carrying the Eddington momentum.
+ These outLows must have speeds ο20.1e (ef eqn 6)).," These outflows must have speeds $v
+\simeq 0.1c$ (cf eqn \ref{vmom}) )."
+ Moreover since the momentum outllow rate is specilied by the quantity plcer (cf eqn 7)). the ionization parameter is also specified (here L;=fjLeaq is the ionizing luminosity of the AGN. with /; given by the spectral shape and ionization threshold).," Moreover since the momentum outflow rate is specified by the quantity $\rho R^2v$ (cf eqn \ref{cont}) ), the ionization parameter is also specified (here $L_i = l_i\le$ is the ionizing luminosity of the AGN, with $l_i$ given by the spectral shape and ionization threshold)."
+" Wine (2010) shows that this condition requires that the outflows should have ionization paramoeters where fo=¢;/10 and go,=η0.1. and thus show lines in the Xray. region of the spectrum."," King (2010) shows that this condition requires that the outflows should have ionization parameters where $l_2 = l_i/10^{-2}$ and $\eta_{0.1} = \eta/0.1$, and thus show lines in the X–ray region of the spectrum."
+ This is indeed what is observed. (Pounds et al.," This is indeed what is observed (Pounds et al.,"
+ 2003a. b: O'Brien et al. 2005).," 2003a, b; O'Brien et al, 2005)."
+ This reasoning shows thatΕΕ., This reasoning shows that.
+.. The papers by Tombesi ct al. (, The papers by Tombesi et al. (
+2010a.b) show that outllows with these properties are extremely common.,"2010a,b) show that outflows with these properties are extremely common."
+ They are detected in a significant fraction (235%) of a sample of more than 50 local AGN., They are detected in a significant fraction $>35\%$ ) of a sample of more than 50 local AGN.
+ This must mean first that the solid angle factor b cannot be small Clombesi et al., This must mean first that the solid angle factor $b$ cannot be small (Tombesi et al.
+ deduce b0.6). and second. that à large number of local AGN have undergone orcontinue to undergo episodes of IEddington accretion.," deduce $b\ga 0.6$ ), and second, that a large number of local AGN have undergone orcontinue to undergo episodes of Eddington accretion."
+ We get more information from the hydrogen. column densities. which are observed to lie in the range Ng<107em 7.," We get more information from the hydrogen column densities, which are observed to lie in the range $10^{22}~{\rm cm^{-2}} < N_H < 10^{24}~{\rm cm^{-2}}$ ."
+ Using eqn. (7)), Using eqn. \ref{cont}) )
+" we find where Ai, is the inner radius of the [low.", we find where $R_{\rm in}$ is the inner radius of the flow.
+ For a continuing outflow this would be of order 100 Schwarzschild radii (corresponding to the escape speed ~Ole)., For a continuing outflow this would be of order 100 Schwarzschild radii (corresponding to the escape speed $\sim 0.1c$ ).
+" Lt is difficult, to. detect. the corresponding very. high column densities ~107!em.7 as the gas is likely to be Lully ionized.", It is difficult to detect the corresponding very high column densities $\sim 10^{24}~{\rm cm^{-2}}$ as the gas is likely to be fully ionized.
+ The highest’ observed columns. presumably. correspond. to continuing steady outflows which have recombined at some distance [rom the black hole., The highest observed columns presumably correspond to continuing steady outflows which have recombined at some distance from the black hole.
+ Lower columns may reveal cases where the outllow is episocic. and last stopped at a time Lap before the observation.," Lower columns may reveal cases where the outflow is episodic, and last stopped at a time $t_{\rm off}$ before the observation."
+" In this case Z, takes a larger value 2efi.", In this case $R_{\rm in}$ takes a larger value $\simeq vt_{\rm off}$.
+ Using (17)) we get where Ade=MIOM.NoyNg/f(1077em7).," Using \ref{NH}) ) we get where $M_8 = M/10^8\msun, N_{23} = N_H/(10^{23}~{\rm cm}^{-2})$."
+ Thus even the lowest observed columns (~1ERcm2 7) correspond to outllows which switched. olf only 2 vears ago., Thus even the lowest observed columns $\sim 10^{22}~{\rm cm^{-2}}$ ) correspond to outflows which switched off only 2 years ago.
+ This sugeests that outflows are even more common than one might expect [rom the simplest interpretation of the results of Tomboesi et al (2010a. b).," This suggests that outflows are even more common than one might expect from the simplest interpretation of the results of Tombesi et al (2010a, b)."
+" ""lTomboesi οἱ al. (", Tombesi et al. (
+2010b) consider three sources in detail. (ος 111. 3€ 120. 3C 390.3) and suggest that their luminosities are below Lea.,"2010b) consider three sources in detail (3C 111, 3C 120, 3C 390.3) and suggest that their luminosities are below $\le$."
+ However this procedure. uses. black hole misses based on assumptions that either the Afσ relation. or the SMDBIIbulge mass relation. hold loraff ΟΛΗΣΙ. at least in a statistical sense. (," However this procedure uses black hole masses based on assumptions that either the $M - \sigma$ relation, or the SMBH–bulge mass relation, hold for SMBH, at least in a statistical sense. ("
+Lhe Al0 relation is used to calibrate the reverberation masses for 3€ 120. ος 390.3: Peterson et al..,"The $M-\sigma$ relation is used to calibrate the reverberation masses for 3C 120, 3C 390.3; Peterson et al.,"
+ 2004. and the SMDLL mass for 3€ 111 uses the stellar bulge luminosity: Marchesini et. aL.," 2004, and the SMBH mass for 3C 111 uses the stellar bulge luminosity; Marchesini et al.,"
+ 2004.), 2004.)
+ We have seen that both relations are likely to give overestimates of the SMDLL mass. and hence overestimates of Ze.," We have seen that both relations are likely to give overestimates of the SMBH mass, and hence overestimates of $\le$ ."
+ Figure, Figure
+After the discovery of a significant. albeit. small. degree of linear. polarisation of the afterglow optical lus of GRB 990510 (Covino et al.,"After the discovery of a significant, albeit small, degree of linear polarisation of the afterglow optical flux of GRB 990510 (Covino et al."
+ 1999: Wijers et al., 1999; Wijers et al.
+ 1999). there have oen several other detections of linear. polarisation in a number of afterglows (see Covino ct al.," 1999), there have been several other detections of linear polarisation in a number of afterglows (see Covino et al."
+ 2002 for review)., 2002 for review).
+ Typically. the polarisation is observed. to be at the 1-3 cent level (but see Bersier ct al.," Typically, the polarisation is observed to be at the 1-3 percent level (but see Bersier et al."
+ 2008 for GRB 020405). with constant or smoothly variable level and position angle when associated. with a relatively smooth lighteurve (c.g. CRB 0920519. Gorosabel et al.," 2003 for GRB 020405), with constant or smoothly variable level and position angle when associated with a relatively smooth lightcurve (e.g. GRB 020813, Gorosabel et al."
+ 2004)., 2004).
+ When deviations rom a smooth power-law cecay in the lightcurve are instead. present. polarisation curves show a certain degree of complexity (e.g. GRB 021004. Lazzati et al.," When deviations from a smooth power-law decay in the lightcurve are instead present, polarisation curves show a certain degree of complexity (e.g. GRB 021004 Lazzati et al."
+ 2003: Rol et al., 2003; Rol et al.
+ 2003: GRB 030329 Cireiner et al., 2003; GRB 030329 Greiner et al.
+ 2003)., 2003).
+ In order το observe polarisation. some kind of asvmmetry is needed: this can be provided by patches of coherent magnetic field. as suggested. by Cruzinoy Waxman (1999) and. Cruzinov (1999).," In order to observe polarisation, some kind of asymmetry is needed: this can be provided by patches of coherent magnetic field, as suggested by Gruzinov Waxman (1999) and Gruzinov (1999)."
+ Lo addition. small regions in which the magnetic field has some degree of order could be amplified by scintillation (Medsyedey Loeb 1999). or bv gravitational mini-lensing (Loeb Perna 1998: loka Nakamura 2001).," In addition, small regions in which the magnetic field has some degree of order could be amplified by scintillation (Medvedev Loeb 1999), or by gravitational mini-lensing (Loeb Perna 1998; Ioka Nakamura 2001)."
+ In these cases the required. degree. of asvmmetry is in the structure of the magnetic field. and not in the overall geometry of the fireball. which could even be spherically symmetric.," In these cases the required degree of asymmetry is in the structure of the magnetic field, and not in the overall geometry of the fireball, which could even be spherically symmetric."
+" Recently Granot Ixónnigl (2003) proposed that the required asymmetry. might be provided. by an orderec magnetic field. embedded: in. the cireum-burst material. possibly amplified when the shock propagates in it. to reach values close to equipartition with the energy density of the shocked. material in the ""pulsar wind bubble” scenario (Ixónnigl Ciranot 2002)."," Recently Granot Könnigl (2003) proposed that the required asymmetry might be provided by an ordered magnetic field embedded in the circum-burst material, possibly amplified when the shock propagates in it, to reach values close to equipartition with the energy density of the shocked material in the “pulsar wind bubble"" scenario (Könnigl Granot 2002)."
+ A cillerent class of models postulates that the firebal is jetted (Le. the ejecta are collimated into a cone with half opening angle @))., A different class of models postulates that the fireball is jetted (i.e. the ejecta are collimated into a cone with half opening angle $\theta_{jet}$ ).
+ In this case the observer Likely sees the fireball oll-axis. since the probability to be exactly. on-axis is vanishinely small (CGhisellini Lazzati 1999: Sari 1909. hereafter GLO anc 899. respectively).," In this case the observer likely sees the fireball off-axis, since the probability to be exactly on-axis is vanishingly small (Ghisellini Lazzati 1999; Sari 1999, hereafter GL99 and S99, respectively)."
+" When the fireball bulk Lorentz factor EU is 1/(8;.,6,) (whore 0, is the viewing angle) the emitting surface starts to be asvmmetrical with respect to the line of sight.", When the fireball bulk Lorentz factor $\Gamma$ is $\sim 1/ (\theta_{jet}-\theta_o)$ (where $\theta_o$ is the viewing angle) the emitting surface starts to be asymmetrical with respect to the line of sight.
+ Moreover it is assumed that a magnetic field is compressed in the plane normal to the motion. analogous to what has been proposed," Moreover it is assumed that a magnetic field is compressed in the plane normal to the motion, analogous to what has been proposed"
+instabilities.),instabilities.)
+" If ve express the scale rey fod in that study iu terms of ourscale ry. we fud Note that t106 Vetioresis Whie AL, increases. aud that its vali cis only weakly depeucdct ona and ou AL."," If we express the scale $r_Q$ found in that study in terms of ourscale $r_s$, we find Note that the ratio$r_Q/r_s$ while $M_{\star}$ increases, and that its value is only weakly dependent on $\alpha$ and on $\dot{M}$."
+" For parameters typical of an ACN, sel-regulation may thus be eusured very far in. while an extrapolation of the above recipe to aralucters typical of a «isk surrounding a T Taun star suggests that. for these latter objects. re should becoue Org)."," For parameters typical of an AGN, self-regulation may thus be ensured very far in, while an extrapolation of the above recipe to parameters typical of a disk surrounding a T Tauri star suggests that, for these latter objects, $r_Q$ should become $O(r_s)$."
+ Some exines of computed on the basis of Eq. (18)) (, Some examples of computed on the basis of Eq. \ref{qprofile}) ) (
+see Fie. 7)),see Fig. \ref{fig:qprofile}) )
+ show how an outer disk dominated by selfgravity can match iu detail with an inner standard Keplerian accretion disk., show how an outer disk dominated by self-gravity can match in detail with an inner standard Keplerian accretion disk.
+ So far we have cosidered the «vase where the nass is all distributed ina cisk (either at t1e center. as a polut lass. or in diffuse form).," So far we have considered the case where the mass is all distributed in a disk (either at the center, as a point mass, or in diffuse form)."
+ In view o: possible applications to ACN configuratiois or to the οeneral galactic context. it 1s portant to coisider au extension of the iiodels to the case where part of the eravitaional field is determined by a diffuse spherical component. which we will call (even if it may jtst correspolk to the ceutral region of an elliptical galaxy).," In view of possible applications to AGN configurations or to the general galactic context, it is important to consider an extension of the models to the case where part of the gravitational field is determined by a diffuse spherical component, which we will call (even if it may just correspond to the central region of an elliptical galaxy)."
+ This wil lead to rotation curves otherwise not accessible bv our models., This will lead to rotation curves otherwise not accessible by our models.
+ Ou the other haud. it ds casily recoguized that this natural exteusiou is come to leave the slow deusitv decline of the disk unaltered.," On the other hand, it is easily recognized that this natural extension is going to leave the slow density decline of the disk unaltered."
+ Therefore. if we are mterested iu producing models with finite total amass. we should be ready to nupose an outer truncation radius or. which might effectively be equivalent. to relax the ποτοσαΊο prescriptio rin the outermost disk.," Therefore, if we are interested in producing models with finite total mass, we should be ready to impose an outer truncation radius or, which might effectively be equivalent, to relax the self-regulation prescription in the outermost disk."
+ We have thus considered a set of models where he field external to the disk is produced by the joint coutribution of a central poiut mass aud of a halo (which. for simplicity. we take to be spherical).," We have thus considered a set of models where the field external to the disk is produced by the joint contribution of a central point mass and of a halo (which, for simplicity, we take to be spherical)."
+ In view of the case of a disk enibedded in au elliptical galaxy. we have modeled he halo as approximately isothermal. with a finite core radius.," In view of the case of a disk embedded in an elliptical galaxy, we have modeled the halo as approximately isothermal, with a finite core radius."
+ Iu this case the dimensionless equation giving the rotation curve (Eq. (13))), In this case the dimensionless equation giving the rotation curve (Eq. \ref{velocita}) ))
+ is modified as follows: We see that now the equations depend on two additional parameters: f. ceiving the relative streneth o he external field. aud ry. which measures the size of the core radius.," is modified as follows: We see that now the equations depend on two additional parameters: $f^2$, giving the relative strength of the external field, and $x_0$, which measures the size of the core radius."
+ m lis case it is easv to demonstrate that a aree radii he deusity deviation p approaches f7 if f<1. aud f if f>>1.," In this case it is easy to demonstrate that at large radii the density deviation $\rho$ approaches $f^2$ if $f\ll 1$, and $f$ if $f\gg 1$."
+ Iu Fig., In Fig.
+" S we show examples of the rotation curve of nodels wit1 a ciffuse halo. for the case €= 0,0)= 1."," \ref{fig:alone} we show examples of the rotation curve of models with a diffuse halo, for the case $\xi=0$ , $x_0=1$ ."
+ For he vertica structure. we have referred to the iprovec vertical prescription of Eq. (A9}).," For the vertical structure, we have referred to the improved vertical prescription of Eq. \ref{improvedprescription}) ),"
+ with )., with $\Omega^2_{ext}\propto f^2/(x_0^2+x^2)$ .
+In the DD scenario. it is assumed that every WD merger exceeding 1.4 Μο results in a SN Ia. In our code. which follows the evolution of the progenitor in. detail throughout its existence. without the use of an analytical formalism. there are two typical formation channels which lead to a SN Ia. We have termed them the RLOF channel and the CE channel.,"In the DD scenario, it is assumed that every WD merger exceeding 1.4 $_{\odot}$ results in a SN Ia. In our code, which follows the evolution of the progenitor in detail throughout its existence, without the use of an analytical formalism, there are two typical formation channels which lead to a SN Ia. We have termed them the RLOF channel and the CE channel."
+ A description as well as a typical example of both is given below. since experience has shown that there exists a great diversity in evolutions obtained with different codes (see also Sect.," A description as well as a typical example of both is given below, since experience has shown that there exists a great diversity in evolutions obtained with different codes (see also Sect."
+ 4)., 4).
+ The numbers in this subsection are given for a somewhat stricter subset of the progenitors of DD SNe la. ie. those in which both merging WDs are of the C-O type.," The numbers in this subsection are given for a somewhat stricter subset of the progenitors of DD SNe Ia, i.e. those in which both merging WDs are of the C-O type."
+ If this requirement is waived. the parameter space of progenitors is slightly extended.," If this requirement is waived, the parameter space of progenitors is slightly extended."
+ The implications of this limitation on the eventual DTDs will be discussed in the Results section., The implications of this limitation on the eventual DTDs will be discussed in the Results section.
+ In this channel. the progenitor system at birth consists of a primary typically between 3.7 and 7.1 Mis. a mass ratio between 0.35 and 1.0 (high mass ratios being strongly favored in case of all but the highest primary masses). and an orbital period of up to 100 days.," In this channel, the progenitor system at birth consists of a primary typically between 3.7 and 7.1 $_{\odot}$, a mass ratio between 0.35 and 1.0 (high mass ratios being strongly favored in case of all but the highest primary masses), and an orbital period of up to 100 days."
+ After the nost massive component evolves off the MS and traverses the Hertzsprung gap. a first phase of mass transfer will take place.," After the most massive component evolves off the MS and traverses the Hertzsprung gap, a first phase of mass transfer will take place."
+ The envelope of the donor ts radiative at that time. resulting in a dynamically stable mass transfer on a nuclear or thermal timescale. which is a canonical RLOF.," The envelope of the donor is radiative at that time, resulting in a dynamically stable mass transfer on a nuclear or thermal timescale, which is a canonical RLOF."
+ The standard model is to treat this RLOF conservatively (6= D). re. all matter lost by the donor ts accreted by the other star.," The standard model is to treat this RLOF conservatively $\beta = 1$ ), i.e. all matter lost by the donor is accreted by the other star."
+ The mass transfer will continue until well after mass ratio reversal. resulting in a system with a more extreme mass ratio (where the donor is now the least massive star) and an increased orbital period. typically a few hundred days.," The mass transfer will continue until well after mass ratio reversal, resulting in a system with a more extreme mass ratio (where the donor is now the least massive star) and an increased orbital period, typically a few hundred days."
+ The evolution of the accretor under the effect of mass accretion (1.9. rejuvenation) has been explicitly calculated by Vanbeverenetal.(1998)., The evolution of the accretor under the effect of mass accretion (i.e. rejuvenation) has been explicitly calculated by \citet{vanbeveren1998}.
+. This has been done with a standard accretion model (seeNeoetal.1977) 1f matter directly impacts the surface of the star. and under the assumption of aceretion induced full mixing (seeVanbeveren&DeLoore1994) ifa Keplerian accretion disk is formed.," This has been done with a standard accretion model \citep[see][]{neo1977} if matter directly impacts the surface of the star, and under the assumption of accretion induced full mixing \cite[see][]{vanbeveren1994} if a Keplerian accretion disk is formed."
+ During the collapse of the former donor to a WD. a fraction of its mass will be expelled. resulting in à somewhat larger period still.," During the collapse of the former donor to a WD, a fraction of its mass will be expelled, resulting in a somewhat larger period still."
+ After the initially least massive star eventually also leaves the MS (a timescale which ts critically affected by the amount of mass it received during the first mass transfer phase). it will also increase in radius until it exceeds its Roche lobe.," After the initially least massive star eventually also leaves the MS (a timescale which is critically affected by the amount of mass it received during the first mass transfer phase), it will also increase in radius until it exceeds its Roche lobe."
+ Thus. a second mass transfer phase. toward the WD. will be initiated.," Thus, a second mass transfer phase, toward the WD, will be initiated."
+ Due to the extreme mass ratio of the system and the fact that the aceretor isa WD. this mass transfer will become dynamically unstable. resulting in a CE phase leading to a spiral-in.," Due to the extreme mass ratio of the system and the fact that the accretor is a WD, this mass transfer will become dynamically unstable, resulting in a CE phase leading to a spiral-in."
+ Eventually. the envelope will be ejected and a double WD binary will emerge.," Eventually, the envelope will be ejected and a double WD binary will emerge."
+ This system consists of two WDs with a mass around one solar mass each and an orbital period which has decreased dramatically to a few hours during the spiral-in., This system consists of two WDs with a mass around one solar mass each and an orbital period which has decreased dramatically to a few hours during the spiral-in.
+ Constant emission of GWR will cause the two components to slowly spiral in over the course of a few billion years., Constant emission of GWR will cause the two components to slowly spiral in over the course of a few billion years.
+ Eventually. they will merge. and since they together exceed the Chandrasekhar mass. will under our assumptions result in a SN Ia. The typical timescale of this event is the sum of the formation timescale of the double WD binary and the GWR spiral-in.," Eventually, they will merge, and since they together exceed the Chandrasekhar mass, will under our assumptions result in a SN Ia. The typical timescale of this event is the sum of the formation timescale of the double WD binary and the GWR spiral-in."
+ This 15 typically on the order of a few Gyr. but can be lower for systems which had very short initial periods. resulting in a delay time varying between 0.22 Gyr up to the Hubble time and beyond.," This is typically on the order of a few Gyr, but can be lower for systems which had very short initial periods, resulting in a delay time varying between 0.22 Gyr up to the Hubble time and beyond."
+ If the same initial system is considered. but the first mass transfer phase is taken to be totally non-conservative (B=0) the system will merge before it can result in a double WD binary. and thus there will not be a SN Ia explosion.," If the same initial system is considered, but the first mass transfer phase is taken to be totally non-conservative $\beta=0$ ) the system will merge before it can result in a double WD binary, and thus there will not be a SN Ia explosion."
+" We start from a system on the ZAMS consisting of 4.04+3.6 M,. stars with an initial orbital period of 5.0 days.", We start from a system on the ZAMS consisting of 4.0+3.6 $_{\odot}$ stars with an initial orbital period of 5.0 days.
+ After the first mass transfer phase at 0.20 Gyr. which is a canonical RLOF and is considered to be conservative. a 0.65+7.0 M. system emerges with a period of 160 days.," After the first mass transfer phase at 0.20 Gyr, which is a canonical RLOF and is considered to be conservative, a 0.65+7.0 $_{\odot}$ system emerges with a period of 160 days."
+ Just prior to the originally least massive star filling its Roche lobe. the other star has become a WD with a mass of 0.57 M.« and the period has increased to 170 days.," Just prior to the originally least massive star filling its Roche lobe, the other star has become a WD with a mass of 0.57 $_{\odot}$ and the period has increased to 170 days."
+ This second mass transfer phase at 0.26 Gyr results in a CE evolution (deseribed with the a- formalism). eventually yielding a system of 0.57 + 0.88 M. WDs with a period of 0.19 days.," This second mass transfer phase at 0.26 Gyr results in a CE evolution (described with the $\alpha$ -formalism), eventually yielding a system of 0.57 + 0.88 $_{\odot}$ WDs with a period of 0.19 days."
+ This requires 3.1 Gyr of GWR in order to Merge. resulting in a DD SN Ia after 3.4 Gyr in total.," This requires 3.1 Gyr of GWR in order to merge, resulting in a DD SN Ia after 3.4 Gyr in total."
+ The evolution of this system is represented in the left panel of Figure 1., The evolution of this system is represented in the left panel of Figure 1.
+ Using the same initial conditions but assuming a totally non-conservative RLOF as first mass transfer phase. the system already merges during that phase. and thus does not result in à SN Ia. Alternatively. one can start with a system which ts quite similar in initial masses. Le. à primary mass between 5.3 and 7.9 M. and a mass ratio between 0.84 and 1.0. but with a much larger initial orbital period. starting at 200 days.," Using the same initial conditions but assuming a totally non-conservative RLOF as first mass transfer phase, the system already merges during that phase, and thus does not result in a SN Ia. Alternatively, one can start with a system which is quite similar in initial masses, i.e. a primary mass between 5.3 and 7.9 $_{\odot}$ and a mass ratio between 0.84 and 1.0, but with a much larger initial orbital period, starting at 200 days."
+ The large orbital separation will ensure that the most massive star does not exceed its Roche lobe until after it has developed a deep convective envelope., The large orbital separation will ensure that the most massive star does not exceed its Roche lobe until after it has developed a deep convective envelope.
+ The result is that this mass transfer phase will be dynamically unstable. and will already here result in a CE spiral-in.," The result is that this mass transfer phase will be dynamically unstable, and will already here result in a CE spiral-in."
+ This will decrease the orbital period by about two orders of magnitude., This will decrease the orbital period by about two orders of magnitude.
+ After the envelope is ejected. a binary with a less extreme mass ratio than in the RLOF scenario (the aceretor has not gained any material due to the fast timescale of the mass transfer) and an orbital period of only a few days is obtained.," After the envelope is ejected, a binary with a less extreme mass ratio than in the RLOF scenario (the accretor has not gained any material due to the fast timescale of the mass transfer) and an orbital period of only a few days is obtained."
+ When the originally most massive star has become a WD and the other has started evolving off the MS. a second mass transfer phase will be initiated.," When the originally most massive star has become a WD and the other has started evolving off the MS, a second mass transfer phase will be initiated."
+ Just as before. this will result in à CE phase and spiral-in.," Just as before, this will result in a CE phase and spiral-in."
+ Systems which do not merge during this phase will eventually emerge as a double WD binary of the order of one solar mass each. with an orbital period of only a few tens of seconds.," Systems which do not merge during this phase will eventually emerge as a double WD binary of the order of one solar mass each, with an orbital period of only a few tens of seconds."
+ GWR has to work only a few tens of thousands of years in order to let, GWR has to work only a few tens of thousands of years in order to let
+of the normal-galaxy contribution from that of other sources (Ando&Pavlidou2009).,of the normal-galaxy contribution from that of other sources \citep{vaso09}.
+". Moreover, because normal galaxies seem to dominate theFermi EGB, their small contribution to anisotropies will fortuitously provide the optimal chance of finding smaller and more exotic sources in the observed signal (Siegal-GaskinsPavlidou2009)."," Moreover, because normal galaxies seem to dominate the EGB, their small contribution to anisotropies will fortuitously provide the optimal chance of finding smaller and more exotic sources in the observed signal \citep{sgp09}."
+". It is a pleasure to thank Marco Ajello, Chuck Dermer, Hai Fu, Troy Porter, and Andy Strong for stimulating discussions."," It is a pleasure to thank Marco Ajello, Chuck Dermer, Hai Fu, Troy Porter, and Andy Strong for stimulating discussions."
+ BDF thanks the Goddard Space Flight Center for hospitality while some of this work was done., BDF thanks the Goddard Space Flight Center for hospitality while some of this work was done.
+" This work was partially supported by NASA, via Fermi GI Program grants NNXO9AT74G and NNX09AUOIG, and the Astrophysics Theory Program through award NNX10AC86G. VP acknowledges NASA support through Einstein Postdoctoral Fellowship grant number PF8-90060 awarded by the Chandra X-ray Center, which is operated by the Smithsonian Astrophysical Observatory for NASA under contract NAS8-03060."," This work was partially supported by NASA, via Fermi GI Program grants NNX09AT74G and NNX09AU01G, and the Astrophysics Theory Program through award NNX10AC86G. VP acknowledges NASA support through Einstein Postdoctoral Fellowship grant number PF8-90060 awarded by the Chandra X-ray Center, which is operated by the Smithsonian Astrophysical Observatory for NASA under contract NAS8-03060."
+ VP would like to thank the Physics Department at the University of Crete for their hospitality during part of this work., VP would like to thank the Physics Department at the University of Crete for their hospitality during part of this work.
+" The work of TP is supported in part by the Provincial Secretariat for Science and Technological Development, and by the Ministry of Science of the Republic of Serbia under project 141002B."," The work of TP is supported in part by the Provincial Secretariat for Science and Technological Development, and by the Ministry of Science of the Republic of Serbia under project 141002B."
+"distorted geometry& of a degenerateRuss groundet state (see,Thee.g. al]2004)... comparisonDavidson of","distorted geometry of a degenerate ground state \citep[see, e.g.,][]{davidson83,russ04}."
+" theBordeni|1983} measured spectrum of ionized HBC with the TD-DFT calculated spectra (bottom panel in Figure B» does not permit to exclude the presence of anion bands, in contrast to the case for DBR."," The comparison of the measured spectrum of ionized HBC with the TD-DFT calculated spectra (bottom panel in Figure \ref{fig4}) ) does not permit to exclude the presence of anion bands, in contrast to the case for DBR."
+ The band system with observed origin around 830 nm in the Ne matrix could have its theoretical counterpart in either cation or anion bands predicted around 740 nm., The band system with observed origin around 830 nm in the Ne matrix could have its theoretical counterpart in either cation or anion bands predicted around 740 nm.
+" Using the real-time, real-space implementation of TD-DFT (see Figure DJ), however, these anion bands are calculated to be much further to the red around 1100 nm (0.88 m1), while the cation bands remain around 760 nm µια” 1)."," Using the real-time, real-space implementation of TD-DFT (see Figure \ref{fig2}) ), however, these anion bands are calculated to be much further to the red around 1100 nm (0.88 $\mu$ $^{-1}$ ), while the cation bands remain around 760 nm (1.31 $\mu$ $^{-1}$ )."
+" As already stated, cation bands usually dominate(1.31 the spectra of matrix-isolated photoionized PAHs."," As already stated, cation bands usually dominate the spectra of matrix-isolated photoionized PAHs."
+" Therefore, we tentatively attribute all observed bands of the photoprocessed matrix (except the OH lines) to the HBC cation and discuss the results accordingly."," Therefore, we tentatively attribute all observed bands of the photoprocessed matrix (except the OH lines) to the HBC cation and discuss the results accordingly."
+" The electronic ground state of the HBC cation transforms according to the By, irreducible representation of the Do, point Electronic transitions to 7Ay, *Bou, and ?Ba, states are dipole-"," The electronic ground state of the HBC cation transforms according to the $_{\text{1g}}$ irreducible representation of the $_{\text{2h}}$ point Electronic transitions to $^2$ $_{\text{u}}$, $^2$ $_{\text{2u}}$, and $^2$ $_{\text{3u}}$ states are dipole-allowed."
+ The theoretical spectrum contains several A-type transitions to the semi-occupied orbital at wavelengths longer than 500 nm., The theoretical spectrum contains several A-type transitions to the semi-occupied orbital at wavelengths longer than 500 nm.
+" The lowest Dj transition is actually calculated to be far (Bog)beyond 2 um. —Do(Big)The observed band pattern between 750 and 835 nm may originate from the following calculated transitions: Dg(B3u)-Do(Big) at 732 nm with oscillator strength f=0.038, D7(Au)<-Do(Big) at 724 nm with f=0.039, and Dg(B3u)+-Do(Big) at 686 nm with f=0.081."," The lowest $_{1}$ $_{\text{2g}}$ $\leftarrow$ $_0$ $_{\text{1g}}$ ) transition is actually calculated to be far beyond 2 $\mu$ m. The observed band pattern between 750 and 835 nm may originate from the following calculated transitions: $_{6}$ $_{\text{3u}}$ $\leftarrow$ $_0$ $_{\text{1g}}$ ) at 732 nm with oscillator strength $f=0.038$, $_{7}$ $_{\text{u}}$ $\leftarrow$ $_0$ $_{\text{1g}}$ ) at 724 nm with $f=0.039$, and $_{8}$ $_{\text{3u}}$ $\leftarrow$ $_0$ $_{\text{1g}}$ ) at 686 nm with $f=0.081$."
+" Another transition Dg(Au)+Do(Big) is predicted at 542 nm, close to the measured broad band with maximum at 531 nm."," Another transition $_{9}$ $_{\text{u}}$ $\leftarrow$ $_0$ $_{\text{1g}}$ ) is predicted at 542 nm, close to the measured broad band with maximum at 531 nm."
+" However, its strength (f= 0.025), if it can be ascribed to the measured band, is considerably underestimated which, e.g., could be related to an underestimation of the Jahn-Teller distortion in the ground state by DFT."," However, its strength $f=0.025$ ), if it can be ascribed to the measured band, is considerably underestimated which, e.g., could be related to an underestimation of the Jahn-Teller distortion in the ground state by DFT."
+" At wavelengths shorter than 400 nm, the density of electronic transitions rises quickly."," At wavelengths shorter than 400 nm, the density of electronic transitions rises quickly."
+ A detailed spectroscopic analysis is hardly practicable anymore., A detailed spectroscopic analysis is hardly practicable anymore.
+ 'The comparison of ionized HBC in two different matrix environments and illustrates an important aspect concerning (Arthe absorptionNe) spectra of PAH cations in this wavelength range., The comparison of ionized HBC in two different matrix environments (Ar and Ne) illustrates an important aspect concerning the absorption spectra of PAH cations in this wavelength range.
+" As it was the case for DBR, the spectrum of ionized HBC contains broad bands which resemble broadened versions of analogue bands found in neutral HBC."," As it was the case for DBR, the spectrum of ionized HBC contains broad bands which resemble broadened versions of analogue bands found in neutral HBC."
+" Especially, the π — «* structure around 217.5 nm is even more apparent for the ionized molecule."," Especially, the $\pi$ – $\pi^{*}$ structure around 217.5 nm is even more apparent for the ionized molecule."
+" If the strong band broadening could be ascribed to a matrix-induced effect, larger bandwidths for the Ar compared to the Ne environment would be expected, as interaction-related effects are usually more pronounced for atoms with higher polarizabilities."," If the strong band broadening could be ascribed to a matrix-induced effect, larger bandwidths for the Ar compared to the Ne environment would be expected, as interaction-related effects are usually more pronounced for atoms with higher polarizabilities."
+" The shapes and widths of these features, however, do not seem to depend on the rare gas atoms surrounding the molecule (for A« 550 nm), indicating an intrinsic effect, e.g., very short lifetimes of the excited electronic states."," The shapes and widths of these features, however, do not seem to depend on the rare gas atoms surrounding the molecule (for $\lambda <$ 550 nm), indicating an intrinsic effect, e.g., very short lifetimes of the excited electronic states."
+" On that supposition, low-temperature gas phase spectra of ionized medium-sized to large PAHs, such as DBR. and HBC, at wavelengths shorter than approximately 400 nm would feature band shapes and widths similar to what has been measured for molecules isolated in the inert gas matrix."," On that supposition, low-temperature gas phase spectra of ionized medium-sized to large PAHs, such as DBR and HBC, at wavelengths shorter than approximately 400 nm would feature band shapes and widths similar to what has been measured for molecules isolated in the inert gas matrix."
+" However, a small wavelength shift, usually to the red, has to be considered in that case."," However, a small wavelength shift, usually to the red, has to be considered in that case."
+ The matrix-induced redshift becomes obvious upon comparison of the Ne and Ar measurements., The matrix-induced redshift becomes obvious upon comparison of the Ne and Ar measurements.
+ We will take a closer look at this in the following subsection., We will take a closer look at this in the following subsection.
+ A connection between large PAHs and the DIBs is often assumed., A connection between large PAHs and the DIBs is often assumed.
+" Nevertheless, DBR can be excluded as possible DIB carrier."," Nevertheless, DBR can be excluded as possible DIB carrier."
+" While the case of the neutral molecule is discussed in another article (Rouilléetlal.| cationic DBR does not exhibit sufficiently narrow (2011),bands - at least for wavelengths shorter than 900 nm."," While the case of the neutral molecule is discussed in another article \citep{rouille10}, cationic DBR does not exhibit sufficiently narrow bands - at least for wavelengths shorter than 900 nm."
+ The first electronic transition of HBC in the gas phase causes some narrow and very weak bands between 410 and 435 nm — 1) not able to reproduce any DIB (Kokkin24390etaatheal|2008)., The first electronic transition of HBC in the gas phase causes some narrow and very weak bands between 410 and 435 nm (24390 – 22990 $^{-1}$ ) not able to reproduce any DIB \citep[][]{kokkin08}.
+. Stronger absorption bands appear further in UV where no DIBs were observed., Stronger absorption bands appear further in the UV where no DIBs were observed.
+" The HBC cation, on the other"," The HBC cation, on the other"
+o» Wangetal.(2001) (andusingtherelationbe-weenNyand from1995).,"by \citet{wangetal2001} \citep[and using the relation between \nh~and
+$A_{\rm v}$ ."
+ The spectrum could be fitted with a power-law uodel with iudex of 2.2+00.3 aud a 0.510 keV uuinosity of LT«10°? (assuningadistanceof2.5kpe:IntZaudctal. 2001).," The spectrum could be fitted with a power-law model with index of 0.3 and a 0.5–10 keV luminosity of $1.7 \times 10^{32}$ \citep[assuming a distance of 2.5
+kpc;][]{intzandetal2001}."
+.. A ack body model could not fit the data accurately QC 'deerees of freedoin = 88.7/25)., A black body model could not fit the data accurately $\chi^2$ /degrees of freedom = 88.7/25).
+ A neutron star atmosphere model (thatofZavlin.Pavlov.&Shibauov1996) could fit the data but with a relatively hieh tempcrature AD of —0.2 keV (for an observer at infinity) aud a neutron star radius of [5200.1 kim (as measured on the surface and using a neutron star mass of 1.1 MJ., A neutron star atmosphere model \citep[that of][]{zps1996} could fit the data but with a relatively high temperature $kT$ of $\sim$ 0.2 keV (for an observer at infinity) and a neutron star radius of 0.1 km (as measured on the surface and using a neutron star mass of 1.4 ).
+ The XN-rav spectra of other quiesceut neutrou star fransieuts cano sometimes be described hy a two coniponeut model (asoftthermalcomipo- but with the soft οςuponeut dominating the spectrin (although the power-law component cau occasionally coutribute nearly half the 0.510 keV flux of the source).," The X-ray spectra of other quiescent neutron star transients can sometimes be described by a two component model \citep[a soft thermal
+component below 1 keV and a power-law-like component above a few keV;
+e.g.,][]{asaietal1996,asaietal1998} but with the soft component dominating the spectrum (although the power-law component can occasionally contribute nearly half the 0.5–10 keV flux of the source)."
+ Although not required by the data. the spectra were also fitted using such a two component model. with either a black body or au atinosphere model for the soft component.," Although not required by the data, the spectra were also fitted using such a two component model, with either a black body or an atmosphere model for the soft component."
+ With hese models. the spectra could be accurately fitted (see Tab. 2)).," With these models, the spectra could be accurately fitted (see Tab. \ref{tab:spectrum}) ),"
+ although the atmosphere us power law combination was very unstable and the errors on the parameters are therefore relatively large., although the atmosphere plus power law combination was very unstable and the errors on the parameters are therefore relatively large.
+ The tempcratures obtained for he soft component were similar to those obtained or other quiescent neutron star trausicuts., The temperatures obtained for the soft component were similar to those obtained for other quiescent neutron star transients.
+ The Hux contribution of the soft component to the 4510 keV flux was only ~25% of the total flux in coutrast to the other svstems n which the soft component dominates., The flux contribution of the soft component to the 0.5–10 keV flux was only $\sim$ of the total flux in contrast to the other systems in which the soft component dominates.
+ This paper discusses the observation of the accretion-driven millisecond X-ray pulsar SAN JLsos.l3658 performed during its 2000 outburst., This paper discusses the observation of the accretion-driven millisecond X-ray pulsar SAX J1808.4–3658 performed during its 2000 outburst.
+ Similar to theDeppoSANX observations performed around the same time (oue of those observations was on the same dav as the onc). the observation revealed onlv a weak N-vav source with a 0.510 keV ⋅⋅ ⋅⋅ ↕∏∐∐∐∪↴∖↴↕↑⋅↖↽∪↕∿−⋗∖↓∩⋟−↸∖↥⋅∶↴∙↴∖↴↓⋜↧↑↑∐↸∖↻∪↴∖↴↕⊓∪∐ ∪↕⋟≋⊀≚⊸∖⊽⋅∐≺∖," Similar to the observations performed around the same time (one of those observations was on the same day as the one), the observation revealed only a weak X-ray source with a 0.5–10 keV luminosity of $\sim2\times10^{32}$ at the position of SAX J1808.4–3658."
+∖∩≺∖∖∙↓∶≩⊓⋅↱↗≺∖∖⋅↕≻↸∖↴∖↴↻↕↑↸∖↕↑↴∖↴↖↖⇁↸∖⋜∐↘↽∐↸∖↴∖∷∖↴∙↑↕∐∖ ↴∖↴≺∏∐⋅≼⊳↸∖↖↖↽⋜↧↴∖↴↑∐↸∖↴⋝↥⋅↕∶↴∙⊾∐↑↸∖↴∖↴↑∪∐↸∖∪∐↑∐↸∖↸⊳↸∖∐⊓⋅⋜↧↕≼⊲≼⊲↕≻ ≋↑↸∖∐⋜⊔∖↑⋜↧↕∙⊔∩∣↭⋜⋯≼↧↖↖⊽↕⋅∣↾∐⋜⊔≼↧↴∖↴↸∖↑⋜↧↕∙⋖⊇∩∩⊇⋟ ↕↴∖↴↕∐≼∐∖↸∖≼↧≋∆⊸∖⊽⋅∐≺∖∖∩≺∖∖∙↓∶≩⊓⋅↱⊐≺∖∖⋜⋯≼∐∪↑⋜∐⊔∐∐⋅↸∖↕⋜↧↑↸∖≼⇂ ↴∎↸∖↕≺↧," Despite its weakness, the source was the brightest one on the central CCD proving that the source detected by \citet{stellaetal2000} and \citet{wijnandsetal2002_bepposax} is indeed SAX J1808.4–3658 and not an unrelated field source."
+↴∖↴⋯∐⋅↸⊳↸∖∙⊺∐↸∖↴∖↴⋅↖↽↴∖↴↑↸∖⋯⋜↧↑↕↸⊳∪⇡↽∟↴∖↴↸∖↑↴⋝↸∖↑↖↖⇁↸∖↸∖∐↑↕∐∖ ⊔↸∖⋜↧↴∖↴↿∐⋅↸∖≺↧⋜⋯≼↧↑∐∖⊓⋅⋯∖↻∪↴∖↴↕↑↕∪∐∪↕⋟≋⋎⊸∖⊽⋅∐≺∖∖∩≺∖∖⋅ ⋝↸∖↸∖⊼↻↕⋜∏∐↸∖≼↧↴⋝∙↖↽↑↖↖↽∪∐↸∖⋜∐⋅↴⋝∙↖↽∱∎∐∖≼↧↴∖↴∪⋯⋅↸⊳↸∖, The systematic off-set between the measured and the true position of SAX J1808.4--3658 in those observations can likely be explained by two nearby field sources.
+↴∖↴∙↽∕∏↓↸∖↴∖↴↸∖ ↴∖↴≺∏∐⋅↸⊳↸∖↴∖↴∐∐∶↴∙⊾∐↑↕⋜↧↖↽↸∖⋜↧↕↴∖↴∪↸⊳∪∐↑⋜⊔∏∐⋜↕↑↸∖≼↧↑∐↸∖∏↿↘ ⋜↧↴∖↴↴∖↴↕∶↴⋁∐↸∖≺⊔∪≋⊀≚⊸∖⊽⋅∐≺∖∖∩≺∖∖∙↓∶≩⊓⋅↱⊐≺∖∖≺↴∖↴↸∖↸∖⋜↧↕↴∖↴∪≼⊲⋜⋯∏≻⋜⋯⋜↧ otal.2002).," These sources might have also contaminated the flux assigned to SAX J1808.4–3658 \citep[see
+also][]{campanaetal2002}."
+. Campanaetal.(2002) reported on a quiesceut observation of SAN J1808.13658 performed with about a vear after the cud of the 2000 outburst (the spectral fit parameters reported by Campanaetal.(2005) are also listed in Tab., \citet{campanaetal2002} reported on a quiescent observation of SAX J1808.4–3658 performed with about a year after the end of the 2000 outburst (the spectral fit parameters reported by \citet{campanaetal2002} are also listed in Tab.
+ 2. for comparison)., \ref{tab:spectrum} for comparison).
+ They found that the source had a luminosity of only ~5o10 1. about a factor of L lower than what was ieasured during the 20080 outburst observation.," They found that the source had a luminosity of only $\sim5 \times 10^{31}$ , about a factor of 4 lower than what was measured during the 2000 outburst observation."
+ This demoustrates that at very low luminosities. SAN JlsO0s.L13658 can exhibit variability aud indicates that the source probably displays variability in its quiescent state.," This demonstrates that at very low luminosities, SAX J1808.4–3658 can exhibit variability and indicates that the source probably displays variability in its quiescent state."
+ However. caution is advised because the 2000 outburst observation was performed when the source exhibited violent behavior (Wijuaudsetal.2001)..," However, caution is advised because the 2000 outburst observation was performed when the source exhibited violent behavior \citep{wijnandsetal2001_rxte}."
+" Therefore. if is still possible that in. ""truc quiescence the source will be observed cousisteutly at the low level reported by Campanactal. 2002)."," Therefore, it is still possible that in “true” quiescence the source will be observed consistently at the low level reported by \citet{campanaetal2002}."
+ The spectrum of SAN. JIsOS.E3658 durug the 2000 outburst appears softer than its mescent spectrum: when the spectiin is fitted with a power-law model. the photon iudex is 224003 diving the 2000 outburst vs. 1.5052 in quiescence.," The spectrum of SAX J1808.4–3658 during the 2000 outburst appears softer than its quiescent spectrum: when the spectrum is fitted with a power-law model, the photon index is 0.3 during the 2000 outburst vs. $^{+0.2}_{-0.3}$ in quiescence."
+ The fact that the 2000 outburst spectrum can be accurately ft with a neutron star atmosphere model or with a combination of a thermal plus power-law compoucut (incon-ctal. 2002). also suggests a difference iu the source spectra between the two epochs.," The fact that the 2000 outburst spectrum can be accurately fit with a neutron star atmosphere model or with a combination of a thermal plus power-law component \citep[in contrast with the quiescent spectrum;][]{campanaetal2002}, also suggests a difference in the source spectra between the two epochs."
+" It is interesting to note that in the fwo-coniponen model, the photou index during the 2000 outburs observation ds very simular to that measured during the quiescent observation. suggesting tha the shape of the power-law component might uo have chaneed considerably between the differcu epochs (although the fiux of this component was still a factor of ~3 higher duriug the 2000 outburs"," It is interesting to note that in the two-component model, the photon index during the 2000 outburst observation is very similar to that measured during the quiescent observation, suggesting that the shape of the power-law component might not have changed considerably between the different epochs (although the flux of this component was still a factor of $\sim$ 3 higher during the 2000 outburst"
+"demonstrate that all simulations end up with stable merger objects, which show a general trend of increasing effective radii with increasing mass.","demonstrate that all simulations end up with stable merger objects, which show a general trend of increasing effective radii with increasing mass."
+ Despite the large range of input Plummer radii of the CCs the effective radii of the merger objects are constrained to values between 10 and 20 pc at the low mass end and to values between 15 and 55 pc at the high mass end., Despite the large range of input Plummer radii of the CCs the effective radii of the merger objects are constrained to values between 10 and 20 pc at the low mass end and to values between 15 and 55 pc at the high mass end.
+ The turnover in the reg vs. Menci space (see Fig. 9)), The turnover in the $r_{\rm eff}$ vs. $M_{\rm encl}$ space (see Fig. \ref{fig_comp_reff_mencl}) )
+ depends on the mass of the initial CC and occurs at larger sizes for higher masses., depends on the mass of the initial CC and occurs at larger sizes for higher masses.
+" The turnover leads to a degeneracy in the reg vs. Mena space of the merger objects, i.e. very different CC parameters can result in a comparable final merger object."," The turnover leads to a degeneracy in the $r_{\rm eff}$ vs. $M_{\rm encl}$ space of the merger objects, i.e. very different CC parameters can result in a comparable final merger object."
+ In addition the turnover leads to a higher probability for merger objects to have intermediate effective radii., In addition the turnover leads to a higher probability for merger objects to have intermediate effective radii.
+" Figure 14 shows r,g as a function of mass of the observed ECs and UCDs (circles) and our models (stars), combining Figs."," Figure \ref{simsandobs} shows $r_{\rm eff}$ as a function of mass of the observed ECs and UCDs (circles) and our models (stars), combining Figs."
+ | and 9.., \ref{fig_ucdgc} and \ref{fig_comp_reff_mencl}.
+ The vast majority of the observed ECs and UCDs are located within the parameter space covered by the modeled merger objects., The vast majority of the observed ECs and UCDs are located within the parameter space covered by the modeled merger objects.
+ Only the very extended objects at Μις~10? Mo and the extremely extended UCDs between Mucp=10’ and 105 Mo are outside the parameter space covered by this study., Only the very extended objects at $M_{\rm EC} \approx 10^5$ $_{\sun}$ and the extremely extended UCDs between $M_{\rm UCD} = 10^7$ and $10^8$ $_{\sun}$ are outside the parameter space covered by this study.
+" In order to verify that less eccentric orbits would produce more extended objects, we calculated two additional models."," In order to verify that less eccentric orbits would produce more extended objects, we calculated two additional models."
+"For the most extended models with the lowest and the highest mass (Rr= 160 pc, M=10°° and 10° Mo) we calculated the evolution on a circular orbit at a galactocentric distance of 60 kpc.","For the most extended models with the lowest and the highest mass $R_{\rm pl}^{\rm CC} =$ 160 pc, $M^{\rm CC} = 10^{5.5}$ and $10^{8}$ $_{\sun}$ ) we calculated the evolution on a circular orbit at a galactocentric distance of 60 kpc."
+ The corresponding merger objects have considerably larger effective radii of 29 pc and 82 pc (see Fig. 14)), The corresponding merger objects have considerably larger effective radii of 29 pc and 82 pc (see Fig. \ref{simsandobs}) )
+" than the merger objects on the eccentric orbit, which have effective radii of 14 and 54 pc."," than the merger objects on the eccentric orbit, which have effective radii of 14 and 54 pc."
+" Due to the lower gravitational field, the masses of the merger objects on the circular orbits are larger than those of the eccentric orbits."," Due to the lower gravitational field, the masses of the merger objects on the circular orbits are larger than those of the eccentric orbits."
+ For the lowest mass model the enclosed mass increases from to50%., For the lowest mass model the enclosed mass increases from to.
+". These results demonstrate that very extended objects like the M31 ECs found by Huxoretal.(2005) and the very extended, high-mass UCDs can be explained by merged cluster complexes in regions with low gravitational fields at large galactocentric radii."," These results demonstrate that very extended objects like the M31 ECs found by \cite{huxor04} and the very extended, high-mass UCDs can be explained by merged cluster complexes in regions with low gravitational fields at large galactocentric radii."
+" The evolution of ECs in a weak gravitational environment has also been studied by Hurley&Mackey (2010), who performed direct N-body models of extended low-mass star"," The evolution of ECs in a weak gravitational environment has also been studied by \cite{hurley}, , who performed direct N-body models of extended low-mass star"
+is given by the following form: where the subscript. 7. represents each absolute / magnitude bin in the hIumninosity funelion: D; and V; are the measured space density and the volume explored in each magnitude bin of the IXCAD dataset. respectively: and 8; is the model space densitv for the 7th J magnitude bin.,"is given by the following form: where the subscript, $i$, represents each absolute $J$ magnitude bin in the luminosity function; $D_i$ and $V_i$ are the measured space density and the volume explored in each magnitude bin of the KCAB dataset, respectively; and $\theta_i$ is the model space density for the $i$ th $J$ magnitude bin."
+ The final unnormalized posterior distribution is given as the natural log; of the product over each magnitude bin of Equation 10: The prior distribution is outside the summation because it is invariant of the specilic data point under consideration., The final unnormalized posterior distribution is given as the natural log of the product over each magnitude bin of Equation 10: The prior distribution is outside the summation because it is invariant of the specific data point under consideration.
+ To caleulate the posterior distribution. we generate a series of mass-age distributions and transform them to the observational plane via the brown dwarl models. as described in $2.," To calculate the posterior distribution, we generate a series of mass-age distributions and transform them to the observational plane via the brown dwarf models, as described in $\S2$."
+ The resolution of the mass-Duncetion model parameters are Ads=0.05 and AL..., The resolution of the mass-function model parameters are $\Delta{\alpha_2} = 0.05$ and $\Delta{m_{12}} = 0.001~M_{\odot}$ .
+" We normalize each iteration to match the space density for objects wilh AM, = 1112.5 in the IKCAD Inminosity function.", We normalize each iteration to match the space density for objects with $M_J$ = 11–12.5 in the KCAB luminosity function.
+ These data sample the most luminous ultracool cdwarfs aud are likely to provide (he most reliable space density estimates., These data sample the most luminous ultracool dwarfs and are likely to provide the most reliable space density estimates.
+ Figure 8. displavs the resultant 2D posterior distribution for o» and mq»., Figure \ref{fig:ka4} displays the resultant 2D posterior distribution for $\alpha_2$ and $m_{12}$.
+ The most probable solutionis ~o»=0.0 and ο~0.08M. with large uncertainties., The most probable solution is $\sim \alpha_2 = 0.0$ and $M_{12} \sim 0.08~M_{\odot}$ with large uncertainties.
+ We also test the effect that varving the a» prior distribution has on the output posterior distribution., We also test the effect that varying the $\alpha_2$ prior distribution has on the output posterior distribution.
+ Figure 9 displays four 1D. posterior distributions derived from four prior distributions for o»: the nominal case. a shifted case. a wider case. aud a narrower case.," Figure \ref{fig:twoplp} displays four 1D posterior distributions derived from four prior distributions for $\alpha_2$; the nominal case, a shifted case, a wider case, and a narrower case."
+ The overall shape ancl peak location of the posterior distribution remains largely the same despite the variations of the prior distribution., The overall shape and peak location of the posterior distribution remains largely the same despite the variations of the prior distribution.
+ The narrow prior distribution produces the biggest differences., The narrow prior distribution produces the biggest differences.
+ However. all the posterior distributions cover similar ranges in à» with similar amplitudes.," However, all the posterior distributions cover similar ranges in $\alpha_2$ with similar amplitudes."
+ Although the posterior distributions differ substantially from their priors. (hev are still affected. by them and only weakly constrain model parameter values.," Although the posterior distributions differ substantially from their priors, they are still affected by them and only weakly constrain model parameter values."
+ Consequently. the different output posterior distributions all fit the data equally well.," Consequently, the different output posterior distributions all fit the data equally well."
+ The choice of a 6wo-segment power law is not clearly required. so we also fit the data with other forms of the mass function in an effort to determine which. if any. provides the," The choice of a two-segment power law is not clearly required, so we also fit the data with other forms of the mass function in an effort to determine which, if any, provides the"
+,.
+ The first integral of gives our final result. which may be used in the code: Al = ," The first integral of gives our final result, which may be used in the code: M = ."
+Not all functional forms of ης) are acceptable choices for example. taking 5 to be a linear function of ¢ leads to a magnetic force which diverges on the polar axis.," Not all functional forms of $\eta(\zeta)$ are acceptable choices — for example, taking $\eta$ to be a linear function of $\zeta$ leads to a magnetic force which diverges on the polar axis."
+ In this paper we will work with two other choices of η., In this paper we will work with two other choices of $\eta$.
+" The first (which we will refer to as ""C7-superconductivity! for brevity) is with gj a constant. which may be varied to adjust the magnetic field strength."," The first (which we will refer to as $\zeta^2$ -superconductivity' for brevity) is ^2 with $\eta_0$ a constant, which may be varied to adjust the magnetic field strength."
+ With this we have ο = , With this we have M = =.
+"This form ο) 44, is the same as that in our normal-MITD toroidal-Deld ease.", This form of $B_\phi$ is the same as that in our normal-MHD toroidal-field case.
+" The second functional form we will work with (hereafter ""C'-superconductivitv) is cla.", The second functional form we will work with (hereafter $\zeta^3$ -superconductivity') is ^3.
+ ςThe magnetic force scalar and magnetic field in this case are given by AL = M b= lip, The magnetic force scalar and magnetic field in this case are given by M = ^2 = ^2.
+s The code we use iteratively solves the equilibrium equations for our neutron star model. and being non-linear is not restricted to the perturbative regime of slow rotation and. weak magnetic fields.," The code we use iteratively solves the equilibrium equations for our neutron star model, and being non-linear is not restricted to the perturbative regime of slow rotation and weak magnetic fields."
+ Our numerical scheme is based on the Hachisu [field (SCE) method (Llachisu 1986).. à more robust extension of an earlier SCE method by Ostriker&Mark(1968).," Our numerical scheme is based on the Hachisu self-consistent field (SCF) method \citep{hachisu}, , a more robust extension of an earlier SCF method by \citet{ostr_mark}."
+. A version of the Llachisu SCE method for magnetised stars was presented by Tomimura&IEriguchi(2005)., A version of the Hachisu SCF method for magnetised stars was presented by \citet{tomi_eri}.
+. Since our scheme is a fairly straightforward. extension of these previous ones. we content ourselves with a summary. here ancl focus on the differences between them: we refer the reader to Hachisu(1986). and Tomimura&Evieuchi(2005). for more details.," Since our scheme is a fairly straightforward extension of these previous ones, we content ourselves with a summary here and focus on the differences between them; we refer the reader to \citet{hachisu} and \citet{tomi_eri} for more details."
+ To find equilibrium models. the user must specify a number of stellar parameters at the outset.," To find equilibrium models, the user must specify a number of stellar parameters at the outset."
+ The major ones are related 1l., The major ones are related 1.
+" the EOS through the polvtropic indices Ny and IN, (related to the compressibility of each 16", the EOS — through the polytropic indices $N_\rmn$ and $N_\rmp$ (related to the compressibility of each 2.
+ shape of the star (related to the rotation rate) specilied through the ratio of polarto equatorial radi rius fru.," the shape of the star (related to the rotation rate) — specified through the ratio of polarto equatorial radii $r_{pole}/r_{eq}$ ,"
+of some multiple image systems.,of some multiple image systems.
+ We have presented new HST//ACS and Keck V/LRIS observations of MACSJJIT49. a massive X-ray. selected galaxy cluster at z=0.544 discovered in the Massive Cluster Survey.," We have presented new /ACS and Keck I/LRIS observations of J1149, a massive X-ray selected galaxy cluster at $z=0.544$ discovered in the Massive Cluster Survey."
+ These data reveal seven robustly identified multiply-imaged galaxies. three of which we have confirmed spectroscopically.," These data reveal seven robustly identified multiply-imaged galaxies, three of which we have confirmed spectroscopically."
+ The most spectacular system is a multiply-imaged face-on disk galaxy at z21.491 that we identify as an L' (Mp~ —20.7) late-type (B/T< 0.5) galaxy with an ⋂∏∶↔⊺∪≣∏⊱↾⋋↾⋅≏∐⋪↑↴⋯⋪⋯∐⊓⋂∏∣⋪∐↾⊜⋂↑↴∿↻↕∖∕↕∙⋝⇁∣≓∶⋔⊜⇂↴∣⋪≣∶↔⊺∣⇈⊖⋋↾ ∣ ≣⋯∐∶↔⊺⊜⋋⋃↑↴⋔≣⋋⊱↾∐∣∐⇀↸⋝⇁⋅≏⊔⋪⊜↕⋯↧∶↔⊺∏≣⊓⊜∐⇂↴⋝⇁∕∣∶⊐∍, The most spectacular system is a multiply-imaged face-on disk galaxy at $z=1.491$ that we identify as an $L^\star$ $M_B\simeq-20.7$ ) late-type $B/T\ls0.5$ ) galaxy with an ongoing star formation rate of $\sim6\Msolpyr$; the brightest images of this galaxy are magnified by $\mu=23$.
+↜∣∶∐↾⋯⋪⊖ observations using integral field spectrographs should probe its properties in exquisite detail. thanks to the combination of lens magnification and fortuitous viewing angle.," Future observations using integral field spectrographs should probe its properties in exquisite detail, thanks to the combination of lens magnification and fortuitous viewing angle."
+ We use the positions and redshifts of robustly identified multiply-imaged galaxies to constrain a detailed model of the mass and structure of the cluster core., We use the positions and redshifts of robustly identified multiply-imaged galaxies to constrain a detailed model of the mass and structure of the cluster core.
+" Our fiducial model contains the main cluster halo plus three group-scale halos: the probability of a model this complex. relative to less complex models is PCNhato=4)/PONSas<4)10' where Ny, IS the number of cluster/group-scale halos."," Our fiducial model contains the main cluster halo plus three group-scale halos; the probability of a model this complex, relative to less complex models is ${\rm P}(N_{\rm
+ halo}=4)/{\rm P}(N_{\rm halo}<4)\ge10^{12}$ where $N_{\rm halo}$ is the number of cluster/group-scale halos."
+" We measure the mass and fraction of mass residing in substructures to be M(<500kpe)=6.7£0.4x10M, and FuC&500kpe)=0.25+0.12 respectively.", We measure the mass and fraction of mass residing in substructures to be $M(\le500\kpc)=6.7\pm0.4\times10^{14}\Msol$ and $\fsub(\le500\kpc)=0.25\pm0.12$ respectively.
+ In summary. JJ1149 is the most complex strong-lensing cluster core studied to date. its relatively dis-assembled nature being qualitatively consistent with the expectation that clusters at high redshifts are on average less mature than those at lower redshifts.," In summary, J1149 is the most complex strong-lensing cluster core studied to date, its relatively dis-assembled nature being qualitatively consistent with the expectation that clusters at high redshifts are on average less mature than those at lower redshifts."
+ A more complete view will emerge from our analysis of the full sample of MACS clusters at z70.5 (Smith et al..," A more complete view will emerge from our analysis of the full sample of MACS clusters at $z>0.5$ (Smith et al.,"
+ in prep.)., in prep.).
+" We also obtain a power law density profile slope of —0.3£0.05 (95% confidence error bars) on scales of r—3—30"". thereby ruling out density profile slopes as flat as those recently proposed by Zitrin&Broadhurst(2009). at z76 confidence."," We also obtain a power law density profile slope of $\gamma=-0.3\pm0.05$ $95\%$ confidence error bars) on scales of $r\sim3-30\arcsec$ , thereby ruling out density profile slopes as flat as those recently proposed by \cite{Zitrin09b} at $\gs7\sigma$ confidence."
+ In summary. Zitrin&Broadhurst's result can be explained by an absence of multiple-image redshifts of any form in their study. and by them not treating the unknown redshifts as free parameters in their model.," In summary, \citeauthor{Zitrin09b}' 's result can be explained by an absence of multiple-image redshifts of any form in their study, and by them not treating the unknown redshifts as free parameters in their model."
+ These issues are probably compounded by them mis-identifying some multiple-image systems., These issues are probably compounded by them mis-identifying some multiple-image systems.
+ Overall. this underlines the critical importance of measuring spectroscopic redshifts of multiply-imaged galaxies for reliable lens models of strong lensing clusters.," Overall, this underlines the critical importance of measuring spectroscopic redshifts of multiply-imaged galaxies for reliable lens models of strong lensing clusters."
+ GPS thanks Keren Sharon and Phil Marshall for assistance with the Keck observations. and acknowledges Paul May. and Chris Berry for helpful discussions.," GPS thanks Keren Sharon and Phil Marshall for assistance with the Keck observations, and acknowledges Paul May, and Chris Berry for helpful discussions."
+ GPS and IRS acknowledge support from the Royal Society and STFC., GPS and IRS acknowledge support from the Royal Society and STFC.
+ AMS acknowledges a RAS Fellowship., AMS acknowledges a RAS Fellowship.
+ HE. JPK and GPS acknowledge support from STSel under grant GO-09722.," HE, JPK and GPS acknowledge support from STScI under grant GO-09722."
+ ML acknowledges the Centre National d'Etudes Spatiales (CNES) and the Danish National Research Foundation for their support., ML acknowledges the Centre National d'Etudes Spatiales (CNES) and the Danish National Research Foundation for their support.
+ JPK acknowledges support from CNRS., JPK acknowledges support from CNRS.
+are stuall.,are small.
+" Dustead of the values listed in Table 2. the fluxes iu the two-component model are 0.98 and 0.20 Jv. respecivelv. bes fgg, differs bv only —5A conrpared to that iu M2."," Instead of the values listed in Table 2, the fluxes in the two-component model are 0.98 and 0.20 Jy, respecively, i.e., $f_{26 \mum}$ differs by only $\sim 5\%$ compared to that in M2."
+ This indicates that the model Πας of a homogeneous density iu cach Fe-vich shell does not iutroduce a major uncertainty., This indicates that the model assumption of a homogeneous density in each Fe-rich shell does not introduce a major uncertainty.
+ Finally. screening and cooling by dust are potential sources of error in our models.," Finally, screening and cooling by dust are potential sources of error in our models."
+ The effects of dust are exanuued in Ioziu Frausson (1998ab)., The effects of dust are examined in Kozma Fransson (1998ab).
+ The screcnineg we use is discussed in I&oziua Frausson (1998b). aud is based on estimates by Lucy et al. (," The screening we use is discussed in Kozma Fransson (1998b), and is based on estimates by Lucy et al. ("
+1991) and Wooden et al. (,1991) and Wooden et al. (
+1993).,1993).
+ Dust from pure iron is unlikely to fori. as that would cool and block out all iron line cussion once the dust las formed.," Dust from pure iron is unlikely to form, as that would cool and block out all iron line emission once the dust has formed."
+ Ou the contrary. there is a wealth of ion lines from the core at late times.," On the contrary, there is a wealth of iron lines from the core at late times."
+ KF99 aud I&ozia {1999) find goo agreement between modeled and observed broad-baud lightcurves for fZ53270 days aud spectra at 2870 days;," KF99 and Kozma (1999) find good agreement between modeled and observed broad-band lightcurves for $t \lsim 3\,270$ days and spectra at $2\,870$ days."
+ As was poiuted out also for the positron leawoe. the models are also able to reproduce optical Fo 1 ines af even ater epochs.," As was pointed out also for the positron leakage, the models are also able to reproduce optical Fe I lines at even later epochs."
+ Based on this. we believe dust effects are simall enough to neglect iu our estimate of AZ(HTi).," Based on this, we believe dust effects are small enough to neglect in our estimate of $\Ti44$."
+ None of he nmocdel approximations we have used appears to be uncertain cuough to allow fotjay to be of the sale magnitude as fog , None of the model approximations we have used appears to be uncertain enough to allow $f_{24 \mum}$ to be of the same magnitude as $f_{26 \mum}$.
+"The best estimate of AL- (within the framework of our modelue) should therefore conie solely frou foo,44.", The best estimate of $\Ti44$ (within the framework of our modeling) should therefore come solely from $f_{26 \mum}$.
+" To estimate the combined error Of fobjan duc to all model approximations (except for tle dust distribution iu the ejecta). we adopt the uncertainties 15. 15%, 15 and 5 for photoionization. cdistauco. atouuca data aid clumping. respectively."," To estimate the combined error of $f_{26 \mum}$ due to all model approximations (except for the dust distribution in the ejecta), we adopt the uncertainties $15\%$, $15\%$, $15\%$ and $5\%$ for photoionization, distance, atomica data and clumping, respectively."
+ For the choice of input model and positron leakage we adopt 15% each., For the choice of input model and positron leakage we adopt $15\%$ each.
+ This gives a coubined uncertzünutv which is ~31%., This gives a combined uncertainty which is $\sim 34\%$.
+ Ou top of this we mist acd the masinuun systematic error of 12% discussed in Sect_, On top of this we must add the maximum systematic error of $12\%$ discussed in Sect.
+ 3.1 for the lifetime of ! Ti;, 3.1 for the lifetime of $^{44}$ Ti.
+ With the line profile cüscussc (iu Sect., With the line profile discussed in Sect.
+ 3.3. we therefore arrive at an upper limit on AL¢UT) whichis ~1.5.10!AL..," 3.3, we therefore arrive at an upper limit on $\Ti44$ which is $\simeq 1.5 \EE{-4} \Msun$."
+ We uote that this lait excludes the upper euds of the allowed ranges of ΑςΤΙ) fouid by Chlueai ct al. (, We note that this limit excludes the upper ends of the allowed ranges of $\Ti44$ found by Chugai et al. (
+1997) and IKE99 (see Sect.,1997) and KF99 (see Sect.
+ 1)., 1).
+ Combining our Πατ with the preliminary results of IKE99 for the broad-band photometry. a likely range for M( “Tis (0.51.5).10LML..," Combining our limit with the preliminary results of KF99 for the broad-band photometry, a likely range for $\Ti44$ is $(0.5 - 1.5)\EE{-4} \Msun$."
+ We cluphasis. however. that the lower nit of this range is probably more unucertaim than the upper (for the reasons meutioned in Sect.," We emphasis, however, that the lower limit of this range is probably more uncertain than the upper (for the reasons mentioned in Sect."
+ 1). which is indeed indicated by the preliminary analysis of Borkowski et al. (," 1), which is indeed indicated by the preliminary analysis of Borkowski et al. ("
+1997).,1997).
+ Models for the yield of ΤΙ give quite different. results., Models for the yield of $^{44}$ Ti give quite different results.
+ This is most likely due to how the explosion is generated in the models. and how fallback onto the neutron star is treated.," This is most likely due to how the explosion is generated in the models, and how fallback onto the neutron star is treated."
+ Tinuues et al. (, Timmes et al. (
+1996: see also Woosley Weaver 1995) use a piston to generate the explosion. and they accountfor fallback in a rather selfcousistcut wav.,"1996; see also Woosley Weaver 1995) use a piston to generate the explosion, and they accountfor fallback in a rather self-consistent way."
+" In their model with zero-age mass Mzayis=20M. (νου, corresponding to SN 1987À) the mass of the initially ejected HT is 1.2<10.|NL. but of this ouly Lb«LOOAL. escapes after fallback."," In their model with zero-age mass $\MZA = 20 \Msun$ (i.e., corresponding to SN 1987A) the mass of the initially ejected $^{44}$ Ti is $1.2\EE{-4} \Msun$, but of this only $1.4\EE{-5} \Msun$ escapes after fallback."
+ This is less than we arene or in Sect., This is less than we argue for in Sect.
+ L1. aud could suggestOO that fallback was not iuportaut for SN 1987TÀ. though we caution again that he lower Iuit fouud by IKE99 (see also I&ozuia 1999) iav iof be very strict.," 4.1, and could suggest that fallback was not important for SN 1987A, though we caution again that the lower limit found by KF99 (see also Kozma 1999) may not be very strict."
+" If fallback is uunuportaut the ejected amouuts of Ni aud ° Ni world be too hiel in this model. vpically by a factors of 2.L. jucleine from the effects of allback iu the 25ΔΕ, model in Woosley Weaver (1995)."," If fallback is unimportant the ejected amounts of $^{56}$ Ni and $^{57}$ Ni would be too high in this model, typically by a factors of $\sim 2-4$, judging from the effects of fallback in the $25 \Msun$ model in Woosley Weaver (1995)."
+ It should be curphasized that the variation of AZ(HT3) with WMzays in Timunes ct al. (, It should be emphasized that the variation of $\Ti44$ with $\MZA$ in Timmes et al. (
+1996) is complex. aud hat for models with Mzajg;=IsAl. and Αννα=22AL... the caleulated. AL(4Ti) comes within the range we propose. albei close to our lower linüt.,"1996) is complex, and that for models with $\MZA = 18 \Msun$ and $\MZA = 22 \Msun$ , the calculated $\Ti44$ comes within the range we propose, albeit close to our lower limit."
+ The models of Thiclemmanu et a. (, The models of Thielemann et al. (
+1996) simulate the explosion by depositing thermal energv iu the core. aud they insert the nass cut artificially so that the right iuouut of ejected “ONL is produced. (,"1996) simulate the explosion by depositing thermal energy in the core, and they insert the mass cut artificially so that the right amount of ejected $^{56}$ Ni is produced. ("
+This effectively means that fallback is imeluced aso in these models.),This effectively means that fallback is included also in these models.)
+ Simulating the explosion in this way eusures Iareer cutropy and thus more alpha-vich freeze-out than in Woosley Weaver (1995)., Simulating the explosion in this way ensures larger entropy and thus more alpha-rich freeze-out than in Woosley Weaver (1995).
+ Accordingly. the ratio ALIET/AFCONI) (where AL(P°N1) is the ass of ejected 7?Ni that does not fall back) is higher iu the models of Thiclemann et al.," Accordingly, the ratio $\Ti44 / \Nii$ (where $\Nii$ is the mass of ejected $^{56}$ Ni that does not fall back) is higher in the models of Thielemann et al."
+ than in pistou-driven simulaions., than in piston-driven simulations.
+ For example. in the 20XI. inodoel of Thiclemann et al. (," For example, in the $20 \Msun$ model of Thielemann et al. ("
+1996) A£(09N1)z0.071AL... ALON)m2.410OAL. and MIT)zL74101M. with M(C*NI) defined in the same way as M(CONi) and ACHT).,"1996) $\Nii \approx 0.074 \Msun$, $\Niii \approx 2.9\EE{-3} \Msun$ and $\Ti44 \approx 1.7\EE{-4} \Msun$, with $\Niii$ defined in the same way as $\Nii$ and $\Ti44$."
+ The values of MONI) and. ALCNi) are close to what have been 1iferred. for SN. LOSTA (Suutzoff Bouchet 1990: Frausson Wozima 1993)., The values of $\Nii$ and $\Niii$ are close to what have been inferred for SN 1987A (Suntzeff Bouchet 1990; Fransson Kozma 1993).
+ The titaniuu miss is slightly larger han the upper Iuuit of the range we estimate in Sect., The titanium mass is slightly larger than the upper limit of the range we estimate in Sect.
+ L1., 4.1.
+ So. while our estimate of ΑΕΤΙ) cannot rule out with certainty any of the two methods used for the explosion (piston-driven or heat gencrated). our results could iudicate that aji intermediate method should be used (sce also the discussion on this in Tinunes et al.," So, while our estimate of $\Ti44$ cannot rule out with certainty any of the two methods used for the explosion (piston-driven or heat generated), our results could indicate that an intermediate method should be used (see also the discussion on this in Timmes et al."
+ 1996)., 1996).
+" Frou models of the chemical evolution of the Galaxy. and especially thle soar abundance of Ca, a value for ACHTi) closer to that of Thiclemamn et al. ("," From models of the chemical evolution of the Galaxy, and especially the solar abundance of $^{44}$ Ca, a value for $\Ti44$ closer to that of Thielemann et al. ("
+1996) may be more correct. at cast for superuovae in eeneral.,"1996) may be more correct, at least for supernovae in general."
+ Iu this context we note that a lügher value of AJCGLUT3) is produced in asviunetrie explosions (Nagataki et al., In this context we note that a higher value of $\Ti44$ is produced in asymmetric explosions (Nagataki et al.
+ 1998)., 1998).
+ The method of calculation oeniploved by Nagataki et al. (, The method of calculation employed by Nagataki et al. (
+sce Nagataki et al.,see Nagataki et al.
+ 1997) is simular to that in Thiecleiiaun et al. (, 1997) is similar to that in Thielemann et al. (
+1996). though the models of Nagatalà et al.,"1996), though the models of Nagataki et al."
+ allow for 2-D instead of just 1-D. With no asviunetry. the models of Nagataki et al.," allow for 2-D instead of just 1-D. With no asymmetry, the models of Nagataki et al."
+ produce M(HTi)~ος10?M. for an explosion sinuülar SN 1987A. when the mass cut has been trinmed ΑΝ)zm0.07 M...," produce $\Ti44 \sim 6\EE{-5} \Msun$ for an explosion similar to SN 1987A, when the mass cut has been trimmed to $\Nii \approx 0.07 \Msun$ ."
+ Tl‘is is sieuificantlyenifiicautly 1less tlthan Thiclemannu ct al. (, This is significantly less than Thielemann et al. (
+1996) despite the similar method of modeling.,1996) despite the similar method of modeling.
+ Applying au asvmunetrv by a factor of 2 between the equator and the poles. the explosion energy," Applying an asymmetry by a factor of 2 between the equator and the poles, the explosion energy"
+distribution agrees with the distribution of the optically thick disks is0.,distribution agrees with the distribution of the optically thick disks is.
+01%.. We next seek correlations between Lynery and stellar or disk parameters.," We next seek correlations between $L_{\rm [Ne\,II]}$ and stellar or disk parameters."
+" We will employ linear regression to compute functions of the form logLien;=a+blogx, but because many objects show upper limits to Linen, we will use ""survival statistics"" that take these values into account."," We will employ linear regression to compute functions of the form $\log L_{\rm [Ne\,II]} = a + b\log x$, but because many objects show upper limits to $L_{\rm [Ne\,II]}$, we will use “survival statistics” that take these values into account."
+" We use the parametric estimation maximization (EM) algorithm in ASURV, which implements methods presented by Isobeetal. (1986)."," We use the parametric estimation maximization (EM) algorithm in ASURV, which implements methods presented by \citet{isobe86}."
+". We also use ASURV to compute correlation coefficients for the same samples, specifically using the Cox hazard model, Kendall’s tau, and Spearman's rho values (where the latter typically requires at least 30 entries to be accurate)."," We also use ASURV to compute correlation coefficients for the same samples, specifically using the Cox hazard model, Kendall's tau, and Spearman's rho values (where the latter typically requires at least 30 entries to be accurate)."
+ A summary of our statistical results is given in Table 6.., A summary of our statistical results is given in Table \ref{table6}.
+" We first discuss a possible dependence between Lpw.g and the intrinsic, unabsorbed Lx."," We first discuss a possible dependence between $L_{\rm [Ne\,II]}$ and the intrinsic, unabsorbed $L_{\rm X}$."
+ Fig., Fig.
+" 2 relates the two quantities for the entire sample, distinguishing between the three object classes (using different symbol shapes and colors), with separate (open) symbols for upper limits (mostly in Liver)."," \ref{fig2} relates the two quantities for the entire sample, distinguishing between the three object classes (using different symbol shapes and colors), with separate (open) symbols for upper limits (mostly in $L_{\rm [Ne\,II]}$ )."
+" We provide error bars for Γ/Νεπῃ as derived from spectral analysis, while for X-rays spectral-fit errors are normally not relevant as the range of uncertainty is dominated by variability on various time scales."," We provide error bars for $L_{\rm [Ne\,II]}$ as derived from spectral analysis, while for X-rays spectral-fit errors are normally not relevant as the range of uncertainty is dominated by variability on various time scales."
+" Such variability is, apart from singular flares, typically characterized by flux variations within a factor of two from low to high levels."," Such variability is, apart from singular flares, typically characterized by flux variations within a factor of two from low to high levels."
+ We therefore adopted error bars defining flux deviations of V2 to both higher and lower values., We therefore adopted error bars defining flux deviations of $\sqrt{2}$ to both higher and lower values.
+" Note the large range now available in both variables, amounting to «2 dex in Lx and «3 dex in Lynery, i.e., much wider ranges than in previous studies (Espaillatetal.,2007;Lahuisetal.,2007;Pascucci 2007)."," Note the large range now available in both variables, amounting to $\approx$ 2 dex in $L_{\rm X}$ and $\approx$ 3 dex in $L_{\rm [Ne\,II]}$, i.e., much wider ranges than in previous studies \citep{espaillat07, lahuis07, pascucci07}."
+". No sharp correlation is found although a statistically significant dependence exists after excluding the four very strong [Nemn]] detections (for T Tau S, DG Tau, Sz 102, and EC 82) that define the upper envelope of the distribution (a correlation still exists if they are included)."," No sharp correlation is found although a statistically significant dependence exists after excluding the four very strong ] detections (for T Tau S, DG Tau, Sz 102, and EC 82) that define the upper envelope of the distribution (a correlation still exists if they are included)."
+" Three of these objects eject prominent jets (T Tau S, DG Tau, and Sz 102) while EC 82 is a little studied object with a relatively strongly absorbed high-inclination/near-edge-on disk (Pontoppidanetal.,2005)."," Three of these objects eject prominent jets (T Tau S, DG Tau, and Sz 102) while EC 82 is a little studied object with a relatively strongly absorbed high-inclination/near-edge-on disk \citep{pontoppidan05}."
+". The best-fit regression for the remaining sample has a slope of 0.50+0.15, with a low probability, PS6%, for this correlation being attained by chance (Table 6))."," The best-fit regression for the remaining sample has a slope of $0.50\pm 0.15$, with a low probability, $P\la 6\%$, for this correlation being attained by chance (Table \ref{table6}) )."
+" As indicated above, the transition disks behave like the optically thick disks without jets."," As indicated above, the transition disks behave like the optically thick disks without jets."
+" The jet sources, in contrast (shown as blue diamonds in Fig."," The jet sources, in contrast (shown as blue diamonds in Fig."
+" 2 and further figures), are systematically more luminous in [Ner]], revealing only modest overlap with the region occupied by the other objects."," \ref{fig2} and further figures), are systematically more luminous in ], revealing only modest overlap with the region occupied by the other objects."
+" A separate regression analysis for the jet sources indicates a significant dependence with a regression slope of 0.77+0.27, i.e., compatible with proportionality."," A separate regression analysis for the jet sources indicates a significant dependence with a regression slope of $0.77\pm 0.27$, i.e., compatible with proportionality."
+" The dependence is less tight for the non-jet objects although still significant, with a shallower slope of 0.51+0.14."," The dependence is less tight for the non-jet objects although still significant, with a shallower slope of $0.51\pm 0.14$."
+" This trend is shallower than what simple theories would predict, i.e. trends close to proportionality (Meijerink 2009,, see Sect. ??))."," This trend is shallower than what simple theories would predict, i.e. trends close to proportionality \citealt{meijerink08,
+hollenbach09}, see Sect. \ref{nedisks}) )."
+ Fig., Fig.
+" 3 relates Liwerrj to the mass accretion rate, as suggested by Espaillatetal.(2007)."," \ref{fig3} relates $L_{\rm [Ne\,II]}$ to the mass accretion rate, as suggested by \citet{espaillat07}."
+". Again, a large range of Mace values is covered, spanning the interval of 9’10:10—1079Μο γι""."," Again, a large range of $\dot{M}_{\rm acc}$ values is covered, spanning the interval of $\approx 10^{-10}- 10^{-6}~M_{\odot}$ $^{-1}$."
+" No correlation is evident among the jet sources or the disks without jets separately, with P«18—7996."," No correlation is evident among the jet sources or the disks without jets separately, with $P \approx 18-79$."
+". However, stronger accretors are predominantly jet sources, and they reveal higher Linett:"," However, stronger accretors are predominantly jet sources, and they reveal higher $L_{\rm [Ne\,II]}$."
+ A dependence between the two variables is therefore ambiguous., A dependence between the two variables is therefore ambiguous.
+" Although a physical dependence may be absent, the segregation into objects with and without jets may produce an apparent correlation."," Although a physical dependence may be absent, the segregation into objects with and without jets may produce an apparent correlation."
+" Jet engines are typically younger and more active objects, and given a rough correlation between accretion rate and outflow rate (e.g., Hartiganetal. 1995)), jet sources typically also show high accretion rates."," Jet engines are typically younger and more active objects, and given a rough correlation between accretion rate and outflow rate (e.g., \citealt{hartigan95}) ), jet sources typically also show high accretion rates."
+ Our separate finding that jet sources are generally more luminous [Νεπ]], Our separate finding that jet sources are generally more luminous ]
+where r; is the unit vector pointing at the position of cach object and N is the total number of such objects within the distance considered.,where $r_i$ is the unit vector pointing at the position of each object and N is the total number of such objects within the distance considered.
+ Phe weight (1) is calculated as shown in equation 3., The weight $(w_i)$ is calculated as shown in equation 3.
+ Using the linear perturbation theory. and the above definition of the dipole. the relation between the observed peculiar. velocity of an observer. Vp. and that. predicted by its gravitational acceleration. 6. as long as the two vectors are well aligned is: where: 3=m- (bis the bias parameter).," Using the linear perturbation theory, and the above definition of the dipole, the relation between the observed peculiar velocity of an observer, Vp, and that predicted by its gravitational acceleration, g, as long as the two vectors are well aligned is: where: $\beta =\frac {\Omega_o ^{0.6}} {b}$, (b is the bias parameter)."
+ Note that Ἐν) decreases with increasing r. therefore the more distant. objects have less contribution to the acceleration. of the LG. but do continue to ace to the peculiar velocity.," Note that $V_p(r)$ decreases with increasing r, therefore the more distant objects have less contribution to the acceleration of the LG, but do continue to add to the peculiar velocity."
+ From the above equation. the three components of the accelaration. generated. by each object on the LG can be calculated (V5.1.Ve).," From the above equation, the three components of the accelaration generated by each object on the LG can be calculated $(V_x,V_y,V_z)$."
+ Phe force can then be smoothed out to guarantee Linearity. where the function Smooth is defined as: ie.," The force can then be smoothed out to guarantee linearity, where the function Smooth is defined as: ie."
+ cach component is multiplied by the function smooth., each component is multiplied by the function smooth.
+ We then sum over shells to build up the three cartesian components of the cumulative dipole (acceleration) at raclius r. Finally we caleulate the amplitude of the cumulative dipole at the same radius. using: where Visios Veege Verein are the three components of the cumulative acceleration.," We then sum over shells to build up the three cartesian components of the cumulative dipole (acceleration) at radius r. Finally we calculate the amplitude of the cumulative dipole at the same radius, using: where $V_{cum-x}$, $V_{cum-y}$, $V_{cum-z}$ are the three components of the cumulative acceleration."
+ rote that the moments of the objects are allected by the cliseretness effects or the shot noise errors. which increase with redshift’ bacause of the rapidly declining. selection unction.," Note that the moments of the objects are affected by the discretness effects or the shot noise errors, which increase with redshift bacause of the rapidly declining selection function."
+ These effects introduce a variance in the dipole amplitude and direction., These effects introduce a variance in the dipole amplitude and direction.
+ To calculate these ellects. we can use the method of Strauss et al. (," To calculate these effects, we can use the method of Strauss et al. ("
+1992). Llucdson (1993) and Branchini Plionis (1996). as follows: ''herefore: and the corresponding error along the line of sight (1D error) is: The resulting dipole growth curves for the EALSS saniple is shown in figure 4.,"1992), Hudson (1993) and Branchini Plionis (1996), as follows: Therefore: and the corresponding error along the line of sight (1D error) is: The resulting dipole growth curves for the EMSS sample is shown in figure 4."
+ As can be seen. there is no signal observed. at small depths (e.<40h.tAlpe). indicating insignificant contribution to the LG motion from the nearby objects.," As can be seen, there is no signal observed at small depths $(ie. < 40 h^{-1} Mpc)$, indicating insignificant contribution to the LG motion from the nearby objects."
+ This is expected. since there are not many objects in the EMSS sample with distances <40.Ape.," This is expected, since there are not many objects in the EMSS sample with distances $< 40 h^{-1} Mpc$."
+ The amplitude of the dipole growth curve saturates at rather large depths ooLI0hAALpe., The amplitude of the dipole growth curve saturates at rather large depths of $> 300 h^{-1} Mpc$.
+ However from distances of about 3005.! Alpe. the amplitude of the dipole increases only by small amounts.," However from distances of about $300 h^{-1}$ Mpc, the amplitude of the dipole increases only by small amounts."
+ Therefore Roane~300.4005.+AZpe. which is the depth from which these AGNSs contribute to the acceleration of the LG via their gravitational attraction.," Therefore $R_{conv} \sim 300-400 h^{-1} Mpc$, which is the depth from which these AGNs contribute to the acceleration of the LG via their gravitational attraction."
+ Note that there are relatively large error bars. particularly at larger depths. which put limits to the accuracy of the value of νο," Note that there are relatively large error bars, particularly at larger depths, which put limits to the accuracy of the value of $R_{conv}$."
+" From this analysis. the value of brass,Ob is found. to be 0.68+0.22."," From this analysis, the value of $b_{EMSS} \Omega_o^{-0.6}$ is found to be $0.68 \pm 0.22$."